Protostellar Jets in Context (Astrophysics and Space Science Proceedings)

Astrophysics and Space Science Proceedings

Protostellar Jets in Context

K. Tsinganos Editor University of Athens, Greece

T. Ray Editor Dublin Institute for Advanced Studies, Ireland

M. Stute Editor University of Athens, Greece

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Editors K. Tsinganos University of Athens Dept. Physics Section of Astrophysics Panepistimiopolis 157 84 Athens Zografos Greece [email protected]

M. Stute University of Athens Dept. Physics Section of Astrophysics Panepistimiopolis 157 84 Athens Zografos Greece [email protected]

T. Ray Dublin Institute for Advanced Studies Astronomy & Astrophysics Section 31 Fitzwilliam Place Dublin Ireland [email protected]

ISSN 1570-6591 e-ISSN 1570-6605 ISBN 978-3-642-00575-6 e-ISBN 978-3-642-00576-3 DOI 10.1007/978-3-642-00576-3 Springer Heidelberg Dordrecht London New York Library of Congress Control Number: 2009928100 c Springer-Verlag Berlin Heidelberg 2009  This work is subject to copyright. All rights are reserved, whether the whole or part of the material is concerned, specifically the rights of translation, reprinting, reuse of illustrations, recitation, broadcasting, reproduction on microfilm or in any other way, and storage in data banks. Duplication of this publication or parts thereof is permitted only under the provisions of the German Copyright Law of September 9, 1965, in its current version, and permission for use must always be obtained from Springer. Violations are liable to prosecution under the German Copyright Law. The use of general descriptive names, registered names, trademarks, etc. in this publication does not imply, even in the absence of a specific statement, that such names are exempt from the relevant protective laws and regulations and therefore free for general use. Cover design: eStudio Calamar S.L. Printed on acid-free paper Springer is part of Springer Science+Business Media (www.springer.com)

Courtesy: Mark McCaughrean and Hans Zinnecker

Preface

It is over a quarter of a century since the discovery of outflows from young stars. The intervening years have led to remarkable advances in our understanding of this phenomenon. Much of the progress can be attributed to advances in facilities and technologies, including not only larger telescopes but also improved instrument and detector performance. In addition protostellar outflows have now been imaged from the ground and space at high spatial resolution, e.g. with HST, and at a wide variety of wavelengths from X-rays to radio waves, revealing more and more about their physics. This veritable revolution in observation has been accompanied by an exponential growth in our ability to numerically simulate the launching and propagation of jets. Codes continue to improve: they now incorporate more physics and are increasingly efficient through, for example, techniques such as adaptive mesh refinement and the use of parallel processing in cluster environments. Simulating the launching and propagation of a jet all the way from the vicinity of the star up to several thousand AU (a size range of 104 ) is now much closer. In more recent times, developments in observation, theory and numerical simulation have been joined by laboratory jet experiments reproducing, on centimetre scales, that which is seen in astrophysics to stretch for several parsecs. It is possible to do this in the lab by reproducing fundamental dimensionless variables such as the Mach number and the ratio of the cooling length to the jet span. Such experiments serve not only to simulate protostellar jets but also to rigorously test codes over wide time domains. The idea of bringing together all these different approaches to the study of protostellar jets led us to form the Jet Simulation, Experiment and Theory (JETSET) network. JETSET is a research–training network, funded by the European Union, involving ten institutions, approximately 100 scientists and 18 directly employed postdoctoral and pre-doctoral researchers. Knowledge was passed on not only through traditional one-to-one supervision and intensive international collaboration but also through a series of dedicated schools open to the wider community. The network was very successful not only in terms of its research output but also in the level of cooperation it engendered. No doubt some of our trainees will be next generation experts in this field. It was thus a great pleasure for JETSET to organise a conference in the beautiful island of Rhodes, home of the ancient Greek astronomer Hipparchus. This meeting

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served not only to showcase some of the network’s achievements but also allowed us to hear from, discuss and debate the recent findings of world-class astrophysicists in our field. Jets from young stars are of course not an isolated astrophysical phenomenon. We now know that objects as diverse as young brown dwarfs, planetary nebulae, symbiotic stars, micro-quasars, AGN, and gamma-ray bursters produce jets. Thus in a series of talks, we also put protostellar jets in context by comparing them with their often much larger brethren and also by considering the ubiquitous accretion disks that seem to be necessary for their formation. The conference itself would not have been possible without the contribution of many people. In particular we would like to thank the Scientific Organising Committee of John Bally, Sylvie Cabrit, Suzan Edwards, Sergey Lebedev, Mario Livio, Mark McCaughrean, Silvano Massaglia, Alex Raga, Kazunari Shibata, Frank Shu, and Xander Tielens for all their help in putting together an excellent programme. All meetings require local support and it is a pleasure to express our gratitude to Titos Matsakos, Perikles Rammos, Petros Tzeferacos and Nektarios Vlahakis for all their hard work. Logistical assistance was provided by Eileen Flood, Emma Whelan, Alexia and Tassos Afentoulides. Finally we would like to acknowledge our Science Project Adviser in Brussels, Dr. Renat Bilyalov, for all his assistance with running the network. July 2009

Kanaris Tsinganos Tom Ray Matthias Stute

Contents

Part I

Introductory Reviews

Astrophysical Jets : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : Mario Livio

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Jets from Young Stars : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : 11 John Bally Part II

The Star/Jet/Disk System

The Star-Jet-Disk System and Angular Momentum Transfer : : : : : : : : : : : : : : : 23 Lee Hartmann Hot Inner Winds from T Tauri Stars : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : 33 Christopher M. Johns-Krull Hot Gas in Accretion Disks and Jets: An UV View of Star Formation : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : 43 Ana I. G´omez de Castro Generalized Multipole X-Wind Model : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : 51 Subhanjoy Mohanty and Frank H. Shu Instabilities in Accretion Disks : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : 57 James M. Stone Theory of Wind-Driving Protostellar Disks : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : 67 Arieh K¨onigl Aspect Ratio Dependence in Magnetorotational Instability Shearing Box Simulations: : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : 77 Andrea Mignone, Attilio Ferrari, Gianluigi Bodo, Paola Rossi, and Fausto Cattaneo

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Advection/Diffusion of Large Scale Magnetic Field in Accretion Disks : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : 83 Richard V.E. Lovelace, David M. Rothstein, and Gennady S. Bisnovatyi-Kogan Magnetic Reconnection in Accretion Disk Systems: From BHs to YSOs : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : 89 Elisabete M. de Gouveia Dal Pino, Pamela Piovezan, Grzegorz Kowal, and Alex Lazarian Part III Jet Launching Self-Collimated Jets from Accretion Discs and Star-disc Interaction Zones : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : 99 Jonathan Ferreira Large-Scale 3D Simulations of Protostellar Jets : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : 111 Jan Staff, Kai Cai, Brian Niebergal, Rachid Ouyed, and Ralph Pudritz Magnetic Field Advection in Weakly Magnetised Viscous Resistive Accretion Disks: Numerical Simulations : : : : : : : : : : : : : : : : : : : : : : : : : : : 117 Gareth C. Murphy, Claudio Zanni, and Jonathan Ferreira Extending Analytical MHD Jet Formation Models with a Finite Disk Radius : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : 123 Matthias Stute, Kanaris Tsinganos, Nektarios Vlahakis, Titos Matsakos, and Jos´e Gracia Magnetohydrodynamic Jets from Different Magnetic Field Configurations : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : 131 Christian Fendt Resistive MHD Jet Simulations with Large Resistivity: : : : : : : : : : : : : : : : : : : : : : : 137 ˇ Miljenko Cemelji´ c, Jos´e Gracia, Nektarios Vlahakis, and Kanaris Tsinganos The X-wind Model : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : 143 Mike J. Cai Disk-Magnetosphere Interaction and Outflows: Conical Winds and Axial Jets : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : 153 Marina M. Romanova, Galina V. Ustyugova, Alexander V. Koldoba, and Richard V.E. Lovelace Simulating the Launching of YSO Jets: : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : 165 Claudio Zanni

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On the Effect of Stellar Wind Braking onto the Central Object : : : : : : : : : : : : : 173 Christophe Sauty, Noemie Globus, Zakaria Meliani, Kanaris Tsinganos, Nektarios Vlahakis, and Edo Trussoni Flaring Activity in Accretion Flows of Young Stellar Objects : : : : : : : : : : : : : : : 179 Fabio Reale Similarities of the Launching Mechanism in Protostellar/AGN Jets : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : 185 Ryoji Matsumoto Formation of Episodic Magnetically Driven Radiatively Cooled Plasma Jets in Laboratory Experiments : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : 195 Francisco Suzuki-Vidal, Sergey V. Lebedev, Andrea Ciardi, Simon N. Bland, Jeremy P. Chittenden, Gareth N. Hall, Adam Harvey-Thompson, Alberto Marocchino, Cheng Ning, Chantal Stehle, Adam Frank, Eric G. Blackman, Simon C. Bott, and Tom Ray Jets in the MHD Context : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : 205 Nektarios Vlahakis Part IV Observational Constraints on Jet Launching Jets from Embedded Protostars : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : 215 Brunella Nisini Accretion Luminosity of Embedded Protostars : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : 225 Simone Antoniucci Resolved Inner Jets from T Tauri Stars : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : 231 Francesca Bacciotti Searching for Jet Rotation Signatures in Class 0 and I Jets: : : : : : : : : : : : : : : : : : 241 Deirdre Coffey, Francesca Bacciotti, Antonio Chrysostomou, Brunetta Nisini, and Chris Davis Observational Constraints to Steady Jet Models in Young Stars : : : : : : : : : : : : 247 Sylvie Cabrit Searching for Brown Dwarf Outflows : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : 259 Emma M. Whelan, Tom Ray, Francesca Bacciotti, Sofia Randich, and Antonella Natta Protostellar Jets Driven by Intermediate- and High-Mass Protostars: An Evolutionary Scenario? : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : 267 Alessio Caratti o Garatti, Jochen Eisl¨offel, Dirk Froebrich, Brunella Nisini, and Teresa Giannini

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General Properties of Jets from Active Galactic Nuclei and Comparison with Protostellar Jets : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : 273 Silvano Massaglia Part V Jet Propagation, Stability, Interaction with the Environment, X-ray Emission The Kelvin-Helmholtz Instability in Stellar Jets: : : : : : : : : : : : : : : : : : : : : : : : : : : : : : 285 Edo Trussoni Radiative Jets from Variable Sources : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : 295 Alejandro C. Raga, Jorge Cant´o, Fabio De Colle, Alejandro Esquivel, Primoz Kajdic, Ary Rodr´ıguez- Gonz´alez, and Pablo F. Vel´azquez Position-Velocity Analysis of HH 111: Physical Structure and Dust Content : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : 305 Linda Podio, Silvia Medves, Francesca Bacciotti, Jochen Eisl¨offel, and Tom Ray Application of Tomographic Techniques to Stellar Jets : : : : : : : : : : : : : : : : : : : : : : 311 Fabio De Colle, Carlos del Burgo, and Alejandro C. Raga Measurement of Magnetic Fields in Stellar Jets : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : 317 Patrick Hartigan Jet Kinematics : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : 329 Alessio Caratti o Garatti and Jochen Eisl¨offel Synthetic Jets – from Models to Observations and Back : : : : : : : : : : : : : : : : : : : : : 341 Jos´e Gracia X-Ray Emission from Young Stellar Jets : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : 347 Manuel G¨udel, Stephen L. Skinner, Sylvie Cabrit, Jochen Eisl¨offel, Catherine Dougados, Roland Gredel, and Kevin R. Briggs The Complex Morphology of the X-Ray and Optical Emission from HH 154: The Pulsed Jet Scenario : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : 353 Rosaria Bonito, Salvatore Orlando, Giovanni Peres, Fabio Favata, and Jochen Eisl¨offel Radiative Shocks in the Context of Young Stellar Objects: A Combined Analysis from Experiments and Simulations : : : : : : : : : : : : : : : : : : 359 Chantal Stehl´e, Matthias Gonz´alez, Edouard Audit, and Thierry Lanz

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X-Ray Imaging Spectroscopy of Planetary Nebulae in the Chandra/XMM Era: New Insight into Stellar Jets : : : : : : : : : : : : : : : : : : : : : : : : : : : 367 Joel H. Kastner 3D Modeling of the 2006 Nova Outburst of RS Ophiuchi: Collimated Outflows and Jet-Like Ejections : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : 373 Salvatore Orlando, Jeremy J. Drake, and J. Martin Laming Part VI

Molecular Outflows and Turbulence Injection by Jets

Molecular Outflows: Observations : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : 381 Rafael Bachiller Driving Mechanisms for Molecular Outflows: : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : 395 Turlough P. Downes Protostellar Jet and Outflow in the Collapsing Cloud Core : : : : : : : : : : : : : : : : : : 405 Masahiro N. Machida, Shu-ichiro Inutsuka, and Tomoaki Matsumoto Outflow Driven Turbulence in Star Forming Clouds : : : : : : : : : : : : : : : : : : : : : : : : : 411 Adam Frank Jet Driven Turbulence? : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : 421 Robi Banerjee, Susanne Horn, and Ralf S. Klessen Prospects for Outflow and Jet Science with ALMA : : : : : : : : : : : : : : : : : : : : : : : : : : 429 John Richer Part VII

JETSET Early Stage Researcher Presentations

Two-component Jet Simulations: Combining Analytical and Numerical Approaches : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : 441 Titos Matsakos, Silvano Massaglia, Edo Trussoni, Kanaris Tsinganos, Nektarios Vlahakis, Christophe Sauty, and Andrea Mignone Jets from Young Stellar Objects: From MHD Simulations to Synthetic Observations : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : 447 Ovidiu Tes¸ileanu, Andrea Mignone, and Silvano Massaglia Molecular Cooling in Large Scale Simulations of Protostellar Jets: : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : 453 Jamie O’Sullivan and Max Camenzind Survival of Molecules in MHD Disk Winds : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : 459 Despina Panoglou, Sylvie Cabrit, Paolo J.V. Garcia, and Guillaume Pineau des Forˆets

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Sheared Magnetic Field and Kelvin Helmholtz Instability : : : : : : : : : : : : : : : : : : : 465 Matteo Bocchi, Hubert Baty, and Max Camenzind Jets from Class 0 Protostars: A Mid-IR Spitzer View : : : : : : : : : : : : : : : : : : : : : : : : 471 Odysseas Dionatos 0.1500 Study of the Atomic and Molecular Jets in DG Tau : : : : : : : : : : : : : : : : : : : 477 Vanessa Agra-Amboage, Catherine Dougados, and Sylvie Cabrit Velocity Resolved IR Diagnostics of Class I Jets : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : 485 Rebecca Garc´ıa L´opez, Brunella Nisini, Teresa Giannini, Jochen Eisl¨offel, Francesca Bacciotti, and Linda Podio Laboratory Astrophysics: Episodic Jet Ejections: : : : : : : : : : : : : : : : : : : : : : : : : : : : : 491 Alberto Marocchino, Jeremy P. Chittenden, Andrea Ciardi, Francisco A. Suzuki-Vidal, and Chantal Stehle Parameter Study in Disk Jet Systems : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : 497 Petros Tzeferacos, Attilio Ferrari, Andrea Mignone, Silvano Massaglia, Gianluigi Bodo, and Claudio Zanni Early Stage Development of the Jetset Database : : : : : : : : : : : : : : : : : : : : : : : : : : : : : 503 Periklis Rammos, Emma T. Whelan, Jos´e Gracia, Stephane Dudzinski, and Philippe Grange Part VIII Posters Shaping Planetary Nebulae by Jets : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : 507 Muhammad Akashi New Herbig-Haro Objects in the Gulf of Mexico : : : : : : : : : : : : : : : : : : : : : : : : : : : : : 511 Tina Armond, Bo Reipurth, and Luiz Paulo R. Vaz Launching Jets from MRI-driven Accretion Discs : : : : : : : : : : : : : : : : : : : : : : : : : : : 515 Steffen Brinkmann and Max Camenzind Properties of Jet Emitting Discs : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : 519 C´eline Combet and Jonathan Ferreira The H2 Velocity Field of Inner Knots in HH 212 : : : : : : : : : : : : : : : : : : : : : : : : : : : : : 523 Serge Correia, Hans Zinnecker, Stephen Ridgway, and Mark McCaughrean Magnetic Fields in Low-Mass Star Forming Regions: Alignment to Jets/Outflows? : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : 527 Rachel L. Curran and Antonio Chrysostomou Interacting Knots in Jets: Simulations vs. Observations : : : : : : : : : : : : : : : : : : : : : 531 Fabio De Colle and Alessio Caratti o Garatti

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Wide Field JCMT HARP-B CO(3-2) Mapping of the Serpens Cloud Core: : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : 535 Odysseas Dionatos, Brunella Nisini, Teresa Giannini, Claudio Codella, John Richer, and Mario Tafalla Numerical Simulations of Herbig Haro Objects: A Low Excitation HH Object : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : 539 Alejandro Esquivel, Alejandro C. Raga, and Fabio De Colle Soft X-rays from DG Tau: A Physical Jet Model : : : : : : : : : : : : : : : : : : : : : : : : : : : : : 543 Hans Moritz G¨unther, Sean P. Matt, and Zhi-Yun Li Multifluid Simulations of the Kelvin-Helmholtz Instability in a Weakly Ionised Plasma : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : 547 Aoife C. Jones, Mohsen Shadmehri, and Turlough P. Downes Large-scale 3D Simulations of Protostellar Jets: Long-term Stability and Jet Rotation : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : 551 Kai Cai, Jan Staff, Brian P. Niebergal, Ralph E. Pudritz, and Rachid Ouyed Extragalactic Jets with Helical Magnetic Fields : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : 555 Rony Keppens and Zakaria Meliani Jets from Collapsing Stars : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : 559 Volodymyr Kryvdyk Outflows in High-Mass Star Forming Regions : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : 563 Ana L´opez-Sepulcre, Claudio Codella, Riccardo Cesaroni, Maite T. Beltr´an, Nuria Marcellino, and Luca Moscadelli Astrophysical Jet Experiment : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : 567 Berenice Loupias, Claire Michaut, Chris D. Gregory, Emeric Falize, Jonathan Waugh, Dono Seiichi, S. Pikuz, Yasuhiro Kuramitsu, Alessandra Ravasio, Serge Bouquet, Wigen Nazarov, Youichi Sakawa, Nigel Woolsey, and Michel Koenig The Angular Momentum of Dense Clumps in Elephant Trunks : : : : : : : : : : : : 571 Veronica Lora, Alejandro C. Raga, and Alejandro Esquivel A Precessing Jet in the NGC 2264 G Outflow : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : 573 Carolyn McCoey, Paula S. Teixeira, Michel Fich, and Charles J. Lada Line Diagnostics of Large Scale Jets from Classical T Tauri Stars: The Case of DG Tau : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : 577 Fiona McGroarty, Linda Podio, Francesca Bacciotti, and Tom Ray

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Relativistic Two-Component Hydrodynamic Jets : : : : : : : : : : : : : : : : : : : : : : : : : : : : 581 Zakaria Meliani and Rony Keppens The Physical Properties of the RW Aur Bipolar Jet from HST/STIS High-Resolution Spectra: : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : 585 Stanislav Melnikov, Jochen Eisl¨offel, Francesca Bacciotti, Jens Woitas, and Tom Ray Stability of Magnetized Spine-Sheath Relativistic Jets : : : : : : : : : : : : : : : : : : : : : : : 589 Yosuke Mizuno, Philip E. Hardee, and Ken-Ichi Nishikawa Chemical Models of Hot Molecules at Shocks in Outflows : : : : : : : : : : : : : : : : : : : 593 Hideko Nomura and Tom J. Millar Survival of H2 and CO in MHD Disk Winds of Class 0, Class I and Class II Stars: : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : 595 Despina Panoglou, Paolo J.V. Garcia, Sylvie Cabrit, and Guillaume Pineau des Forˆets Three-Fluid Magnetohydrodynamics in Star Formation : : : : : : : : : : : : : : : : : : : : 597 Cecilia Pinto and Daniele Galli Physical Conditions of the Shocked Regions of Planetary Nebulae : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : 601 Angels Riera, Alejandro C. Raga, Garrelt Mellema, Alejandro Esquivel, and Pablo F. Vel´azquez The Jets of the Proto-Planetary Nebula CRL 618 : : : : : : : : : : : : : : : : : : : : : : : : : : : : 603 Angels Riera, Alejandro C. Raga, Pablo F. Vel´azquez, Sinhue Haro-Corzo, and Primoz Kajdic The Formation of Filamentary Structures in Radiative Cluster Winds : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : 605 Ary Rodr´ıguez-Gonz´ales, Alejandro Esquivel, Alejandro C. Raga, and Jorge Cant´o Hydrodynamic Modeling of Accretion Shock on CTTSs : : : : : : : : : : : : : : : : : : : : : 607 Germano G. Sacco, Constanza Argiroffi, Salvatore Orlando, Antonio Maggio, Giovanni Peres, and Fabio Reale MRI and Outflows: Angular Momentum Transport in Protoplanetary Disks : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : 611 Raquel Salmeron Analysis of the Central X-ray Source in DG Tau : : : : : : : : : : : : : : : : : : : : : : : : : : : : : 615 P. Christian Schneider and J¨urgen H.M.M. Schmitt Verification of Candidate Protostellar Outflows in GLIMPSE : : : : : : : : : : : : : : 619 Bringfried Stecklum, Alessio Caratti o Garatti, Chris Davis, Hendrik Linz, Thomas Stanke, and Hans Zinnecker

Contents

xix

Young Stellar Jets and Outflows in the Massive Star Forming Complex W5 : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : 623 Guy S. Stringfellow, John Bally, and Adam Ginsburg Water Masers and Radio Continuum Emission Tracing Thermal Radio Jets : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : 627 M.A. Trinidad Effects of Flaring Activity on Dynamics of Accretion Disks in YSOs : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : : 631 Tatiana G. Yelenina, Salvatore Orlando, Fabio Reale, Giovanni Peres, Andrea Mignone, and Titos Matsakos

Index . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .635 A Color Figures. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .641

The contributions include a picture of the first author where available. Most of the photographs have been taken by Sotiria Fotopoulou and Athina Pouri during the conference, the conference dinner, coffee breaks or the excursion. Some pictures have been kindly provided by the author himself/herself.

Contributors

Vanessa Agra-Amboage Laboratoire d’Astrophysique de l’Observatoire de Grenoble, UMR5521 du CNRS, 38041 Grenoble Cedex, France, [email protected] Muhammad Akashi Department of Physics, Technion–Israel Institute of Technology, [email protected] Simone Antoniucci INAF - Osservatorio Astronomico di Roma, Via di Frascati, 33, I-00040 Monte Porzio Catone (RM), [email protected] Constanza Argiroffi, DSFA, Universit´a di Palermo, Piazza del Parlamento, 1, Palermo, Italy Tina Armond Centro de Astrof´ısica da Universidade do Porto, Rua das Estrelas, 4150-762 Porto, Portugal, [email protected] Edouard Audit Service d’Astrophysique, DSM/IRFU/SAp, CEA/Saclay, 91191 gif-sur-yvette Cedex, France, [email protected] Francesca Bacciotti Istituto Nazionale di Astrofisica (INAF) – Osservatorio Astrofisico di Arcetri, Largo E. Fermi 5, I-50125, Firenze, Italy, [email protected] Rafael Bachiller Observatorio Astronomico Nacional (IGN), Calle Alfonso XII, 3, 28014 Madrid, Spain, [email protected] John Bally Department of Astrophysical and Planetary Sciences, University of Colorado, Boulder, CO 80309, USA, [email protected] Robi Banerjee Institute for Theoretical Astrophysics, Zentrum f¨ur Astronomie, University of Heidelberg, Albert-Ueberle-Str. 2, 69120 Heidelberg, Germany, [email protected] Hubert Baty Observatoire Astronomique de Strasbourg, 67000 Strasbourg, France, [email protected] Maite T. Beltr´an Dep. d’Astronomia i Meteorologia, Facultat de F´ısica, UB Barcelona, Spain

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Contributors

Gennady S. Bisnovatyi-Kogan Space Research Institute, Russian Academy of Sciences, Moscow, Russia, [email protected] Eric G. Blackman Department of Physics and Astronomy, University of Rochester, Rochester, NY, USA Simon N. Bland Imperial College London, The Blackett Laboratory, Plasma Physics Group, London SW7 2BW, UK Matteo Bocchi ZAH - Landessternwarte, Koenigstuhl 10, 69117 Heidelberg, Germany, [email protected] Gianluigi Bodo Osservatorio Astronomico di Torino, Viale Osservatorio 20, 10025 Pino Torinese, Italy Rosaria Bonito INAF-Osservatorio di Palermo-COMETA, Italy, [email protected] Simon C. Bott Center for Energy Research, University of California, San Diego, CA 92093-0417, USA Serge Bouquet LUTH, Observatoire de Paris, CNRS, Universite Paris-Diderot, France and D´epartement de Physique Th´eorique et Appliqu´ee, CEA-DIF, France Kevin Briggs Institute of Astronomy, ETH Z¨urich, 8092 Z¨urich, Switzerland Steffen Brinkmann Landessternwarte (ZAH), Universit¨at Heidelberg, K¨onigstuhl 12, 69117 Heidelberg, Germany, [email protected] Sylvie Cabrit LERMA, Observatoire de Paris, 61 Av. de lObservatoire, 75014 Paris, [email protected] Kai Cai Department of Physics, McMaster University and Astronomy, ABB-241, 1280 Main St. W, Hamilton, ON, Canada, L8S 4M1 and Department of Physics and Astronomy, University of Kentucky, Lexington, KY 40506, USA, [email protected] Mike J. Cai Academia Sinica, Institute of Astronomy and Astrophysics, Taiwan, [email protected] Max Camenzind Landessternwarte (ZAH), University of Heidelberg, K¨onigstuhl 12, 69117 Heidelberg, Germany, [email protected] Jorge Cant´o Instituto de Astronom´ıa, Universidad Nacional Aut´onoma de M´exico, Ap. 75-264, 04510 D.F., M´exico, [email protected] Alessio Caratti o Garatti Th¨uringer Landessternwarte Tautenburg, Sternwarte 5, 07778 Tautenburg, Germany, [email protected] Fausto Cattaneo Department of Astronomy and Astrophysics, The University of Chicago, 5640 S. Ellis ave., Chicago, IL 60637, USA

Contributors

xxiii

ˇ Miljenko Cemelji´ c TIARA, Academia Sinica, National Tsing Hua University, No. 101, Sec. 2, Kuang Fu Rd., Hsinchu 30013, Taiwan, [email protected] Riccardo Cesaroni INAF - Osservatorio Astrofisico di Arcetri, Largo E. Fermi, 5, 50125 Firenze, Italy Jeremy P. Chittenden Imperial College, Prince Consort Road, London SW7 2BW, UK, [email protected] Antonio Chrysostomou University of Hertfordshire, Hatfield, UK, [email protected] Andrea Ciardi Observatoire de Paris, LUTH, Meudon 92195, France Claudio Codella INF – Istituto di Radioastronomia, Sezione di Firenze, Firenze, Italy, [email protected] Deirdre Coffey The Dublin Institute for Advanced Studies, Dublin 2, Ireland, [email protected] C´eline Combet Department of Physics and Astronomy, University of Leicester, LE17RH, Leicester, UK, [email protected] Serge Correia AIP, An der Sternwarte 16, 14482 Potsdam, Germany, [email protected] Rachel L. Curran Osservatorio Astronomico di Palermo, Piazza del Parlamento, 1, 90134 Palermo, Italy, [email protected] Chris Davis Joint Astronomy Centre, Hilo, Hawaii, USA, [email protected] Fabio De Colle Dublin Institute for Advanced Studies, 31 Fitzwilliam Place, Dublin 2, Ireland, [email protected] Elisabete M. de Gouveia Dal Pino IAG-USP, Kua do Matao 1226, Cidade Universitaric Sao Paulo, Brazil, [email protected] Carlos del Burgo Dublin Institute for Advanced Studies, 31 Fitzwilliam Place, Dublin 2, Ireland, [email protected] Odysseas Dionatos INAF - Osservatorio Astronomico di Roma, Italy, [email protected] Catherine Dougados Laboratoire d’Astrophysique de l’Observatoire de Grenoble, UMR5521 du CNRS, 38041 Grenoble Cedex, France, [email protected] Turlough P. Downes Dublin City University and Dublin Institute for Advanced Studies, Dublin 2, Ireland, [email protected] Jeremy J. Drake Harvard-Smithsonian Center for Astrophysics, 60 Garden Street, Cambridge, MA 02138, USA, [email protected]

xxiv

Contributors

Stephane Dudzinski Dublin Institute for Advanced Studies, Dublin 2, Ireland Jochen Eisl¨offel Th¨uringer Landessternwarte Tautenburg, Sternwarte 5, 07778 Tautenburg, Germany, [email protected] Alejandro Esquivel Instituto de Ciencias Nucleares, Universidad Nacional Aut´onoma de M´exico, A. Postal 70-543, M´exico D.F. 04510, M´exico, [email protected] Emeric Falize LUTH, Observatoire de Paris, CNRS, Universite Paris-Diderot, France Fabio Favata ESA, Community Coordination and Planning Office, Paris Christian Fendt Max Planck Institute for Astronomy, K¨onigstuhl 17, 69117 Heidelberg, Germany, [email protected] Attilio Ferrari DFG, University of Turin, via P. Giuria 1, 10125 Torino, Italy Jonathan Ferreira Laboratoire d’Astrophysique de Grenoble, Grenoble, France, [email protected] Michel Fich University of Waterloo and University of Western Ontario Adam Frank Dept. of Physics and Astronomy, University of Rochester, Rochester, NY 14627, [email protected] Dirk Froebrich Centre for Astrophysics and Planetary Science, University of Kent, Canterbury, CT2 7NH, UK, [email protected] Daniele Galli INAF-Osservatorio Astrofisico di Arcetri, Largo E. Fermi 5, 50125 Firenze, Italy, [email protected] Paulo J. V. Garcia Faculdade de Engenharia, Universidade do Porto, Portugal, [email protected] Teresa Giannini INAF-Osservatorio Astronomico di Roma, Via Frascati, 33, 00040 Monte Porzio Catone, Italy, [email protected] Adam Ginsburg Center for Astrophysics and Space Astronomy, Department of Astrophysical and Planetary Sciences, University of Colorado, Boulder, C. USA Noemie Globus Observatoire de Paris, LUTH, 92190 Meudon, France Ana I. G´omez de Castro Astronom´ya y Geodesia, Fac. de CC Matem´aticas, Universidad Complutense de Madrid, 28040 Madrid, Spain, [email protected] Matthias Gonz´alez Instituto de Fusi´on Nuclear Universidad Polit´ecnica de Madrid - ETSII, calle Jos´e Guti´errez Abascal 2 28006 Madrid, Spain, [email protected] Jos´e Gracia School of Cosmic Physics, Dublin Institute of Advanced Studies, 31 Fitzwilliam Place, Dublin 2, Ireland, [email protected] Philippe Grange Dublin Institute for Advanced Studies, Dublin 2, Ireland

Contributors

xxv

Roland Gredel Max-Planck-Institute for Astronomy, 69117 Heidelberg, Germany, [email protected] Chris D. Gregory LULI, Ecole Polytechnique, 91128 Palaiseau, Cedax, France ¨ Manuel Gudel Institute of Astronomy, ETH Z¨urich, 8092 Z¨urich, Switzerland, [email protected] ¨ Hans Moritz Gunther Hamburger Sternwarte, Gojenbersgweg 112, 21029 Hamburg, Germany, [email protected] Gareth N. Hall Imperial College London, The Blackett Laboratory, Plasma Physics Group, London SW7 2BW, UK Philip E. Hardee Department of Physics and Astronomy, University of Alabama, Tuscaloosa, AL 35487, USA Sinhue Haro-Corzo Instituto de Ciencias Nucleares, Universidad Nacional Aut´onoma de M´exico, M´exico, [email protected] Patrick Hartigan Physics and Astronomy Department, Rice University, 6100 S. Main, Houston, TX 77005, USA, [email protected] Lee Hartmann University of Michigan, 830 Dennison, 500 Church St., Ann Arbor, MI 48105, USA, [email protected] Adam Harvey-Thompson Imperial College London, The Blackett Laboratory, Plasma Physics Group, London SW7 2BW, UK Susanne Horn Institute for Theoretical Astrophysics, Zentrum f¨ur Astronomie, University of Heidelberg, Albert-Ueberle-Str. 2, 69120 Heidelberg, Germany, [email protected] Shu-ichiro Inutsuka Department of Physics, Graduate School of Science, Kyoto University, Sakyo-ku, Kyoto 606-8502, Japan, [email protected] Christopher M. Johns-Krull Department of Physics and Astronomy, Rice University, Houston, TX 77005, USA, [email protected] Aoife C. Jones School of Mathematical Sciences, DCU, Ireland, [email protected] Primoz Kajdic Institute de Geoffsica, Universidad Nacional Aut´onoma de M´exico, 04510 D. F., M´exico, [email protected], [email protected] Joel H. Kastner Laboratoire d’Astrophysique de Grenoble Universit´e Joseph Fourier – CNRS, BP 53, 38041 Grenoble Cedex, France, [email protected] Rony Keppens Centre for Plasma-Astrophysics, K.U. Leuven, Belgium, [email protected] Ralf S. Klessen Institute for Theoretical Astrophysics, Zentrum f¨ur Astronomie, University of Heidelberg, Albert-Ueberle-Str. 2, 69120 Heidelberg, Germany, [email protected]

xxvi

Contributors

Michel Koenig LULI, Ecole Polytechnique, 91128 Palaiseau, Cedax, France Alexander V. Koldoba Institute for Mathematical Modeling RAS, Moscow 125047, Russia, [email protected] Arieh K¨onigl University of Chicago, Chicago, IL 60637, USA, [email protected] Grzegorz Kowal University of Wisconsin-Madison, 716 Langdon St., Madison, WI 53706-1481, USA Volodymyr Kryvdyk Department of Astronomy, Faculty of Physics, Taras Shevchenko Kyiv National University, av. Glushkova 2/1, Kyiv 03022, Ukraine, [email protected] Yasuhiro Kuramitsu Graduate School of Engineering, Osaka University, Osaka 565-0871, Japan Charles J. Lada Harvard-Smithsonian Center for Astrophysics now ESO-Garching, Cambridge, MA 02138, USA J. Martin Laming Space Science Division, Naval Research Laboratory, Code 7674L, Washington DC 20375, USA, [email protected] Thierry Lanz Department of Astronomy, University of Maryland, College Park, MD 20742 USA and LERMA, Observatoire de Paris, CNRS and UPMC, 5 Place Jules Janssen 92195 Meudon, France, [email protected] Alex Lazarian University of Wisconsin-Madison, 716 Langdon St., Madison, WI 53706-1481, USA Sergey V. Lebedev Imperial College London, The Blackett Laboratory, Plasma Physics Group, London SW7 2BW, UK Zhi-Yun Li Department of Astronomy, University of Virginia, P.O. Box 400325, Charlottesville, VA 22904, USA, [email protected] Hendrik Linz MPIA, K¨onigstuhl, Heidelberg, Germany Mario Livio Space Telescope Science Institute, 3700 San Martin Drive, Baltimore, MD 21218, USA, [email protected] Rebecca Garc´ıa L´opez INAF-Osservatorio Astronomico di Roma, Via Frascati, 33, 00040 Monte Porzio Catone, Italy, [email protected] Ana L´opez-Sepulcre INAF - Osservatorio Astrofisico di Arcetri, Largo E. Fermi, 5, 50125 Firenze, Italy, [email protected] Veronica Lora Instituto de Astronom´ıa, Universidad Nacional Aut´onoma de M´exico, Ap. 70-264, 04510 D.F., M´exico, [email protected]

Contributors

xxvii

Berenice Loupias LULI, Ecole Polytechnique, 91128 Palaiseau, Cedax, France Richard V.E. Lovelace Department of Astronomy, Cornell University, Ithaca, NY 14853, USA, [email protected] Masahiro N. Machida Department of Physics, Graduate School of Science, Kyoto University, Sakyo-ku, Kyoto 606-8502, Japan, [email protected] Antonio Maggio INAF-Osservatorio Astronomico di Palermo, Piazza del Parlamento, 1, Palermo, Italy N. Marcelino DAMIR, CSIC, Madrid, Spain Alberto Marocchino Imperial College, Prince Consort Road, London SW7 2BW, UK, [email protected] Silvano Massaglia Dipartimento di Fisica Generale dell’Universit´a, Via Pietro Giuria 1, 10125 Torino, Italy, [email protected] Titos Matsakos Dipartimento di Fisica Generale, Universit´a degli Studi di Torino, via Pietro Giuria 1, 10125 Torino, Italy, [email protected] Ryoji Matsumoto Department of Physics, Graduate School of Science, Chiba University, 1-33 Yayoi-Cho, Inage-ku, Chiba 263-8522, Japan, [email protected] Tomoaki Matsumoto Faculty of Humanity and Environment, Hosei University, Fujimi, Chiyoda-ku, Tokyo 102-8160, Japan, [email protected] Sean P. Matt Department of Astronomy, University of Virginia, P.O. Box 400325, Charlottesville, VA 22904, USA, [email protected] Mark McCaughrean Astrophysics Group, School of Physics, University of Exeter, Exeter EX4 4QL, UK, [email protected] Carolyn McCoey University of Waterloo and University of Western Ontario Fiona McGroarty Department of Physics and Astronomy, Queen’s University Belfast, Belfast BT7 1NN, Northern Ireland, [email protected] Silvia Medves Universit`a di Pisa, Dipartimento di Fisica, Largo B. Pontecorvo 3, 56127 Pisa, Italy, [email protected] Zakaria Meliani Centre for Plasma-Astrophysics, K.U. Leuven, Belgium, [email protected] Garrelt Mellema Stockholm University, Sweden, [email protected] Stanislav Melnikov Th¨uringer Landessternwarte Tautenburg (TLS), Sternwarte 5, 07778 Tautenburg, Germany Claire Michaut LUTH, Observatoire de Paris, CNRS, Universite Paris-Diderot, France

xxviii

Contributors

Andrea Mignone Universit´a degli Studi di Torino, via P. Giuria 1, I-10125 Turin, Italy, [email protected] and Osservatorio Astronomico di Torino, via Osservatorio 20, I-10025 Pino Torinese (TO), Italy, [email protected] Tom J. Millar ARC, School of Mathematics and Physics, Queen’s University Belfast, UK Yosuke Mizuno Center for Space Plasma and Aeronomic Research, The University of Alabama in Huntsville, 320 Sparkman Drive, NSSTC 2104, Huntsville, AL 35805, USA, [email protected] Subhanjoy Mohanty Imperial College London, South Kensington Campus, London SW72A7, UK, [email protected] L. Moscadelli INAF - Osservatorio Astrofisico di Arcetri, Largo E. Fermi, 5, 50125 Firenze, Italy Gareth C. Murphy Dublin Institute for Advanced Studies, 31 Fitzwilliam Place, Dublin 2, Ireland, [email protected] Antonella Natta Osservatorio Astrofisico di Arcetri, Italia Wigen Nazarov University of St. Andrews, School of Chemistry, UK Brian P. Niebergal Department of Physics and Astronomy, University of Calgary, 2500 University Drive NW, Calgary, AB, Canada T2N 1N4 Cheng Ning Imperial College London, The Blackett Laboratory, Plasma Physics Group, London SW7 2BW, UK Ken-Ichi Nishikawa CSPAR, The University of Alabama in Huntsville, Huntsville, AL 35805, USA Brunetta Nisini INAF-Osservatorio Astronomico di Roma, Via di Frascati, 33, I-00040 Monte Porzio Catone (RM), [email protected] Hideko Nomura Department of Astronomy, Kyoto University, Japan, [email protected] Salvatore Orlando INAF - Osservatorio Astronomico di Palermo “G.S. Vaiana”, Piazza del Parlamento 1, 90134 Palermo, Italy; Consorzio COMETA, via Santa Sofia 64, 95123 Catania, Italy, [email protected] Jamie O’Sullivan Landessternwarte (ZAH), Heidelberg, Germany, [email protected] Despina Panoglou Faculdade de Ciˆencias, Universidade do Porto, Portugal Universit´e Pierre et Marie Curie–Paris 6, France, [email protected] Rachid Ouyed Department of Physics and Astronomy, University of Calgary, SB 605, 2500 University Drive NW, Calgary, AB, Canada T2N 1N4, [email protected]

Contributors

xxix

S. Pikuz Multicharged Ions Spectra Data Center of VNIIFRTI, Mendeleevo, Russia Guillaume Pineau des Forˆets Institut d’Astrophysique Spatiale, Orsay, France, [email protected] Cecilia Pinto Dipartimento di Astronomia e Scienza dello Spazio, Universit´a di Firenze, Largo E. Fermi 5, 50125 Firenze, Italy, [email protected] Pamela Piovezan MPA, Garching Germany, [email protected] Giovanni Peres DSFA, Universit´a di Palermo, Piazza del Parlamento, 1, Palermo, Italy Linda Podio School of Cosmic Physics, Dublin Institute for Advanced Studies, 31 Fitzwilliam Place, Dublin 2, Ireland, [email protected] Ralph E. Pudritz Department of Physics and Astronomy, McMaster University, 1280 Main St.W., Hamilton, ON, Canada, L8S 4M1, [email protected] Alejandro C. Raga ICN, Universidad Nacional Aut´onoma de M´exico, Ap. 75-543, 04510 D. F., M´exico, [email protected] Periklis Rammos Dublin Institute for Advanced Studies, 31 Fitzwilliam Place, Dublin 2, Ireland, [email protected] Sofia Randich Osservatorio Astrofisico di Arcetri, Italia Alessandra Ravasio LULI, Ecole Polytechnique, 91128 Palaiseau, Cedax, France Tom Ray School of Cosmic Physics, Dublin Institute for Advanced Studies, 31 Fitzwilliam Place, Dublin 2, Ireland, [email protected] Fabio Reale Dipartimento di Scienze Fisiche and Astronomiche, Universit`a di Palermo, Sezione di Astronomia, Piazza del Parlamento 1, 90134 Palermo, Italy and INAF - Osservatorio Astronomico di Palermo “Giuseppe S. Vaiana”, Piazza del Parlamento 1, I-90134 Palermo, Italy, [email protected] Bo Reipurth Institute for Astronomy, University of Hawaii, 640 N. Aohoku Place, Hilo, HI 96720, USA, [email protected] John Richer Cavendish Laboratory, JJ Thomson Avenue, Cambridge, UK, [email protected] Stephen Ridgway NOAO, PO Box 26732, Tucson, AZ 8526, USA, [email protected] Angels Riera Departament de F´ısica i Enginyeria Nuclear, Universitat Polit‘ecnica de Catalunya, Spain, [email protected] Ary Rodriguez-Gonz´alez ICN, Universidad Nacional Aut´onoma de M´exico, Ap. 75-543, 04510 D. F., M´exico, [email protected]

xxx

Contributors

Marina M. Romanova Department of Astronomy, Cornell University, Ithaca, NY 14853, USA, [email protected] Paola Rossi INAF/Osservatorio Astronomico di Torino, via Osservatorio 20, 10025 Pino Torinese, Italy David M. Rothstein Department of Astronomy, Cornell University, Ithaca, NY 14853, USA, [email protected] Germano G. Sacco Consorzio COMETA, Via S. Sofia, 64, 95123, Catania, Italy, [email protected] Youichi Sakawa Graduate School of Engineering, Osaka University, Osaka 565-0871, Japan Raquel Salmeron Research School of Astronomy and Astrophysics and Research School of Earth Sciences, The Australian National University, Canberra, Australia, [email protected] Christophe Sauty Observatoire de Paris, LUTH, 92190 Meudon, France, [email protected] ¨ Jurgen H.M.M. Schmitt Hamburger Sternwarte, 21029 Hamburg, Germany, [email protected] P. Christian Schneider Hamburger Sternwarte, Gojenbergsweg 112, 21029 Hamburg, Germany, [email protected] Dono Seiichi Graduate School of Engineering, Osaka University, Osaka 565-0871, Japan Frank H. Shu University of California at San Diego, LaJollal, CA 92093, USA, [email protected] Stephen L. Skinner CASA, University of Colorado, Boulder, CO 80309, USA, [email protected] Jan Staff Department of Physics, Purdue University, 525 Northwestern Avenue West Lafayette, IN 47907-2036, USA and Department of Physics and Astronomy, Louisiana State University, 202 Nicholson Hall, Tower Dr., Baton Rouge, LA 70803-4001, USA, [email protected] Thomas Stanke ESO, Garching, Germany Chantal Stehl´e LERMA, Observatoire de Paris, CNRS and UPMC, 5 Place Jules Janssen, 92195 Meudon, France, [email protected] Bringfried Stecklum TLS Tautenburg, Sternwarte 5, 07778 Tautenburg, Germany, [email protected] James M. Stone Department of Astrophysical Science, Princeton University, Princeton, NJ 08544, USA, [email protected]

Contributors

xxxi

Guy S. Stringfellow Center for Astrophysics and Space Astronomy, Department of Astrophysical and Planetary Sciences, University of Colorado, Boulder, Colorado, [email protected] Matthias Stute IASA and Section of Astrophysics, Astronomy and Mechanics, Department of Physics, University of Athens, Panepistimiopolis, 157 84 Zografos, Athens, Greece, [email protected] Francisco Suzuki-Vidal Imperial College London, The Blackett Laboratory, Plasma Physics Group, London SW7 2BW, UK, [email protected] Mario Tafalla Observatorio Astronmico Nacional, Madrid, Spain, [email protected] Paula S. Teixeira Harvard-Smithsonian Center for Astrophysics now ESO-Garching, Cambridge, MA 02138, USA, [email protected] Ovidiu Tes¸ileanu Universit´a degli Studi di Torino, via P. Giuria 1, I-10125 Turin, Italy and Research Centre for Atomic Physics and Astrophysics, RO-077125, Bucharest, Romania, [email protected] M.A. Trinidad Department of Astronomy, University of Guanajuato, Guanajuato, Mexico, [email protected] Edo Trussoni INAF/Osservatorio Astronomico di Torino, via Osservatorio 20, 10025 Pino Torinese, Italy, [email protected] Kanaris Tsinganos IASA and Section of Astrophysics, Astronomy and Mechanics, Department of Physics, University of Athens, Panepistimiopolis, 157 84 Zografos, Athens, Greece, [email protected] Petros Tzeferacos Dipartimento di Fisica Generale, Universit´a degli Studi di Torino, Via Giuria 1, 10125 Torino, Italy, [email protected] Galina V. Ustyugova Keldysh Institute of the Applied Mathematics RAS, Moscow 125047, Russia, [email protected] Luiz Paulo R. Vaz Depto. de F´ısica, ICEx, UFMG, CP 702, 30123-970 Belo Horizonte, MG, Brazil, [email protected] Pablo F. Vel´azquez ICN, Universidad Nacional Aut´onoma de M´exico, Ap. 75-543, 04510 D. F., M´exico, [email protected] Nektarios Vlahakis IASA and Section of Astrophysics, Astronomy and Mechanics, Department of Physics, University of Athens, Panepistimiopolis, 157 84 Zografos, Athens, Greece, [email protected] Jonathan Waugh Departement of Physics, University of York, Heslington, York Y0105DD, UK Emma T. Whelan Dublin Institute for Advanced Studies, [email protected]

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Contributors

Jens Woitas Th¨uringer Landessternwarte Tautenburg (TLS), Sternwarte 5, 07778 Tautenburg, Germany Nigel Woolsey Department of Physics, University of York, Heslington, York Y0105DD, UK Tatiana G. Yelenina, INAF Osservatorio Astronomico di Palermo, Piazza del Parlamento 1, 90134 Palermo, Italy, [email protected] Claudio Zanni INAF Osservatorio Astronomico di Torino, Via dell’Osservatorio 20, 10025 Pino Torinese, Italy, [email protected] Hans Zinnecker AIP, An der Sternwarte 16, 14482 Potsdam, Germany, [email protected]

Part I

Introductory Reviews

Astrophysical Jets Mario Livio

1 Introduction Highly collimated jets are observed in many classes of astrophysical objects, ranging from active galactic nuclei (AGN) to young stellar objects (YSOs). In the present paper, like in a couple of previous reviews [1,2], I will make the assumption that the jet formation mechanism, namely, the mechanism for acceleration and collimation, is the same in most if not all of the different classes of objects which exhibit jets (see [2] for details and relevant references). Adopting a mostly phenomenological approach, I will then attempt to determine to which constraints such an assumption can lead. However, with the discovery of new classes of objects which produce jets (see Sect. 2 below) and with recent developments in theoretical work, the constraints become more meaningful. It should be noted right away that the emission

M. Livio () Space Telescope Science Institute, 3700 San Martin Drive, Baltimore, MD 21218, USA e-mail: [email protected]

K. Tsinganos et al. (eds.), Protostellar Jets in Context, Astrophysics and Space Science Proceedings, DOI 10.1007/978-3-642-00576-3 1, c Springer-Verlag Berlin Heidelberg 2009 

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mechanisms which render jets observable in the different classes of objects, are very different in objects like, for example, YSOs and AGN. Here, I therefore concentrate only on acceleration and collimation.

2 The Disk-Jet Connection In this section, I present all the classes of objects which exhibit jets, and discuss some aspects of the observational evidence for a connection between accretion disks and jets.

2.1 Systems Producing Collimated Jets In Table 1, I give a list of all the types of objects in which collimated jets have been (at least tentatively) observed, and the nature of the physical system involved. A few of these objects (symbiotic stars; low mass x-ray binaries with a neutron star accretor) require a little explanation, one class (planetary nebulae), has not yet routinely made it into the jet literature, another class (supersoft x-ray sources) is relatively new, and two classes (recurrent novae and pulsars) are still only tentative. The evidence for jets in gamma-ray bursts (GRBs) is indirect, but quite compelling. Systems which have been traditionally associated with jets are: many AGN and YSOs and some massive x-ray binaries (HMXBs), such as SS 433, Cyg X–3, and the Galactic center source 1E140.7–2942. More recently, black hole x-ray transients have been added as a class, as were GRBs. So far, the only low mass x-ray binary (LMXB) with a neutron star accretor in which a jet has been observed is Cir X–1, and even in that case it is not clear how collimated the flow really is.

Table 1 Systems which exhibit collimated jets Object Physical System Stellar Young Stellar Objects Accreting young star Massive X-Ray Binaries Accreting neutron star or black hole Black Hole X-Ray Transients Accreting black hole Low Mass X-Ray Binaries Accreting neutron star Symbiotic Stars Accreting white dwarf Planetary Nebulae Nuclei Accreting nucleus (or “interacting winds”) Supersoft X-Ray Sources Accreting white dwarf Recurrent Novae(?) Accreting white dwarf Pulsars(?) Spinning neutron star(?) Extragalactic Active Galactic Nuclei Accreting supermassive black hole Gamma-Ray Bursts Accreting black hole

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The only symbiotic system in which a jet has been unambiguously observed, both in the optical and in the radio, is R Aqr. Spectroscopic evidence suggests the possible presence of a jet also in MWC 560, and optical images suggest the existence of a jet in He 2–104. I now turn to the new classes of objects which should, in my opinion, from now on routinely be included in any discussion of jets. In planetary nebulae (PNe), jets have now been directly observed (in the optical) in NGC 6543. Other systems in which the data are less conclusive include K1–2, M1–92, He 2–104, and NGC 7009. In addition, several “point-symmetric” PNe have been interpreted as resulting from precessing and wobbling jets. A relatively new, exciting addition to the classes of objects which produce jets are the supersoft x-ray sources (SSS). These are luminous (Lbol  1037 –1038 erg s1 ) objects, with a characteristic radiation temperature of (1–10)105 K, in which probably a white dwarf accretes mass from a subgiant companion at such a high rate that it burns hydrogen steadily. Recent spectroscopic observations of the LMC source RX J0513.9-6951 reveal what is probably a bipolar collimated outflow with a velocity of 3,800 km s1 , through the presence of blue- and red-shifted satellite emission features to the optical Hydrogen and Helium recombination lines. Similar features corresponding to a projected velocity of 850 km s1 have now been observed also in the SSS RX J0019.8C2156. In the latter case satellite lines to P , Pˇ, and Br  have also been observed. Such “jet lines” may also be present in CAL 83. The similarity of the spectral features corresponding to the outflow to those observed in SS 433 is striking. The latest class of systems which observations indicate may produce jets is that of recurrent novae (RNe). The hydrogen emission lines in the recurrent nova U Sco show a triple structure, with red- and blue-shifted satellite peaks corresponding to line-of-sight velocities of ˙1800 km s1 . These satellite peaks could correspond to an outflow with an opening angle of 6ı . Similar “jet” satellite lines were seen in the (probably recurrent) nova Nova Oph 1998. There is increasing evidence, in the form of kinks in the afterglow light curve, and in polarization, that GRBs are collimated into narrow jets. Finally, intriguing x-ray images of the Crab and Vela pulsars show features that may be interpreted as jets (1999; e.g. Chandra X-ray Observatory Center press release, NASA PR 99-109). An examination of Table 1 reveals that all the objects which exhibit jets (with the possible exception of the Crab pulsar) contain accreting central objects (some models for jets in PNe and YSOs do not involve accretion, see Sect. 2.2, but others do); this leads us naturally to the question in the next section.

2.2 Does the Formation of Jets Require an Accretion Disk? Clearly a complete answer to this question is difficult, since it requires both a demonstration that disks can produce jets in all the different classes of objects and

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that other mechanisms cannot produce them. Since I have adopted a phenomenological approach, I will rather attempt to answer the simpler question: has an accretion disk been observed in all of the classes of objects which produce jets? In the case of YSOs, the answer is clearly: yes, with the most dramatic manifestation being the disks and jets recently observed in the Herbig-Haro object HH 30, in DG Tau B, and in Haro 6–5B. Similarly, disks have unambiguously been observed in all the classes of x-ray binaries (HMXBs, LMXBs, SSS, and black hole x-ray transients). Furthermore, in the case of the black hole x-ray transients, it has been shown that most likely the IR and radio emitting plasma is ejected from the inner disk (see also Sect. 19). The situation with AGN is somewhat more frustrating. Although almost all of the researchers in this field agree that there are accretion disks in AGN, the evidence is somewhat circumstantial, and every now and then there are even attempts to cast doubt on their existence. Here I would merely like to mention a few recent pieces of evidence for the presence of disks in AGN, which are fairly convincing: 1. The iron K˛ line in MCG-6-30-15, which is consistent with emission from a disk and a similar line from NCG-5-23-16 and other AGNs. 2. The fact that the fit to the double peaked Balmer lines in 3C 390.3 with an accretion disk, and the superluminal motion observed in the same source, give an inclination angle for the disk and the jet which shows that the jet is exactly perpendicular to the disk. 3. The dust torus observed in NGC 4261, which is remarkably consistent with AGN unification schemes containing an accretion disk. 4. The warped subparsec-scale molecular disk observed in the maser emitting LINER NGC 4258. 5. The fact that velocity-delay maps of optical and ultraviolet emission lines in objects like NGC 5458 and NGC 4151 appear much more consistent with disk kinematics than with spherical freefall. Incidentally, for some time there has been a question whether the double-peaked Balmer emission lines observed in some (mostly radio-loud) AGN originate in an accretion disk, or in two line emitting cones (formed by two-sided jets. However, [3] have shown that at least in the case of 3C 390.3 the double-peaked lines cannot be produced in a two-sided jet, because the emitting region on the receding jet is expected to be obscured from view by the accretion disk, which is optically thick up to radii of R  1018 cm .MBH =108 Mˇ /. The fact that the red wing of a line produced in a two-sided jet may be obscured from view by the accretion disk is well known from YSOs (see e.g. [OI]  6,300 profiles for T Tauri stars). For GRBs, obviously, there is no direct indication for the presence of a disk. Nevertheless, the formation of such a disk around the black hole that results from the collapse of massive stars is very likely. In the case of PNe, until recently, only theoretical arguments for the presence of disks in these systems existed. These relied on one hand on the fact that following a common envelope phase (which is required, to form the observed close binary nuclei), the somewhat bloated secondary companion is likely to fill its Roche lobe.

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On the other hand, in binary systems in which the secondary star accretes from the wind of an AGB star primary, a disk can form around the secondary. Large dust disks have been observed in the optical and infrared, and in the optical in the “Red Rectangle” nebula. A word of caution is needed in relation to YSOs (and PNe). While the presence of accretion disks in the former systems is unquestionable, some models for the collimation of jets in these systems (and indeed in PNe), suggest that refraction through oblique shocks is sufficient to produce highly collimated jets, without an active role for the accretion process. In these models, a fast and dilute wind interacts with a slowly moving or stationary torus in the equatorial plane, and collimation is achieved via refraction through the oblique shocks in the interaction region. Further work on these models will be required, to establish whether they can indeed produce long-lived, highly collimated jets. Here, however, I will not discuss such models further, since, as explained in the introduction, I am interested in a universal model for all the classes of objects, while this mechanism (“shock focused inertial confinement”) requires the presence of a torus which is not expected to exist at least in some of the systems. To conclude this section therefore, my answer to the question: do jets require an accretion disk is: probably yes, although inertial collimation and the processes operating in pulsars certainly deserve more attention.

3 Clues on the Jet Formation Mechanism Since we have determined that the formation of jets most probably requires the presence of an accretion disk, we can now examine some of the properties of jets, in an attempt to determine which basic ingredients must be associated with the accretion disk, for the acceleration and collimation mechanisms to operate.

The Jet Origin An important conclusion can be drawn from the observed jet velocities. In Table 2, I give examples for the ratio Vjet =Vescape (where Vescape is the escape velocity from the central object) for the different classes of objects. It is immediately clear that in all cases the jet velocity is of the order of the escape velocity from the central object (or the Keplerian speed near its surface). This immediately indicates that most of the outflow originates from the center of the accretion disk, from the vicinity of the central object. This general inference has received impressive observational confirmation by the HST images of HH 30 and the DG Tau B, which show clearly the jet emanating from the center of the accretion disk. Evidence for the fact that jets originate in the inner disk is provided also by multiwavelength observations of the black hole x-ray transient (“microquasar”) system

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M. Livio Table 2 The ratio of jet velocity to the escape velocity from the central object Object Vjet =Vescape Example Young Stellar Objects 1 HH 30, HH 34 Vjet  100–350 km s1 < Active Galactic Nuclei 1 Radio sources,   10 > M87,   3 Gamma-Ray Bursts 1   300 X-Ray Binaries 1 SS 433, Cyg X–3 Vjet  0:26c Black Hole X-Ray 1 GRO 1655-40, GRS 1915+105 > Transients Vjet  0:9c Planetary Nebulae

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FLIERS, Ansae, hot winds V  200–1000 km s1 RX J0513.9-6951, RX J019.8+2156 Vjet .projected/  3800 km s1 , Vjet .projected/  850 km s1 U Sco Vjet .projected/  1800 km s1

GRS 1915+105. These observations have demonstrated convincingly that there exists a one-to-one correspondence between x-ray and IR (and probably radio) flares, and the constant time delay between the x-ray/IR peaks indicates that these are triggered by the same event. This, in turn, implies that initially, the emitting regions of the x-ray and IR are in close proximity to each other. The fact that subsequently the IR and x-ray emission appear to decouple suggests that the emitting regions separate significantly at later times. A picture in which the inner disk ejects a relativistic plasma which produces the IR and radio flares by synchrotron emission is consistent with the existing data (especially since GRS 1915+105 has actually been observed to eject relativistic blobs which produce synchrotron emission, although not all the details have been clarified). A thorough examination of a variety of acceleration and collimation mechanisms has led many to the conclusion that the only mechanism that could work involves hydromagnetic acceleration and collimation. At least some fraction of the magnetic flux has to be in open field lines, which form an angle with the disk surface (see [2] for a review and [4, 5] for some important details).

References 1. Livio, M.: Astrophysical jets: a phenomenological examination of acceleration and collimation. Phys. Rep. 311, 225–245 (1999) 2. Livio, M.: Astrophysical jets. In Cosmic Explosions, S. S. Holt & W. W. Zhang (eds.), pp. 275– 297. American Institute of Physics, Melville (2000) 3. Livio, M., Xu, C.: On the Observational Evidence for Accretion Disks in Active Galactic Nuclei. Astrophys. J. Lett. 478, L63–L65 (1997)

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4. Ogilvie, G. I., Livio, M.: On the Difficulty of Launching an Outflow from an Accretion Disk. Astrophys. J. 499, 329–339 (1998) 5. Ogilvie, G. I., Livio, M.: Launching of Jets and the Vertical Structure of Accretion Disks. Astrophys. J. 553, 158–173 (2001)

Part II

The Star/Jet/Disk System

Jets from Young Stars John Bally

Abstract Most stars produce spectacular jets during their formation. There are thousands of young stars within 500 pc of the Sun and many power jets. Thus, protostellar jets may be the most common type of collimated astrophysical outflow. Shocks powered by outflows excite emission lines throughout the spectrum, exhibit a rich variety of structure, and motions with velocities ranging from less than 20 to over 500 km s1 . Due to their relative proximity, proper motions and structural changes can be observed in less than a year. I review the properties of classical Herbig-Haro objects, irradiated jets, and outflows emerging from the nearest massive star forming regions. Protostellar outflows are ideal laboratories for the exploration of the jet physics.

J. Bally () Department of Astrophysical and Planetary Sciences, University of Colorado, Boulder, CO 80309, USA e-mail: [email protected]

K. Tsinganos et al. (eds.), Protostellar Jets in Context, Astrophysics and Space Science Proceedings, DOI 10.1007/978-3-642-00576-3 2, c Springer-Verlag Berlin Heidelberg 2009 

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1 Introduction Protostellar jets and outflows produce the visual-wavelength shocks knows as Herbig-Haro (HH) objects [17]. HH objects are among the most beautiful astronomical objects in the sky and provide deep insights into jet physics in general. Because most forming and young stars produce bipolar jets and outflows, they are abundant with many examples located within a few hundred parsecs of the Sun. HH jets provide crucial insights into the launch, collimation, propagation, and the physical properties of jets in general. Bipolar outflows and jets are powerful probes of various aspects of star formation. The detection of HH objects or other signatures of an outflow provide one of the easiest means by which to identify the presence of a young star. Over the past decades, many young stellar objects (YSOs) were first identified by the detection of their outflows. The structure, velocity field, and symmetries of outflows provide powerful diagnostics of protostellar accretion processes and dynamical interactions in multiple star systems and in clusters. The sizes of gaps between major shocks point to strong variations in the ejection velocity and mass-loss rates of the source YSOs. The giant, parsec-scale outflows provide constraints on the mass-loss histories of their source stars extending from 104 to over 105 years, a time-scale comparable to the formation time-scale of young stars. The spacing of major shocks indicate that major eruptions occur roughly every few thousand years. The closeconnection between accretion and mass-loss implies that accretion onto YSOs is episodic. The terminal shocks in outflows probe the interaction zone between protostellar ejecta and the ambient medium. Thus the most distant shocks from a source serve as mechanical probes of the interstellar medium with which they interact. The observed properties of the shocks provide information about the density and velocity structure, ionization state, and chemical composition of the medium. Protostellar outflows have a profound impact on the star formation environment. In the absence of massive stars, their momentum and energy injection can be a major source of turbulence generation and cloud disruption. Thus, outflows in low- to intermediate-mass star forming regions may dominate the mechanism by which star formation self-regulates. In such environments, outflows may be the most important source of feedback. Because of their proximity and large numbers, the time-evolution of protostellar outflows can be studied in a variety of ways. Thus, they provide powerful lessons that can be applied to the study of all classes of astrophysical collimated outflow.

2 Classical Herbig-Haro Objects YSO accelerate winds and jets to velocities of 10s to over 500 km s1 , several times the escape-speed from the innermost regions of protostellar disks. Internal shocks form where faster ejecta overruns slower material. Low relative collision velocities

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(less than about 40 km s1 ) tend to excite the near-IR lines of H2 if the medium is molecular and the near-IR lines of [FeII] and the visual-wavelength [OI], and [SII] lines if the medium consists of mostly weakly ionized atoms. Faster relative velocities (>60 km s1 ) can lead to the ionization of hydrogen. Charge exchange and collisional excitation form thin zones that radiate only in hydrogen recombination lines, the so-called “Balmer filaments”. The layer of fully ionized hydrogen is followed by a recombination and cooling zone where both Balmer and forbidden lines are produced. This zone tends to have temperatures of order 104 K, set by the thermostating effects of the common, visual-wavelength forbidden lines. At most only one H˛ photon can be produced by each recombining H atom because collisions do not have sufficient energy to excite the n D 2; 3, or higher energy levels of H. However, the 2 eV forbidden transitions of species such [SII] can be readily excited by collisions. As hydrogen recombines, the low electron density insures that trace ions have long lifetimes so that their forbidden lines can be collisionally excited over and over before they recombine. Thus, the intensity of the forbidden emission lines can become comparable to or greater than H˛. Shocks with speeds higher than about 150 km s1 excite species such as [OIII]. Shocks with speeds lager than about 300 km s1 can sometimes be detected in X-rays and non-thermal radio emission. When launched from Class 0 or young Class I sources, primary flows tend to be molecular and are often traced by species such as H2 , CO, and SiO which typically exhibit radial velocities of 20 to 100 km s1 . Outflows from Class 0 sources are very dense with n(H2 ) in the range 104 to over 107 cm3 , have large mass loss rates of order 106 to more than 105 Mˇ yr1 , and high mechanical luminosities. Weak maser emission in species such as H2 O are occasionally seen in the youngest outflows from low-mass stars. However, bright maser emission is generally associated only with high-mass protostars. Somewhat more evolved Class I YSO tend to drive faster jets dominated by HI and low-ionization potential metals rendered visible by their forbidden line emission, have lower densities around 102 to over 104 cm3 , and higher speeds in the range 100 to 400 km s1 . More evolved Class II YSOs (classical T-Tauri stars) tend to have much fainter and lower mass-loss rate jets. Primary jets and winds transfer momentum to and entrain their surroundings by means of low-velocity shocks propagating into the medium. These shocks can sometimes be seen in H2 emission when the interaction is with the molecular cloud. Most high-velocity molecular emission observed at sub-mm, mm, and cm wavelengths is produced by gas entrained and accelerated by such secondary shocks. Species such as CO and other molecular transitions probe the mass and radial velocity of sweptup and entrained gas in an outflow, but only in the molecular cloud. When primary jets and winds blow out of their parent clouds, they no longer entrain molecules. Sometimes, the 21 cm line of HI can be used to trace the entrained atomic gas. The mm-wave transitions and HI are excited by collisions at the ambient gas temperature of the cloud and therefore do not require shocks to be observable. CO, other easy-to-excite molecular transitions, and HI when it can be discerned from Galactic emission, trace the total amount of momentum injected into the cloud and the amount of mass accelerated by an outflow over its lifetime.

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Protostellar jets vary in their ejection velocities, mass-loss rates, orientations, and probably degree of collimation. Ejection velocity variations with amplitudes of tens of km s1 on time-scales ranging from months to decades produce low-excitation shocks that render jets visible close to (within less than 0.1 pc) their sources. Larger but less frequent variations in ejection velocity, and presumably mass-loss rates, are responsible for the outer shocks and large gaps in between. Large eruptions associated with big velocity increases over-take older, slower ejecta to produce the giant, usually chaotic shock complexes located from 0.1 to many parsecs their driving YSOs. The giant HH 34 outflow complex in Orion provides a beautiful example [7]. Dozens of parsec-scale outflows, some of which have lengths exceeding 10 pc have been identified [16]. Internal shocks often splash sideways and contribute to the creation of relatively wide-angle outflow cavities filled with slower-moving material. Cavities may also be filled and formed by wide-angle winds launched at larger disk radii with lower velocities than the axial jets. Gas displaced and accelerated by these jets and winds may constitute the bulk of low-velocity material in bipolar molecular outflows which is traced by species such as CO. Many outflows exhibit bends indicating C-shaped deflections or point symmetries. Such symmetries provide clues about the dynamical environment of the engine; S-and Z-shaped symmetries indicate that the outflow axis has changed over time, perhaps due to precession induced by a companion, or interactions with sibling stars in a cluster. C-shaped bends indicate motion of surrounding gas (side-winds), or the motion of the outflow source itself.

3 Irradiated Jets and Outflows The discovery of irradiated jets embedded in HII regions and in UV-rich environments [2] permit the measurement of flow properties using the standard theory of photo-ionized plasmas which can provide a more robust method of density measurement than the analysis of the highly non-linear theory of shocks. External radiation renders visible much weaker jets and outflows than the flows seen in dark clouds where only shock processed gas can be seen at visual and near-IR wavelengths. Many irradiated jets have been discovered in the Orion Nebula and in NGC 1333 [2], in M43 [19], in the Carina Nebula [18] - see Figs. 1 and 2), and some other HII regions such as W5 (Stringfellow et al., this volume). A significant subset of irradiated jets show extreme bends indicatting deflection by a side-wind, radiation pressure, or the rocket effect. The sub-arcsecond resolution of HST was required to identify dozens of irradiated jets in the core of the Orion Nebula [1, 3]. Some, such as HH 514 emerging from the proplyd HST 2, exhibit pronounced kinematic and intensity asymmetry. HH 508, emerging from one of the four companion stars to  1 Ori B, the northern member of the Trapezium, is a onesided microjet which has the highest surface brightness of any known HH object in H˛ because it is located within 103 AU of an OB star. Bally et al. [1] and Bally

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Fig. 1 Irradiated jets in the Carina Nebula illuminated by the Trumpler 14 cluster. A color version of this plot can be found in the electronic version of this book and in Appendix A (Fig. A.1)

Fig. 2 The HH 666 irradiated jet emerging from dust pillars located southwest of  Car. The image was obtained with the ACS camera on the Hubble Space Telescope (see [18]). A color version of this plot can be found in the electronic version of this book and in Appendix A (Fig. A.2)

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and Reipurth [2] noted that most low-mass stars located in the southwestern part of the Orion Nebula are surrounded by parabolic arcs of emission indicating deflection of circumstellar material away from the nebular core. These were dubbed ‘LL Ori’ objects after the prototype first noted by Gull and Sofia [10] who though that LL-Ori was an example of a wind-wind collision front. HH 505 contains both a jet and an LL Ori bow. Masciadri and Raga [12] modeled the HH 505 parabolic bow as a jet deflected by a side-wind. They were able to reproduce the H˛ morphology of HH 505 and its source jet, showing that the bow is produced by weak secondary shocks formed where the side-wind interacts with jet material moving away from the jet axis after passing through the bow shock at the head of the jet. The ACS images demonstrate that many LL Ori objects, including LL Ori itself, contain jets which are frequently asymmetric. Bally et al. [3] show that while most LL Ori-type bows and bent jets in the southwestern quadrant of the Orion Nebula may be deflected by a large-scale outflow of plasma from the nebular core, even in the absence of such a side-wind, radiation pressure acting on dust, and the asymmetric photo-ablation of a neutral jet beam can also deflect irradiated jets. As the neutral jet beam emerges from the circumstellar environment into the irradiated environment of the HII region, the photo-ionized skin of the jet expands away from the jet core. For most irradiated jets, the radiation field is highly anisotropic. Thus, the photo-ablation flow deflects the jet away from the illuminating star.

4 Outflows from Massive Stars While most low-mass stars produce highly collimated jets and outflows during their formation, stars with masses above about 105 Mˇ sometimes generate wide-angle and explosive flows. The closest and best known example of a massive star forming complex is the BN/KL region in the OMC1 cloud core located immediately behind the Orion Nebula. A powerful and poorly collimated outflow emerges from a group of massive protostars embedded in BN/KL. Multi-epoch radio-frequency images show that the three brightest radio-emitting stars in OMC1, sources BN, I, and n, have proper motions (motions in the plane of the sky) of 26, 15, and 24 km s1 away from a region less than 500 AU in diameter from which they were ejected about 500 years [8, 9]. Apparently, a non-hierarchical multiple star system containing at least 4 massive members experienced a dynamical interaction resulting either in the formation of a tight binary or possibly a stellar merger whose (negative) gravitational binding energy ejected these stars from the OMC1 core. With estimated stellar masses of 10, 20, and 10 Mˇ for BN, I, and n respectively, the kinetic energy of the stars is 2  1047 ergs, comparable to the kinetic energy in the CO outflow emerging from this region. This energy must be generated by the infall of two or more stars into a deeper gravitational potential well. Assuming that source I is a binary containing two 10 Mˇ stars, its members must be separated by less than 11 AU, the orbital period must be shorter than 7 years, and the perihelion velocity of the stars must be

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at least 70 km s1 . The H2 fingers and their optical counterparts (HH 201, HH 205 through 210) consist of hundreds of individual bow shocks that give the outflow the appearance of a wide-angle explosion. Proper motion measurements of the fingers indicate a dynamical age of between 500 to 1,000 years. This outflow is clearly not powered by a jet. Rather, it appears to have been created by an explosive event in the OMC1 cloud core. My group has been developing a model in which at least 4 massive stars formed within the OMC1 cloud core. Assume that they started to accrete as low-mass protostars at random locations relatively far from each other within OMC1. As they orbited within the gravitational potential of the core they accreted gas by the Bondi-Hoyle process. Low-mass protostars and cores that happen to move toward the denser center will tend to accrete more material and experience orbital damping due to the combined effects of accretion and the gravitational drag of the wake that forms behind the stars (dynamic friction). Nick Moeckel has been modeling this process and finds that within a few hundred thousand years, the protostellar seeds grow into massive stars and sink to the center of the core where they form a non-hierarchical system of massive stars. Such systems are subject to three and four body encounters that eventually result in the formation of a compact binary and the dynamical ejection of the least massive members. Prior to the dynamical decay, matter located within about 300 AU of the cluster, comparable to the interstellar separation, will either be accreted onto individual circumstellar disks that have outer radii smaller than about 1/2 to 1/3 times the typical interstellar separation, or be expelled to beyond 300 AU by gravitational torques. Taking the mean interstellar spacing before decay to be about 100 AU, the disk outer radii would be 30 AU where for a stellar mass of 10 Mˇ , the Kepler speed is about 17 km s1 . Since non-hierarchical multiples likely have chaotic orbits, disks may be truncated at somewhat smaller radii. During the final penetrating encounter that led to the formation of an AU-scale binary whose gravitational potential energy expelled the stars from the region, the circumstellar disks of the stars that formed the binary would be destroyed. The interaction would eject their contents at roughly the Kepler velocity. If the binary consists of a pair of 10 Mˇ stars separated by less than 6 AU as required by the energetics, the outer radius of any surviving disk around either star can be no greater than about 1 to 2 AU. Matter ejected from a 1 AU orbit around a 10 Mˇ star would have a velocity of order 100 km s1 . Gravity will also accelerate the individual stars to about this speed if the periastron is a few AU. In a prograde (head-on) encounter between a star and the other star’s disk, matter can be ejected with speeds of 200 to 400 km s1 . It is proposed that the high velocities associated with the fastest ejecta in the BN/KL outflow were generated by the disruption of the inner-most portions of circumstellar disks around OMC1’s most massive stars. The total mass of high-velocity ejecta is expected to be comparable to the initial masses of the circumstellar matter located between 1 to 30 AU; this could easily be comparable to the observed mass in the fast ejecta in Orion, about 0.1 Mˇ .

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Orbital motion of the massive stars prior to the dynamical decay may have amplified magnetic fields within the OMC1 cloud core. If the shear-dynamo process if efficient, magnetic fields within a few hundred AU of the massive proto-cluster could reach equipartition with the gravitational potential. For the BNKL core, this implies mean magnetic field strengths of order 10 gauss in the inner few hundred AU of the cluster. The total amount of energy stored in such a magnetic field is around 1047 to 1048 ergs, comparable to the kinetic energies of the stars ejected by the dynamical decay and the OMC1 outflow. After the stars were ejected about 500 years ago, the magnetic stress in the region would have exceeded the gravitational potential energy of gas and stars left behind. Thus, the magnetic fields would drive supersonic expansion of the magnetized medium at about the local Alfven speed which for conditions appropriate for OMC1 would have been about 20 km s1 , comparable to the observed velocity of the bulk of the mass in the OMC1 outflow traced by CO and other molecules. In summary, it is proposed that orbital decay of a small group of accreting massive stars led to the formation of non-hierarchical cluster and the amplification of the ambient magnetic field to equipartition values. The dynamical decay of the stellar cluster ejected the stars from OMC1 (radio sources BN, I, and n). The disruption of the innermost circumstellar environments during the final close encounter that led to the formation of a compact binary (presumably source I) launched the fastest ejecta that produced the high-velocity H2 and [Fe II] “fingers”. The expansion of the magnetized medium accelerated the 10 Mˇ associated with the lower velocity molecular outflow. Located at a distance of 725 pc, the Cepheus A (Cep A) region is the second closest region of active high mass star formation. Observations also provide evidence

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Fig. 3 The precessing beam of the HW#2 jet in Cepheus A as imaged in the 2.12 m emission line of H2 . Taken from Cuningham, Moeckel, and Bally [4]

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for dynamical interactions in that region. Recent near-IR H2 observations have produced evidence that the 15 Mˇ (B0.5; [15] and references therein) protstellar source HW2 produces a pulsed, precessing jet that has changed is orientation by about 45 degrees in roughly 104 years. Figure 3 shows that the eastern side of Cep A contains at least four chains of shock-excited H2 emission that terminate in four distinct bow shocks. The bow shock orientations indicate that these flows emerge from the cluster of IR sources embedded within the Cep A cloud core, with the most luminous being HW2. HH 174, located 5’ due east of HW2, is the most distant shock. Additional terminal bow shocks and shock trains are found closer to HW2, but displaced progressively towards the north with smaller position angles (PAs) with respect to HW2. The axes of these flows change in 10 degree increments from PAD90 to 45 degrees. Observations with the VLA indicate that today, the thermal radio jet emerging from HW2 is oriented towards the northeast at PA  45 degrees and has motions of order 500 km s1 [6]. The orientations of the visual and near-IR reflection nebulosity is consistent with this latter outflow axis. An intriguing possibility is that since the ejection responsible for the HH 174 shock, the orientation of the HW2 outflow has changed by 45c i rc in discrete events during which outflow activity from HW2 increased dramatically. Thus, HW2 may drive a pulsed and precessing jet. What could cause the orientation changes and periodic eruptions indicated by these observations? A likely possibility is forced precession of the accretion disk surrounding HW2 triggered by the motion of a binary companion in an eccentric orbit whose orbital-plane is NOT co-planar with the disk. Patel et al. [15] found direct evidence for a massive circumstellar disk surrounding HW2. Additionally, Mart´ın-Pintado et al. [11] found evidence for a massive “hot core” displaced from HW2 by about 0.6”. The observed properties require that it is internally heated by a star with a projected separation of less than 500 AU from HW2. A third companion star may be responsible for the remarkable “water maser arc” found in Cep A [5]. Moeckel and Bally [14] have shown that, unlike low-mass (1 Mˇ ) stars, massive stars surrounded by massive disks have a relatively high probability of capturing a sibling cluster member into a highly eccentric elliptical orbit. We speculate that such capture occurred in Cep A HW2 about 10 to 20 thousand years ago. The resulting misalignment of the orbital plane and the HW2 disk can cause the disk to precess. Periastron passages of the captured companion may drive quasi-periodic episodes of accretion and mass-loss responsible for the pulsed, precessing jet evidenced by the observations. The effects of repeated passages of a non-coplanar companion in an eccentric orbit were investigated by Moeckel and Bally [13].

References 1. Bally, J., O’Dell, C. R., & McCaughrean, M. J. 2000, AJ, 119, 2919 2. Bally, J., & Reipurth, B. 2001, ApJ, 546, 299 3. Bally, J., Licht, D., Smith, N., & Walawender, J. 2006, AJ, 131 4. Cuningham, N., Moeckel, N. & Bally, J. 2009, AJ, (in press)

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5. Curiel, S. et al. 2002, ApJ, 564, L35 6. Curiel, S., et al. 2006, AJ, 638, 878 7. Devine, D., Bally, J., Reipurth, B., & Heathcote, S. 1997, AJ, 114 , 2095 8. Gomez, L., Rodriguez, L. F., Loinard, L., Lizano, S., Poveda, A., & Allen, C. 2005, ApJ, 635, 1166 9. Gomez, L., Rodriguez, L. F., Loinard, L., Lizano, S., Allen, C., Poveda, A., & Menten, K. M. 2008, arXiv:0805.3650v1 10. Gull, T. R., & Sofia, S. 1979, ApJ, 230, 782 11. Mart´ın-Pintado, J., Jim´enez-Serra, I., Rodr´ıguez-Franco, A., Mart´ın, S., & Thum, C. 2005, ApJ, 628, L61 12. Masciadri, E., & Raga, A. C. 2001, ApJ, 121, 408 13. Moeckel, N., & Bally, J. 2006, ApJ, 653, 437 14. Moeckel, N., & Bally, J. 2007, ApJ, 656, 275 15. Patel, N. A., Curiel, S., Sridharan, T. K., Zhang, Q., Hunter, T. R., Ho, P. T. P., Torrelles, J. M., Moran, J. M., G´omez, J. F., & Anglada, G. 2005, Nature, 437, 109 16. Reipurth, B., Bally, J., & Devine, D. 1997, AJ, 114, 2708 17. Reipurth, B., & Bally, J. 2001, ARA&A, 39, 403 18. Smith, N., Bally, J., & Brooks, K. J. 2004, AJ, 127, 2793 19. Smith, N., Bally, J., Licht, D., & Walawender, J. 2005, AJ, 129, 2308

The Star-Jet-Disk System and Angular Momentum Transfer Lee Hartmann

Abstract The interaction between the stellar magnetic field and the accretion disk in T Tauri stars is a complex, poorly-understood region. Both accretion and angular momentum loss are driven in some manner from this region. I discuss some of the models of disk-star interaction with a focus on angular momentum regulation, suggesting that time-dependent field line structure is essential to explain the slow rotation of many accreting T Tauri stars.

1 Introduction Many T Tauri stars rotate slowly, at rates 10% of breakup velocity. This is surprising, as it is thought that stars accrete most of their mass from disks. In addition, there is a tendency - most pronounced among 0:8 Mˇ stars - for stars with inner disks to rotate slowly. Following the initial suggestion by K¨onigl [29], many efforts have been made to explain this slow rotation by magnetically coupling the star to its disk at a sufficiently large radius (e.g., [5, 47]; see review by [3]). However, the L. Hartmann () University of Michigan, 830 Dennison, 500 Church St., Ann Arbor, MI 48105, USA e-mail: [email protected]

K. Tsinganos et al. (eds.), Protostellar Jets in Context, Astrophysics and Space Science Proceedings, DOI 10.1007/978-3-642-00576-3 3, c Springer-Verlag Berlin Heidelberg 2009 

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problem is made much more difficult by the observation that disk-coupled stars are accreting, and thus must be continually adding angular momentum at the same time. Dissatisfaction with the disk-stellar magnetosphere model of spindown has lead Matt and Pudritz [39, 41, 42] to reintroduce a form of magnetically-coupled stellar wind spindown instead; and recent observations suggest that stellar winds of T Tauri stars (as opposed to disk winds or jets) are more substantial than previously thought [9, 30]. In view of the complexity of the problem, and the perceived problems with the disk braking scenario, it is worthwhile reviewing some of the basic issues involved, taking perhaps a different perspective than usually adopted in the hope that it can result in new approaches.

2 The Energy Problem To begin, it is important to emphasize that the production of a slowly-rotating star from an accretion disk implies that a substantial fraction of the accretion energy must go into accelerating some material outward. It is easiest to see this by making the simple assumption that we wish to produce a star with essentially negligible angular momentum compared with the angular momentum of the Keplerian disk at its inner radius Ri . This means that the energy loss from disk material must be of order EP  .1=2/MP v2 D GMMP =.2Ri / ;

(1)

where MP is the mass accretion rate. The energy requirement has the consequence that the spindown of low-mass stars almost certainly involves the accretion process rather than a stellar wind driven by magnetic activity. The reason is that protostellar accretion of low-mass stars, which occurs over timescales of a few times 105 yr, involves the time-averaged release of much more accretion luminosity than stellar luminosity (e.g., [27]). Even though the magnetic activity of young stars is much greater than that of the Sun, this energy release is fundamentally the result of magnetic tapping into the energy resources of the star, which over time must be at most a fraction of the stellar luminosity. This makes it very hard for a pure stellar wind to account for spindown on the necessary short timescales, although it may make a modest contribution in the post-T Tauri phase (see below).

3 The Protostellar Phase The second point to emphasize is that, as most of the stellar mass is accreted during the protostellar phase, most of the angular momentum loss must occur then as well, with only a small fraction to deal with during T Tauri evolution. This point is emphasized by recent infrared spectroscopy of Class I sources, which indicate that

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these objects may rotate about twice as fast as typical T Tauri stars - still well below breakup [6]. There is evidence that protostellar accretion cannot be steady in general, but must alternate between high and low states of accretion [27, 19]. In particular, very high accretion rates, such as in FU Ori systems, might generally result in crushing magnetospheres with more or less direct accretion onto the central stars. Such accretion rates lead to relatively thick inner disks which may advect substantial amounts of energy into the young star [45, 28]. During this process, one might expect that the onslaught of so much high-angular momentum material causes the stellar surface to spin up, with consequent generation of large shears in the upper layers of the star. Such large velocity gradients might lead to the production of especially strong magnetic fields which could then drive strong outflows with major angular momentum loss. Alternatively, such episodes of rapid accretion might simply result in spin up, with spin down dominating during the longer timescales of low accretion, where the magnetosphere can reestablish itself and attach to the disk at large radii. Even if angular momentum fluxes are large during protostellar evolution, it appears likely that at least some spin braking must occur during late protostellar evolution and the T Tauri phase (see below).

4 Stellar Wind Braking In addition to the energetic problems with stellar wind braking mentioned above, there is the problem of the rapidly rotating stars in young clusters (e.g., [53,48,49]), which seemingly can only be explained if the fastest-spinning T Tauri stars contract to the main sequence with little or no angular momentum loss [20, 2]. There is still room, however, for some stellar wind braking at late times. The frequency of inner disks decreases dramatically over several Myr, reaching very low levels at ages of 5 Myr [24]. Some additional angular momentum loss by winds would help explain the slow rotators also seen in clusters like the Pleiades and ˛ Per. As John Stauffer and I pointed out long ago [20], the observations are consistent with stellar wind braking as long as the angular momentum loss “saturates” at rapid rotation (i.e., that the angular momentum losses of rapid rotators are not much higher than those of slow rotators, in analogy with the saturation of coronal activity for rapid rotators). In this case, slow rotators, with small rotational angular momentum, can be significantly braked while the rapid rotators, by virtue of their higher angular momentum, are much less affected. Bouvier et al. [2] estimated that a modest amount of angular momentum loss, of order 1035 erg, would be consistent with observations of rotational evolution. This corresponds, for example, to a wind mass loss rate of roughly 1010 Mˇ yr1 and a wind Alfven radius of order 10R  10 Rˇ . Scaling from the detailed results for winds in dipole field geometries by Matt and Pudritz [41] suggests that the surface dipole magnetic field would need to be of order 100 G, reasonably consistent with T Tauri upper limits [25]. The energy flux required to power a wind of this magnitude is of order 1031 erg s1 ,

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which is a bit higher than X-ray luminosities but not unreasonably so. Thus, wind braking can be important at ages  5 Myr for slow rotators, but is unlikely to account for the slow rotation of accreting T Tauri stars, which would require a wind mass loss rate closer to 109 Mˇ yr1 [41, 42].

5 Magnetospheric Braking Figure 1 shows the paradigmatic stellar magnetosphere-disk coupling, with field lines channeling disk material as it falls onto the star, adding angular momentum, while other (outer) field lines potentially transfer angular momentum outward to the disk. As the moment of inertia of the disk material that can be connected to the stellar magnetic field is quite small relative to the star, the only way that spindown can be effected is if either the angular momentum is transferred to the outer disk via viscous stresses or is put into a disk or X-wind. Initially, K¨onigl [29] adopted the Ghosh and Lamb [12, 13] model for the magnetic field interaction with the accretion disk. In this model, the field lines move in a steady state through disk material, being somewhat bent by the inertia of the disk matter. A similar analytic model was constructed by Collier Cameron and Campbell

Wind from disk Magnetically-heated accretion columns add angular momentum coronal loops? dusty disk

gas disk inside dust evaporation radius

Connection to outer diskoutside co-rotation?

Fig. 1 Schematic view of disk-stellar magnetosphere interaction. The magnetosphere truncates the disk at some point, probably inside the radius at which dust sublimates, and then material accretes supersonically along magnetic flux tubes onto the star, where it creates excess continuum emission as it shocks. The accreting material adds angular momentum to the star, which must be taken away to prevent spinup, either by coupling to the disk outside of corotation and/or driving a wind (the stellar wind is probably unimportant). A color version of this plot can be found in the electronic version of this book and in Appendix A (Fig. A.3)

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[5]. More recently, numerical simulations of this process have been undertaken (e.g., [33,34,51]; and references therein). In these works a quasi-steady state can be set up by the adoption of a large turbulent diffusivity, which allows the magnetic field to sweep through the disk (although even in this case the behavior is time-dependent). Shu et al. [47] criticized application of the Ghosh and Lamb model by arguing that it is extremely unlikely that the field can slip smoothly and steadily through disk material. To avoid this difficulty, and to construct a steady model, Shu et al. assumed that the stellar magnetic field only interacts with the disk at precisely the corotation radius - this is the X-wind model. However, the assumption of disk truncation at corotation is problematic. Accretion rates and magnetic fields are at least somewhat variable, over timescales much shorter than the star can change its rotation speed in response. More generally, we know that stellar magnetic fields are not axisymmetric, because we observe rotational modulation of starlight (e.g., [22], and references therein). Thus, it is implausible that the stellar field truncates the disk precisely at corotation at every longitude. An even bigger problem for the X-wind model has arisen from the recognition that the inner disk radii estimated from infrared excesses correspond to the region where dust is sublimated by the stellar radiation field, and does not necessarily correspond to the inner disk edge [44], which if at smaller radii would be inconsistent with corotation. Recent interferometric results suggest emission interior to the dust destruction radius [11]. Even more problematic is the estimate of inner disk radii from CO profile observations, which suggest that on average the inner gas disk radius is at 0:7 that of corotation [4].

6 Twist and Shout The most plausible general case is one in which many magnetic field lines connect to disk regions with differing angular velocity than that of the star. This precludes a steady state. Initially poloidal field lines become twisted up, balloon outward, and then open up and reconnect [52, 36, 35]. Matt and Pudritz [40] argued that this process of field line blow out strongly reduces the magnetic angular momentum transfer from the star to the disk, because the field lines cannot be twisted up beyond a reasonable limit. Matt and Pudritz [39, 41, 42] then argued that a stellar wind coupled to the stellar magnetic field must be responsible for T Tauri spindown. However, the required mass loss rates are so large as to require that a significant fraction of accretion energy be used to drive the stellar wind (Sect. 2), and by some unspecified mechanism. The calculations of Matt and Pudritz [40] did not take into account angular momentum loss accompanying the driving of mass loss as the field lines balloon outwards, as repeatedly shown in numerical simulations ***(e.g., [50]; Shibata and Uchida 1985; [21, 15, 16]; Fig. 2). As shown in the calculations of Goodson et al. [15], Matt et al. [38], Ustyugova et al. [51], and others, a centrifugal flow is driven from the field lines connecting with the disk as they bend outward, while a hot-

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coronal gas

disk wind C IV, OVI

Hα, etc.

Fig. 2 Twister model of magnetic coupling and ejection. As the field lines rooted in the star and disk twist up, they bulge upward and outward. These bulging loops result in an enhanced geometry outside of the field line to drive a disk wind (i.e., magnetic field forced outward). As the field expands, material on both sides of loops starts to either fall toward the star or back toward the disk; eventually the density drops sufficiently that the magnetic heating drives the material to coronal temperatures, at which point the gas stops falling and becomes hydrostatic. Eventually the twisting becomes so large that the field opens up and reconnects, driving material that originated in the disk outward along field lines that connect to both the disk and the star, drawing angular momentum from both

ter flow can result from material near the top of the loop. The bending outward of the field lines connected to the disk provide a favorable geometry for centrifugally driving outflow, even for regions interior to corotation. If there is sufficient non-axisymmetry in the field geometry, one might even imagine accretion at some longitudes where the field lines are tilted inward at the same time that the field lines at other longitudes are connecting to the disk and transferring angular momentum to the disk and to a “propeller flow” (e.g., [51]). Such a model would imply modulation of mass loss signatures on the stellar rotation period, an effect observed in SU Aur [14].

7 “Twister” Model with Magnetic Heating I would like to add an additional effect to the field line twisting model of winds which could enhance angular momentum loss. The accreting columns in T Tauri stars (Fig. 1) must be mechanically heated to temperatures of order 8,000–10,000 K to reproduce the observed Balmer emission line profiles [43]; adiabatic heating is

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insufficient [37]. The amount of energy input needed to explain this heating is significant. Using the Ly˛ flux of TW Hya [23], I estimate that somewhere between 1 and 5% of the accretion energy is emitted in this line alone (the uncertainty is due to the variability of the accretion rate; [1]). Now consider what would happen if one of these accreting loops begins to twist up (Fig. 2). The portion of the field line connecting to the disk will start to bulge outward; this makes it energetically unfavorable for material in the outer loop to fall toward the star, and instead will begin to fall back toward the disk. Meanwhile, material in the inner loop will continue to fall toward the star. The result will be an evacuation of material in the loop. However, if we assume that the magnetic heating continues during the twisting (and the twisting might even increase the amount of magnetic energy dissipation), the density in the loop can drop only so far- perhaps an order of magnitude or so - before the decrease in radiative cooling leads to a runaway increase in temperature to coronal values. The same heating responsible for infall seen in the chromospheric lines of T Tauri stars can drive emission measures of two orders of magnitude smaller - densities an order of magnitude smaller - to coronal temperatures. This increase in temperature will eventually halt the evacuation of material as the gas pressure becomes sufficient to support the gas against the stellar gravity. One is then left with an immense, dense, coronal loop. The subsequent twisting and the opening up of the loop will result in a massive version of a solar coronal mass ejection. The disadvantage of this coronal twister model of angular momentum loss is that it is difficult to estimate its efficiency, both in terms of how much material is involved in the mass ejection and the duty cycle of such ejection - the frequency of field line opening. Investigation of these issues really depend upon understanding just how gas gets onto the stellar magnetic field lines, as emphasized to the author by J. Stone. Nevertheless, it seems worthwhile to consider this complex problem further as the coronal twister model has several potential advantages over stellar winds or disk braking. First, the material to be ejected is already at a considerable distance from the stellar surface, and thus starts its journey already some way out of the star’s gravitational potential well. The gas already has significant angular momentum and rotational kinetic energy, and has a higher density than a pure stellar wind would likely have. These last features are basically due to the fact that the twisting of the field lines, and probably much of the magnetic heating, are powered ultimately by accretion energy of the disk. Thus, the coronal twister model provides a way of tapping into the accretion energy with a wind that at least in part is connected to stellar magnetic field lines, providing a mechanism for the “stellar wind” powering suggested by Matt and Pudritz. After the conference had ended, but before publication of the proceedings, Cranmer [7] made the first quantitative attempt to develop a mechanism by which accretion energy can be transferred to a stellar wind. In Cranmer’s model, timedependent accretion of clumps within closed magnetospheric loops excite magnetic waves which propagate to open field lines and then add momentum and energy to the stellar wind. The calculations are complex and difficult, but the idea is intriguing.

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In some sense this is similar to the qualitative model presented above, except that here I propose heating of field lines that naturally open up. Matt and Pudritz pointed out that the transfer of angular momentum from star to disk is weakened by the twisting up by field lines. The essence of the proposal here is that disk braking is not the only mechanism of angular momentum loss; as the field lines balloon out and open up, they can provide further stellar angular momentum loss by transfer to an enhanced wind.

8 Observational Consequences The other advantage of the coronal twister model is that it potentially provides explanations of some otherwise puzzling observational phenomena discovered in recent years. For example, optical and near-IR studies have begun to suggest that some accreting T Tauri stars have relatively warm or even hot winds (e.g., [9, 10, 30]), i though not necessarily strong coronal winds. More suggestive are the observations of very broad, highly-asymmetric emission line profiles of transition-region ions in the older accreting T Tauri star TW Hya, in particular C III, C IV, and O VI [8]. While Dupree et al. argue that these line profiles, which show predominant red-shifted emission, are P Cygni profiles associated with winds, Johns-Krull and Herczeg [26] argue that the absence of absorption in the red wing of the short-wavelength component by the blue wing of the longwavelength member of the C IV doublet argues against wind absorption providing the asymmetry. Instead, they argue that this extended red emission of ions formed at 1–3105 K arises in the magnetospheric infall region, as suggested by Lamzin et al. [31]. In support of this idea, Lamzin et al. [32] found that formation of this emission in the accretion shock at the stellar surface does not provide sufficiently wide emission profiles. G¨unther and Schmitt [18] tried to get around this problem by postulating an enormous turbulent velocity, which is less than convincing to this author. The asymmetries in the transition region lines of TW Hya are quite striking; there is much less blueshifted emission than typically seen in T Tauri line profiles (e.g., [43]). This blueshifted component comes from emission of material rising off the disk as it falls toward the star (see Fig. 1). As TW Hya is observed nearly poleon, the blue wing will be especially pronounced unless the emission is suppressed in the outer loop near the disk. This may simply represent a temperature gradient along the emitting loops, which have lower densities and thus higher temperatures than the accreting columns seen in optical lines. The twister model has an advantage in this regard; infall is enhanced on the star-side of the loop, while the bending outward of the field will tend to drive material up into the loop, cooling that portion. The appearance of emission from ions with characteristic temperatures of order 1–3105 K (e.g., C IV, O VI) is a natural part of the transition from chromospheric to coronal temperatures. It should be noted that this temperature regime is thermally unstable unless heating mechanisms compensate; conductive flux is unlikely to be

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important over the large distance scales involved in these loops. Thus, it is difficult to see how these ions can be present in a steady-state flow. Another prediction of the coronal twister model is that material will stop draining out of the loop once coronal temperatures are achieved. A characteristic temperature, whose isothermal sound speed corresponds to the Keplerian velocity at 3R D 6 Rˇ from a 0:8 Mˇ star, is roughly 1:8  106 K. G¨udel and Telleschi [17] and G¨unther and Schmitt [18] note that there appears to be an excess of O VII X-ray emission relative to both O VIII and O VI, implying an excess of emission measure at temperatures 1–2106 K in accreting stars. The emission measures of O VII-emitting gas seem more consistent with mass loss rates of 1010 Mˇ yr1 or lower, rather than the 109 Mˇ yr1 suggested by Matt and Pudritz [41, 42]. However, as this material is already at 2  3R , it is perhaps not unreasonable to assume an Alfven radius of order 10R , in which case a mass loss rate of 1010 Mˇ yr1 could compensate for angular momentum addition of accretion at 108 Mˇ yr1 . The above scenario is obviously very uncertain. What is clear is that T Tauri stars do accrete from disks, and there must be some way of getting rid of the excess angular momentum. It seems to this author that while global simulations of accretion and ejection are valuable, the interaction of the magnetic field with the disk is necessarily parameterized. To make real progress, it seems necessary to focus more carefully on the details of how disk matter gets loaded onto stellar magnetic field lines. Acknowledgements The ideas expressed here were in substantial part the result of conversations with Richard Lovelace and especially Jim Stone. This work was supported in part by the University of Michigan and the dwindling Hartmann-Calvet trust fund.

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Hot Inner Winds from T Tauri Stars Christopher M. Johns-Krull

Abstract P-Cygni components in the H˛ and other strong permitted lines have been known in Classical T Tauri Stars for several decades. These features trace mass loss from the source and, in particular, diagnose the mass loss close to the star. However, the exact nature of this mass loss including the driving mechanism and its thermal structure are still not well known. Recently, renewed effort has been expended to constrain the thermal structure of this inner wind in these sources, and there is now evidence that the temperature in these flows reach a few 10; 000 K. However, there is considerable debate over just how hot these winds get. This contribution provides a review of the observational evidence related to the inner winds of T Tauri stars, focussing on the evidence which constrains the thermal structure of the flow.

C.M. Johns-Krull () Rice University, Department of Physics and Astronomy, Houston, TX 77005, USA e-mail: [email protected]

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1 Introduction Strong mass loss is commonly observed around classical T Tauri stars (CTTSs). This mass loss is typically diagnosed by blue-shifted absorption components observed in permitted atomic lines (e.g. [27]) such as the Balmer lines of hydrogen, as well as in lines such as the Na D doublet and in Ca II and Mg II. Mass loss is also clearly traced by broad, blue-shifted emission in forbidden lines (e.g. [1,7]). This forbidden line emission is typically associated with jets from young stars and is extensively discussed elsewhere in this volume. The forbidden line emission traces conditions that extend to very large distances from the star. Here, we are concerned with mass loss originating closer to the central star, and specifically, we would like to understand the temperature structure of these flows. This can be a difficult task as it can be hard to clearly distinguish emission from a smoothly accelerating wind from that which might be produced in shocks present in an unresolved jet. Nevertheless, a number of observational diagnostics can be brought to bear on the issue, and these are reviewed in the current contribution. It should also be noted that we are interested in energetically or dynamically significiant flows in this review. For example, it is now generally established that CTTSs are accreting disk material at a typical rate of 108 Mˇ yr1 (e.g. [32, 13]). In order for this accretion to occur, angular momentum must be carried away, and a magnetized outflow is a natural means to accomplish this provided the outflow rate is 10% of the accretion rate (e.g. [22]). Using this as rough guide, mass loss rates of 109 –1010 Mˇ yr1 would be significant. Since TTSs are bright X-ray sources indicative of hot coronae (e.g. [10] ), they may well have solar-like winds; however, such a wind with a mass loss rate of 1014 Mˇ yr1 (e.g. [11]) is probably dynamically insignificant.

2 Diagnostics of the Cool Component of the Wind As mentioned above, mass loss has been diagnosed from permitted atomic lines for some time in CTTSs, and a number of attempts have been made to model the observed line profiles and determine mass loss rates and temperatures, as well as attempting to constrain the size of the wind acceleration region. An incomplete collection of these attempts is given briefly here. Kuhi [23] used spherical wind models to analyze Balmer and Ca II lines of CTTSs, finding typical mass loss rates of 108 Mˇ yr1 , with a temperature of 4500 K. DeCampli [5], also assuming spherically symmetric winds, analyzed Balmer line fluxes and found they could be explained by relatively cool winds with mass loss rates up to 107 Mˇ yr1 . Mundt [27] used an analysis of the Na D line profiles to argue that the winds from CTTSs are cool and accelerated close to the central star. Hartmann et al. [16], again assuming spherically symmetric wind models, computed line fluxes and profiles for several members of the Balmer series and found typical mass loss rates of a few 108 Mˇ yr1 with wind temperatures typically 8; 000 K. Natta and Giovanardi [31] analyzed the Na D profiles of CTTSs with spherically symmetric

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35

wind models and found mass loss rates in the range of 108 –107 Mˇ yr1 with temperatures typically in the range of 6,000–6,500 K. In one of the first attempts at abandoning the assumption of spherically symmetric winds, Calvet et al. [4] computed models of a wind originating from the boundary layer bewteen the disk and the star. These models were originally motivated to try to provide better fits to the observed Balmer line profiles, and the resulting fits again implied mass loss rates of a few 108 Mˇ yr1 with characteristic temperatures of 8; 000 K. Returning to spherical wind models, Johns and Basri [19] fit the Balmer lines of the CTTS SU Aur and found a mass loss rate of 5109 Mˇ yr1 with a characteristic temperature of 8,000 K. As a group, the above mass loss rates suggest a typical value of 108 Mˇ yr1 with typical wind temperatures of 8; 000 K. By the criteria set out in the introduction, all of the above mass loss rate determinations qualify as dynamically significant. Only the analysis of Mundt [27] directly addresses the issue of the proximity of the wind acceleration region to the central star; however, the vast majority of the line emission computed in all the models above originates within 10 R of the central source. Another means of constraining the distance of the wind line formation region from the star is to examine the line profile variability. The CTTSs SU Aur is particularly interesting in this regard. The wind absorption signature in the H˛ and Hˇ line profiles of this CTTS displays periodic variations, and the period of this wind signature is equal to the rotation period of the star [12, 19]. This suggests the winds from CTTSs are launched either from the stellar surface, or from a point directly influenced by the star, perhaps through its magnetic field. The general picture of CTTS winds that emerged at this time from the analysis of optical permitted emission lines was one of relatively high mass loss rate flows with temperatures <10;000 K that are launched from close proximity to the stellar surface. As new wavelength regimes were opened up and closer inspection of the optical line profiles took place, this picture started to change. Based on the winds used to fit the Na D line profiles of CTTSs, Natta and Giovanardi [31] predicted Brackett line profiles for these stars which should show P-Cygni like shapes. Najita et al. [30] surveyed the Br- line in several CTTSs and failed to find P-Cygni line shapes, instead observing primarily symmetric emission line profiles. Najita et al. argued that the winds from CTTSs were actually cooler than previously suspected. At the same time, Hartmann et al. [17] noted that the asymmetries and occasional redshifted absorption components observed in the higher Balmer series lines suggested an origin in an accretion flow. Hartmann et al. [17] calculated model line profiles for magnetospheric accretion flows, and these model calculations were expanded on and improved by Muzerolle et al. [28, 29]. These models could explain many aspects of the observed line profiles, and with their success a bit of a paradigm shift occurred. It was now generally believed that the permitted emission lines formed primarily in magnetospheric accretion flows, and that relatively cool, low mass loss rate winds contributed only some blue-shifted absorption in lines like H˛. Assuming that most of the emission in these lines did not originate in the winds made it difficult to determine mass loss rates from the winds probed by these blue-shifted absorption features.

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3 Is There a Hot Component to the Winds of CTTSs? The current interest in a potential high temperature wind component from CTTSs has largely grown out of the analysis of optical and near infrared emission line profiles of He I, and more recently from the analysis of far ultraviolet emission line profiles of a number of species. Beristain et al. [2] analyzed optical line profiles of ˚ line He I and He II in a sample of CTTSs. They noted a number of He I 5876 A profiles which showed blueshifted broad components (BC) to the line profile that appeared to be correlated with the source accretion rate as diagnosed by the optical veiling. These authors interpreted this as an accretion powered wind with a high temperature component. The exact temperature is hard to deduce, however, due to the nature of the He I line excitation. If collisionally dominated, these lines typically form in material with a temperature of 30; 000 K; however, in the presence of strong X-ray emission, these lines can have a significant contribution from recombination and therefore the temperature of the line formation region can be only 10,000 K or so. In addition to the confusion over the temperature of the material traced by the He lines, the line profile shapes do leave some room for interpretation. The blueshifted BC is typically observed as a pure emission feature. This can result from an overall blue-shift of the emitting material, but can also result from red-shifted self absorption in an accreting flow. This potential abiguity was lifted by observations of ˚ line from many CTTSs [9, 8]. As shown in Fig. 1, this line often the He I 10,830 A shows a classic P-Cygni line profile shape with absorption going below the stellar continuum. Such a shape must form in a wind, likely tracing material close to the ˚ line definitely traces wind material, central star (see also [24]). While the 10,830 A the same concerns over the temperature of the gas described above apply to this line. Nevertheless, these observations do show a wind component with a temperature of 10; 000 K or warmer. A much higher temperature component for the winds from CTTSs was first suggested by Dupree et al. [6] based on the analysis of far ultraviolet line profiles from the CTTSs TW Hya and T Tau observed by FUSE. The FUSE data for T Tau are quite noisy, so the assertion is based primariy on the C III and O VI lines of TW Hya observed by FUSE. These lines are shown in Fig. 2. This figure shows the essence of the Dupree et al. [6] analysis. The red side of the line profiles are fitted with a Gaussian and the observed profile on the blue side of the line is well below this Gaussian fit. This deficit is interpreted as self absorption in the wind, and the observation of this profile shape in the O VI line is taken as evidence that the temperature of the wind extends up to at least 300,000 K, the line formation temperature for O VI. Dupree et al. [6] derive a minimum mass loss rate for this hot wind of 2  1011 Mˇ yr1 for TW Hya. Johns-Krull and Herczeg [20] challenged the wind interpretation of Dupree et al. [6] for TW Hya, and by implication the reality of hot winds for all CTTSs. The key objections raised by Johns-Krull and Herczeg are that if the hot wind hypothesis is true for TW Hya, this predicts certain absorption components in the HST STIS spectrum of this star that are not observed. One such absorption that should be

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˚ line profile from six CTTSs showing clear P-Cygni line profiles Fig. 1 The He I 10830 A indicative of a strong wind. This metastable line forms at relatively high temperature; however, the excitation conditions of this line are strongly affected by X-ray emission, making a firm determination of the temperature in the wind difficult. Taken from Edwards et al. [9]

present concerns the narrow H2 emission line seen in the blue wing of the C IV line profile shown at the bottom Fig. 2. Herczeg et al. [18] showed that the H2 emission lines of TW Hya form by fluoresence in the surface layers of the disk of this CTTS within only a few AU of the star. As a result, any H2 lines which have wavelengths that place them in the blue wings of lines forming in the wind will suffer absorption by material in the wind. This is clearly seen for a H2 line in the blue wing of the C ˚ The H2 line in the bottom panel of Fig. 2 should be absorbed by C II line at 1334 A. IV in the wind, but Johns-Krull and Herczeg [20] show that this is not the case. Another key clue is the shape of different members of multiplets which are close enough in wavelength that their wings should overlap. The members of the C II multiplet in the top panel of Fig. 3 show what is expected in this case. The red wing of the blue member of the multiplet is clearly absorbed by the blue wing of the red member of the multiplet. With a characteristic formation temperature of 100; 000 K, this same behavior should be seen in the C IV lines if the wind really does reach temperatures of 300,000 K or more. As shown in the middle panel of Fig. 3 and in Johns-Krull and Herczeg [20], this is not the case. A potential complaint against the Johns-Krull and Herczeg [20] analysis of the H2 line in the blue wing of C IV is that perhaps the wind starts out quite hot near the star and cools with distance so that by the time the material is between the observer and the regions of the disk where the H2 is emitted, the wind no longer contains C IV.

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Fig. 2 High temperature lines from the CTTS TW Hya. The C III and O VI lines are from FUSE observations, while the C IV line profile is from HST STIS observations. These so-called transition region lines trace material with temperatures from a few 104 K up the 300; 000 K in the case of O VI. Also shown is a Gaussian fit to the red side of the line profiles. Dupree et al. [6] have interpreted the deficit of flux on the blue side of the line profiles relative to the Gaussian fit as evidence for self absorption in a high temperature wind. This interpretation has been challenged by Johns-Krull and Herczeg [20]. Taken from Johns-Krull and Herczeg [20]

This is a different wind structure than proposed by Dupree et al. [6] and is unlikely to be generally true for reasons shown in Fig. 4 for the CTTS RU Lup. These profiles show that the acceleration region like starts out fairly cool and that the wind heats up as it is accelerated.

4 Summary and Future Work The case for very high temperature winds originating from the star or very close circumstellar environment in CTTSs is not strong at this point. There is clear evidence for wind absorption from species such as He I, O I, C II, Si II, and Si III. All of

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Fig. 3 Emission line profiles from TW Hya testing the wind hypothesis. The upper panel shows ˚ observed with HST STIS. The black curve is the observed profile and the C II doublet at 1335 A the red curve is the observed profile shifted so that the red member of the doublet overlays the blue member. The doublet lines should have very similar shapes and the figure shows that the red wing of the blue member is significantly extincted by the wind component in the blue wing of the red member of the doublet. The middle panel shows the same test for the higher temperature C IV line. The two members have the same shape, indicating no such wind absorption at the characteristic temperature of C IV. Taken from Johns-Krull and Herczeg [20]. A color version of this plot can be found in the electronic version of this book and in Appendix A (Fig. A.4)

these lines can form in relatively cool gas with temperatures below 30,000 K, and possibly well below this temperature. This is still quite warm compared to the initial estimates of 8,000 K from optical permitted lines. It is not currently known how these winds are heated to such temperatures; however, the energy requirements for heating to these temperatures are well below what would be required to heat a massive stellar wind to 300,000 K or more. Another outstanding issue is the origin of the hot ultraviolet lines seen in CTTSs. They are too strong to be simply the product of stellar magnetic activity [21]. A natural source for the energy needed to produce these lines is in the accretion shocks that form as disk material accretes onto the stellar surface (e.g. [25, 3]). While shocks easily provide the energy required to

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Fig. 4 HST STIS line profiles showing the wind from the CTTS RU Lup. For the lower temperature lines (O I, Si II, and C II) the wind absorption starts very close to the stellar rest velocity and continues out to high velocity, indicating these lines form throughout the wind including the acceleration region near the star. On the other hand, the Si III line which probes only sightly higher temperatures shows little wind absorption near 0 velocity, indicating that the temperatures in the initial acceleration region near the star are likely low and that the wind is heated as it is acelerated

produce these hot emission lines, simplistic models of line profiles calculated from such shocks do not reproduce the observed line profiles very well [26,14]. Thus, the ultimate origin of the hot lines in CTTSs remains poorly understood.

References 1. Appenzeller, I., Oestreicher, R., Jankovics, I.: Forbidden-line profiles of T Tauri Stars. A&A, 141, 108–115 (1984) 2. Beristain, G., Edwards, S., Kwan, J.: Helium emission from classical T Tauri Stars: Dual origin in magnetospheric infall and hot wind. ApJ, 551, 1037–1064 (2001) 3. Calvet, N., Gullbring, E.: The structure and emission of the accretion shock in T Tauri Stars. ApJ, 509, 802–818 (1998) 4. Calvet, N., Hartmann, L., Hewett, R.: Winds from T Tauri Stars. II – Balmer line profiles for inner disk winds. ApJ, 386, 229–238 (1992) 5. DeCampli, W.M.: T Tauri winds. ApJ, 244, 124-146 (1981) 6. Dupree, A. K., Brickhouse, N. S., Smith, G. H., Strader, J.: A hot wind from the classical T Tauri Stars: TW Hydrae and T Tauri. ApJ, 625, L131–L134 (2005) 7. Edwards, S., Cabrit, S., Strom, S. E., Heyer, I., Strom, K. M., Anderson, E.: Forbidden line and H-alpha profiles in T Tauri Star spectra – A probe of anisotropic mass outflows and circumstellar disks. ApJ, 321, 473–495 (1987) 8. Edwards, S., Fischer, W., Hillenbrand, L., Kwan, J.: Probing T Tauri accretion and outflow with 1 micron spectroscopy. ApJ, 646, 319–341 (2006) 9. Edwards, S., Fischer, W., Kwan, J., Hillenbrand, L., Dupree, A. K.: He I 10830 as a probe of winds in accreting young stars. ApJ, 599, L41–L44 (2003)

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10. Feigelson, E., Townsley, L., G¨udel, M., Stassun, K., in Protostars and Planets V, eds. B. Reipurth, D. Jewitt, K. Keil, (Univ. of Arizona, Tucson, 2007), p. 313 11. Foukal, P.V.: Solar Astrophysics, 1st edn. (Wiley, New York, 1990) 12. Giampapa, M. S., Basri, G. S., Johns, C. M., Imhoff, C.: A synoptic of H-alpha line profile in the T Tauri Star SU Aurigae, ApJS, 89, 321–344 (1993) 13. Gullbring, E., Hartmann, L., Briceno, C., Calvet, N.: Disk accretion rates for T Tauri Stars, ApJ, 492, 323–341 (1998) 14. G¨unther, H. M., Schmitt, J. H. M. M.: Where are the hot ion lines in classical T Tauri Stars formed? A&A, 481, 735–745 (2008) 15. Hartmann, L.: Accretion Processes in Star Formation, 1st edn. (Cambridge, Cambridge, 1998) 16. Hartmann, L., Avrett, E. H., Loeser, R., Calvet, N.: Winds from T Tauri Stars. I – Spherically symmetric models. ApJ, 349, 168–189 (1990) 17. Hartmann, L., Hewett, R., Calvet, N.: Magnetospheric accretion models for T Tauri Stars. 1: Balmer line profiles without rotation. ApJ, 426, 669–687 (1994) 18. Herczeg, G. J., Wood, B. E., Linsky, J. L., Valenti, J. A., Johns–Krull, C. M.: The far-ultraviolet spectra of TW Hydrae. II. Models of H2 fluorescence in a disk. ApJ, 607, 369–383 (2004) 19. Johns, C.M., Basri, G.: The line profile variability of SU aurigae. ApJ, 449, 341–364 (1995) 20. Johns-Krull, C.M., Herczeg, G.J.: How hot is the wind from TW Hydrae? ApJ, 655, 345–350 (2007) 21. Johns-Krull, C. M., Valenti, J. A., Linsky, J. L.: An IUE atlas of pre-main-sequence Stars. II. Far-ultraviolet accretion diagnostics in T Tauri Stars. ApJ, 539, 815–833 (2000) 22. K¨onigl, A., Pudritz, R.E.: in Protostars and Planets IV, eds. V. Mannings, A. P. Boss, S. S. Russell (Univ. of Arizona, Tucson, 2000), p. 759 23. Kuhi, L. V.: Mass loss from T Tauri Stars. ApJ, 140, 1409–1433 (1964) 24. Kwan, J., Edwards, S., Fischer, W.: Modeling T Tauri winds from He I 10830 profiles. ApJ, 657, 897–915 (2007) 25. Lamzin, S.A.: The structure of shock waves in the case of accretion onto low-mass young stars. Astronomy Rep., 42, 322–335 (1998) 26. Lamzin, S.A.: Calculation of profiles of the CIV 1550 doublet formed in an accretion shock in a T Tauri Star: Axially symmetric radial accretion. Astronomy Rep., 47, 540–550 (2003) 27. Mundt, R.: Mass loss in T Tauri Stars – Observational studies of the cool parts of their stellar winds and expanding shells. ApJ, 280, 749–770 (1984) 28. Muzerolle, J., Calvet, N., Hartmann, L.: Magnetospheric accretion models for the hydrogen emission lines of T Tauri Stars. ApJ, 492, 743–753 (1998) 29. Muzerolle, J., Calvet, N., Hartmann, L.: Emission-line diagnostics of T Tauri magnetospheric accretion. II. Improved model tests and insights into accretion physics. ApJ, 550, 944–961 (2001) 30. Najita, J., Carr, J.S., Tokunage, A.T.: High-resolution spectroscopy of BR gamma emission in young stellar objects. ApJ, 456, 292–299 (1996) 31. Natta, A., Giovanardi, C.: Sodium lines in T Tauri Stars – Diagnostics of pre-main-sequence winds. ApJ, 356, 646–661 (1990) 32. Valenti, J.A., Basri, G., Johns, C.M.: T Tauri Stars in blue. AJ, 106, 2024–2050 (1993)

Hot Gas in Accretion Disks and Jets: An UV View of Star Formation Ana I. G´omez de Castro

1 Introduction: The Jets Engine During the T Tauri phase, there is a differentially rotating region attached to the top of the stellar convective layer connecting the star with the accretion disk that rotates significantly faster than the stellar surface; rotation periods during the TT phase are about 7–8 days (˝ D 0:8  0:9 day1 ) while the Keplerian frequency  1=2  3 r . This shear region is feed by turbulent, is: ˝k D 11:1day1 MMsun 3Rsun magnetized material from the accretion disk. The turbulent disk dynamo is feed by the magneto-rotational instability in the acretion disk. Shear amplifies the field producing a strong toroidal component; an external dynamo sets in. This toroidal field and the associated magnetic pressure push the field lines outwards from the disk rotation axis, inflating and opening them in a butterfly-like pattern reminiscent of

A.I. G´omez de Castro () Astronom´ıa y Geodesia, Fac. de CC Matem´aticas, Universidad Complutense de Madrid, 28040 Madrid, Spain e-mail: [email protected]

K. Tsinganos et al. (eds.), Protostellar Jets in Context, Astrophysics and Space Science Proceedings, DOI 10.1007/978-3-642-00576-3 5, c Springer-Verlag Berlin Heidelberg 2009 

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I II

III

Fig. 1 The interaction between the stellar magnetic field and the disk twists the stellar field lines due to the differential rotation. The toroidal magnetic field generated out of the poloidal flux and the associated pressure tends to push the field lines outwards, inflating them, and eventually braking the magnetic link between the star and the disk (boundary between regions I and II). Three basic regions can be defined: Region I dominated by the stellar wind, Region II dominated by the disk wind and Region III dominated by stellar magnetospheric phenomena. The dashed line traces the boundaries between this three regions. The continuous lines indicate the topology of the field and the shadowed areas represent regions the where magnetic reconnection events are likely to occur, producing high energy radiation and particles (from G´omez de Castro 2004)

the helmet streamers in the solar corona, so producing a current layer between the stellar and the disk dominated regions as displayed in Fig 1. Magnetic field dissipation in the current layer produces high energy radiation and particles. The magnetic link between the star and the disk is broken and reestablished continuously by magnetic reconnection. The opening angle of the current layer, as well as its extent, depends on the stellar and disk fields, the accretion rate and the ratio between the inner disk radius and the stellar rotation frequencies. Hot, pressure driven outflows are produced from the region closer to the rotation axis while cool centrifugally driven flows are produced by the disk; plasmoids are ejected from the current layer generating a third outflowing component. The bulk of the energy produced in this engine is released at UV and X-ray wavelenghts as in the Sun atmosphere. In the very early epochs, when extinction is high (AV  3), only the X-ray radiation from the engine is detected. Later on, about 1 Myr, extinction drops and the engine can be studied in the UV. Only in the UV, the various components of the engine can be defined and studied as well as their evolution, starting during the T Tauri phase all the way down into the main sequence. The UV emission region is unresolved; its physical size is expected to be smaller than 1 AU, including the largest structures reminiscent of the helmet streamers. As a result, there are only two observational techniques that allow distangling the various contributions: monitorings and high resolution UV spectroscopy. A detail accounting on the UV properties of the components of the engine can be found in G´omez de Castro (2008). In this brief contribution, I should highlight some of the UV properties of three key components: the accretion shocks (see Sect. 2), the recent evidence of plasma rings around T Tauri stars (see Sect. 3) and the interaction of the outflows with the circumstellar medium (see Sect. 4).

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2 Accretion Shocks The first evidence of UV radiation from accretion shocks comes from monitorings with the International Ultraviolet Explorer (IUE) which found rotational modulation of the UV continuum and line fluxes [4, 5, 17]. An important result of these campaigns is that only 50% of the UV continuum excess is rotationally modulated.Thus, a significant fraction of the UV excess is not produced by accretion shocks even in sources where rotational modulation has been detected. This is nicely shown in the Si III] and C III] profiles of RY Tau (see Fig. 2). Two observations of the same star obtained in 1993 and 2001 show that there is a variable redshifted component. From this single observation three important properties are learnt: 1. UV radiation from accretion shocks is produced on scales significantly smaller than the stellar radius as expected from accretion shocks models [7, 12]. Thus, it is expected that only matter falling onto the visible hemisphere can be detected at UV wavelengths. The fact that the variable flux component is redwards shifted supports these theoretical expectations. Moreover, the broadening shows that infalling matter should cover a significant fraction of the hemisphere to account for the broad distribution of projected velocities in the infalling gas 2. The red wing extends to velocities of 250 km/s which corresponds to the free-fall velocity1 from 1.7 R which is much smaller than the fiducial values derived for the inner disk radius

F(10–14 erg / s/ cm2 / A)

4 3

SiIII]

SiIII]

CIII]

CIII]

2 1 0 3 2 1 0

–400 –200

0

V(km/s)

200

400

–400 –200

0

200

400

V(km/s)

Fig. 2 SiIII] and CIII] UV lines observed in RY Tau (from G´omez de Castro & Verdugo 2007); RSDP processed data are plotted with a thin line and the 3-pixels average profile with a thick line. The rest wavelength of the lines and the velocity of the unresolved jet at ' 80 km/s are marked with dashed lines. Left panel: Observations obtained in Dec. 31st, 1993. Right panel: Observations obtained in March 27th, 2001. Both lines show an excess of flux in the red wing compared with the 1993 observations; this excess is shaded in the figure

1

RY Tau mass is 1.63 Msun and radius 2.4 Rsun according to Hartigan et al 1995.

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3. UV line radiation is not only produced by accretion; also the wind contributes to it. Thus accretion rates derived from the UV excess, assuming that it is caused just by magnetospheric infall, are overestimated

3 Ion Belts ˚ has been detected recently from Strong continuum FUV emission (1300–1700 A) some stars with bright molecular disks including GM Aur, DM Tau, and LkCa 15, together with inner disk gaps of few AUs [1]. This emission is likely due to energetic photoelectrons mixed into the molecular layer that likely indicates the existence of a very hot component in the inner disk. In addition, there is increasing evidence of the existence of ion belts/rings around some TTSs. An ion belt has been detected around the TTS, RW Aur [8]. A corotation radius of 4.4 R is derived and a log Te .K/ ' 4:7 and log ne .cm3 / D 11:6 are estimated. This was the first detection of such an structure around a classical TTS. Moreover, there are indications of a similar structure around AB Dor, a weak line TTS (see Fig. 3). The structure is resolved, as in AB Dor, because there is an inner hole that allows separating the stellar/wind contribution from the belt. However in a 5.7 hours time lapse the double peaked profile is lost, and the inner part of the profile is filled in again with emission [9].

8

Flux (10–14 erg s–1 cm–2 A–1)

6

f = 1.329

4 2 0 6

f = 1.794

4 2 0

–400

–200

0

200

400

Radial Velocity (Km/s)

Fig. 3 SiIII] profiles of AB Dor obtained with the HST/GHRS (see Gomez de Castro [9] for more details). In the bottom panel, the profile at phase () 0.329 is overplotted (dashed line) on the profiles at  D 0:794 (continuous line), for comparison

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Flux (erg s–1 cm–2 A–1)

2x10–12

1.5x10–12

10–12

5x10–13

0 –500

0 500 Vr (km/s)

1000:500

0 500 Vr (km/s)

1000

˚ profile of AB Dor during a normal stellar flare (left) and a transient Fig. 4 The C IV 1548 A feature probably associated with a CIR (right). Both events lasted several kiloseconds. The left profile is typical of three events that occured during the short monitoring time, while the profile on the right was observed only once. Note the presence of a narrow absorption and the very broad line wings in the right panel profile (see Gomez de Castro [9] for more details)

4 Interaction of the Outflows with the Environment The UV signature of outflows can be detected both, at small scales, where the stellar solar-like wind interacts with the material in the young planetary disk and, at large scales, where the bullets of plasma ejected from the helmet streamers interact/collide with the molecular cloud environment. At small scales, the most clear example comes from the UV monitoring of AB Dor, a very bright nearby 30 Myr old star [9]. Nine events were detected during 10.63 hours of monitoring with HST/GHRS. The C IV and Si IV UV line profiles produced by most of the events are narrow and redshifted, indicating hot gas falling onto the star during the flare. However, the strongest event produced a very broad profile with a slightly blueshifted narrow absorption. This profile lasted a few kiloseconds and thus the broad wings are most likely tracing the front shock of a corotating interaction regions or CIRs (shock fronts formed by the interaction between the slow and the fast component of the solar wind). Thus, high-resolution UV spectroscopic monitoring can be used in to study the impact of the early Sun in young planetary disks evolution as well as on planetary atmospheres embryos. In the large scales, it is known since the first spectra of Herbig-Haro objects (HHOs) obtained with the International Ultraviolet Explorer (IUE) that the working surface of jets in the cloud, the bow shocks, have a higher degree of ionisation than previously inferred from optical data (Bohm-Vitense et al. [2]; Schwartz et al. [16]; see also Gomez de Castro & Robles [6] for a compilation of all IUE observations). Studies revealed that some HHOs, like HH 1 or HH 2, produce strong emission lines of C IV[uv1], O III], Si III], and C III] [14]. These were deemed high-excitation HHOs (HE-HHOs), and were identified as regions of strong shocks (usually located

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at bow-shocks), which display very high temperatures and densities (Te D 105 K and ne D 106 cm3 ). Conversely, low-excitation HHOs (LE-HHOs), like HH43 or HH47, are characterised by the presence of the H2 Lyman band emission lines ˚ with dominant emission features at 1258, 1272, 1431, 1446, 1505, 1547 and 1562 A (Schwartz 1983). LE-HHOs were found to trace weaker shocks, lying mainly within the jet itself. However, more recent UV observations show that this classification is not as clean as it may look like. On the one hand, H2 Lyman band emission has been detected in the low resolution spectra of HH 2 obtained with the Hopkins Ultraviolet Telescope (HUT) whereas neither NV nor OVI emission were detected; thus strong standard radiative shock waves fail to reproduce the degree of ionization observed [15]. On the other hand, a high excitation dense component (T D 105 K; ne D 106 cm3 ) was detected to co-exist with a dominant low excitation component (T D 104 K; ne D 103 cm3 ) in HH 29 though the hot component filling factor is 0.1%-1% [13]. The observed co-existence of both high and low excitation gas could just result from the very low spatial resolution of the spectroscopic observations since both objects are very clumpy. However the recent detection of X-ray emission from some protostellar jets (L1551 IRS 5 and DG Tau jets by Favata et al. [3] and Guedel et al [11] respectively) suggests that important aspects of jet physics are being neglected. Our current understanding of jet formation involves the co-existence of two components: a warm high density component associated with a centrifugally driven wind from the disk and a hot, low density component associated with plasmoids ejection from the magnetic interface between the stellar field and the disk field (see Fig. 1). X-ray radiation could be caused by magnetic reconnection and dissipative processes involving magnetic fields could be responsible of the observed deviations from pure radiative shocks. The first UV images of HHOs and protostellar jets are still to be obtained to map the spatial distribution of the hot plasma and further understand the details of jet physics. Acknowledgements This work has been supported by the Ministry of Education and Science of Spain through grant AYA2007-67726.

References 1. Bergin et al.: ApJ, 614, L133, (2004) 2. Bohm-Vitense et al.: ApJ, 262, 224 (1982) 3. Favata et al.: A&A, 386, 204 (2002) 4. Gomez de Castro, A.I. and Fernandez, M.: MNRAS, 283,55, (1996) 5. Gomez de Castro, A.I. and Franqueira, M., ApJ, 482, 465, (1997) 6. Gomez de Castro, A.I. and Robles, A., INES Access Guide No. 1: Herbig-Haro Objects, ESA Scientific Publication, ESA-SP 1237, (1999) 7. Gomez de Castro, A.I. and Lamzin,S., MNRAS, 304, L41, (1999) 8. Gomez de Castro, A.I. and Verdugo, E., ApJ, 597,443, (2003) 9. Gomez de Castro, A.I. MNRAS, 332, 409, (2002)

Hot Gas in Accretion Disks and Jets: An UV View of Star Formation 10. Gomez de Castro, A.I., ApSS, in press 11. Guedel et al.: A&A, 478, 797 (2008) 12. Lamzin, S.A., Astronomy Reports, 42, 322,(1998) 13. Liseau et al.: A&A, 306, 255 (1996) 14. Ortolani & d’Odorico: A&A, 83, L8 (1980) 15. Raymond et al.: ApJ 489, 314 (1997) 16. Schwartz et al.: AJ, 90, 1820 (1985) 17. Simon et al.: AJ, 100, 1957 (1990)

49

Generalized Multipole X-Wind Model Subhanjoy Mohanty and Frank H. Shu

Abstract The X-wind model for magnetospheric accretion and outflow in classical T Tauri stars (CTTS) has gained credence in recent years for a variety of theoretical and observational reasons. However, both this model as well as other theoretical scenarios for explaining magnetospheric disk accretion assume that the stellar field, were it not perturbed by an electrically conducting accretion disk, would have a dipolar geometry (e.g., [5]; OS95 hereafter). Observations of accretion hot spot sizes and net field polarization on the surface of CTTS, however, clearly indicate that the stellar field has a complex multipolar structure. To overcome this discrepancy between theory and data, we reformulate X-wind theory without the dipole constraint. This contribution represents a brief summary of the paper by Mohanty and Shu [6]. In Sect. 1 we present the fundamental physical ideas of the generalized theory, and the associated equations; in Sect. 2 we compare the resulting theoretical prediction to recent observations, and provide some illustrative numerical simulations with multipole stellar fields.

1 Equations of the Generalized X-Wind Model The central idea of X-wind theory is that of trapped flux: In steady state, all the stellar flux threading the disk is trapped in a narrow annulus at the inner edge of the disk, called the X-region. The strong pinch means that the field in the immediate vicinity of the X-region must be nearly force-free above and below the midplane (which is not force-free). In the cold limit, where the X-region corresponds to essentially the single radius RX , the field configuration then achieves the shape of a complete

S. Mohanty () Imperial College, London e-mail: [email protected] F.H. Shu University of California at San Diego e-mail: [email protected]

K. Tsinganos et al. (eds.), Protostellar Jets in Context, Astrophysics and Space Science Proceedings, DOI 10.1007/978-3-642-00576-3 6, c Springer-Verlag Berlin Heidelberg 2009 

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fan. The outer 1/3 of these field lines bend sufficiently outward to launch a mass outflow (which keeps these field lines open – i.e., they reconnect with their stellar counterparts at infinity); the central 1/3 are too steep for disk material to climb up, and thus constitute a dead zone; and the inner 1/3 bend sufficiently inward to carry gas onto the stellar surface in an accretion funnel flow. The associated magnetic torques transfer angular momentum from the infalling gas to the footpoints of the funnel flow field lines in the X-region, and conversely impart angular momentum to the outflowing wind at the expense of the material in the X-region. The net effect in the X-region is thus an outward pinch on the interior field lines and an inward pinch on the exterior field lines towards the common mid-point RX , which is what keeps the flux trapped in this region in a fan-shape and gives the X-wind model its name. At the same time, there is a net transport of angular momentum from the accreting gas to the wind (via the X-region), which allows originally high specific angular momentum disk gas to accrete onto the star without spinning up the latter. In steady state, the radius RX is both in Keplerian rotation and in co-rotation with the star (so the field lines do not constantly wrap up), a condition known as disk-locking. Let us quantify these statements. We adopt cylindrical coordinates ($ , ', z) with the origin at the center of the star. Say the amount of trapped flux in an equatorial ring at $ D RX is ˚t . Now, the funnel field lines that are loaded with infalling gas in the X-region are the same ones that carry the gas all the way to the star. Let Fh be the fraction of the surface area 2R2 of the stellar upper hemisphere covered by hot spots of mean field strength BN h . Then the magnetic flux connected to the hot spots, Fh .2R2 /BN h , must equal the same 1/3 of the trapped flux ˚t that links to the base of the funnel flow in the upper surface of the X-region: 1 Fh .2R2 /BN h D ˚t : 3

(1)

Next, consider angular momentum balance in the X-region. With a Keplerian angular velocity of ˝X , disk gas at the disk inner edge RX has specific angular mo2 ˝X . Let JN be the average fraction of this that lands on the star mentum equal to RX via the funnel flow, and JNW the average fraction carried away by the wind. Further let the total accretion rate through the disk into the X-region be MP D , with the wind and funnel flow carrying away a fraction f and 1-f of this respectively. Finally, let T > 0 be the viscous torque exerted by disk material interior to the X-region on disk matter exterior to it. Then angular momentum conservation demands 2 2 2 MP D RX ˝X D f MP D JNW RX ˝X C .1  f /MP D JN RX ˝X C T ;

(2)

where the left-hand term is the angular momentum flowing into the X-region through disk accretion, and the three right-hand terms the angular momenta leaving the X-region via the wind, funnel flow and viscous torque respectively. Thus f D

1  JN  ; JNW  JN

(3)

Generalized Multipole X-Wind Model

53

2 where  T =.MP D RX ˝X / > 0 is the dimensionless viscous torque acting across a circle beyond RX . With the guesses that JN  0 for a small slowly rotating star, and  0 as well, (3) yields JNW  1=f . To make further progress, and obtain the location of the X-point in the meridional plane RX , we need to input the dynamics of the X-wind. With 1/3 of the trapped flux driving the wind, we have (see [2.3] of OS95 and [3.10b] of [7]).

Fig. 1 Four numerical simulations of magnetospheric accretion in the generalized X-wind model, using an arbitrary multipole stellar surface field that reproduces the observed small net polarization (i.e., loopy field configuration) on the surface of CTTS (see [6] for details). X- and y-axes in units of stellar radius R . RX D 10 R (top left), 7.5 R (top right), 5 R (bottom left) and 2 R (bottom right). Shaded quarter circle is the star; black lines show the steady-state field (lines of constant flux); grey lines show the subset of these field lines participating in the funnel flow; dashed line is the wind interface (wind field lines not shown). The uppermost funnel flow field line always lies below the equipotential curve, shown as the black curve originating at the X-point; the lowermost funnel flow line skims the disk surface inwards of RX and finally rises above it at the kink point RK , at an angle of 120ı , marked by the inclined grey line. Note the very small hot spot covering fraction on the stellar surface, in agreement with CTTS observations. Also note that the field configuration near the X-point is quite similar in all cases, but the funnel flow shape varies considerably as one approaches the star for the different RX , with the hot spot location varying as well. See Mohanty and Shu [6] for details

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1 N 1=2 .GM MP 2 R3 /1=4 ; ˚t D 2 ˇf D X 3

(4)

where ˇN is a dimensionless (inverse mass-loading) parameter measuring the ratio of magnetic field to mass flux in the frame that corotates with the footpoint of a field line labeled by the streamline averaged over X-wind streamlines, D 0  1. R1 That is, B D ˇ. / u with ˇN  0 ˇ. / d , and the normalization for the stream-function is the mass-loss rate in the wind MP W D f MP D . The quantity 2 3 1=4 RX / in (4) then provides the correct units when GM , MP D and RX .GM MP D are the fundamental dimensional quantities of the physical problem. Combining (1) and (4), we get N 1=2 Bnorm Fh BN h D ˇf



RX R

3=4 ;

(5)

where we have defined a fiducial field strength,

Bnorm 

2 GM MP D R5

!1=4 :

(6)

Finally, we impose the disk-locking condition for steady-state:  ˝ D ˝X D

GM 3 RX

1=2 :

(7)

With RX defined in terms of the stellar rotation rate ˝ by the above, substitution into (5) gives the desired general relationship: N 1=2 .GM MP D =˝ /1=2 : Fh R2 BN h D ˇf

(8)

Equation (8) encapsulates the idea of flux trapping, because it relates the amount of measured flux in hot spots on the left-hand side to independently observable quatities on the right-hand side without any assumptions about the specific multipolar character of the stellar magnetic field ultimately responsible for the funnel flow. We have not yet explicitly specified the values for ˇN and f , however. Unfortunately, these cannot be explicitly derived within our current framework. Our assumption of the cold limit (thermal speed near the inner edge RX is small compared to its Kepler speed) allows us to shrink the X-region to a mathematical point (single radius RX ) in the meridional plane, with infinitesimal disk thickness there. This makes the problem tractable by imposing the fan field configuration at RX (leading to 1/3 trapped flux fractions in the funnel flow, dead zone and X-wind each). How the vanishingly small diffusivities  and  then load field lines is however not addressable in this limiting procedure. We can make a rough guess, though.

Generalized Multipole X-Wind Model

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On the basis that 1/3 of the trapped field lines participate in the wind, it has been proposed (e.g., OS95) that the wind mass fraction f  1/3 as well. Furthermore, (3) – see discussion above – then implies JNW  1=f  1=3. For the somewhat ad N hoc X-wind loading law ˇ D .2ˇ=3/.1  /1=3 , Table 3 of Cai et al. [1] tabuN N lates ˇ versus JW ; interpolating the latter yields ˇN D 1:21 if we adopt JNW D 1=3. Please note, however, that the theory does not demand precisely these values of N Indeed, if the quantity ˇf N 1=2 in (8) is roughly constant for CTTS, the f and ˇ. equation – which is the central prediction of X-wind theory – can still be checked from the slope of a log-log plot by measuring all the other observable quantities N 1=2 can further be simultaneously without knowing ˇN and f , and the value of ˇf observationally derived as the intercept of the plot.

2 Comparison to Observations Johns-Krull and Gafford ([4]; hereon JG02) have compared CTTS observations to the dipole model predictions of various groups (including the X-wind one of OS95), as well as to the general X-wind prediction encapsulated in (8). To make N 1=2 ; in the restricted case described progress, they have assumed a constant BN h =ˇf above, where f D 1=3 and ˇN D 1:21, this amounts to postulating a constant BN h for all CTTS hot spots (roughly true, within a factor of a few, for many sources with measured BN h ). They find the observations bear little relation to all dipole predictions including OS95’s. However, the scatter is greatly reduced and a significant correlation found with the generalized X-wind prediction without the dipole constraint. Further support for the theory comes from two recent sets of spectropolarimetric observations, wherein Fh , BN h and MP D are all measured simultaneously. For both V2129 Oph and BP Tau, Donati et al. [2, 3] find strong multipole surface field components, and measure Fh BN h  100 G and 180 G respectively for the two stars. With the same stellar and accretion parameters they use (see [6] for details), and with adopted f D 1=3 and ˇN D 1:21, (8) predicts Fh BN h D 79 G and 170 G for V2129 Oph and BP Tau respectively, in very good agreement with the measured values. The consistency with both stars, and the good correlation in the study above, signal that the generalized theory can reproduce the broad quantitative features of current CTTS data. Lastly, while there is scatter in JG02’s study due to assumed constant BN h and non-simultaneously measured Fh and MP D , the agreement with Donati et al.’s data suggests our adopted values of ˇN and f are roughly valid. Upcoming coeval measurements of BN h , Fh and MP D for a large CTTS sample, analogous to the V2129 and N 1=2 , crucial BP Tau data above, should enable a good observational estimate of ˇf for understanding microphysics in the X-region.

56

References 1. Cai, M., Shang, H., Lin, H.H. & Shu, F.H.: ApJ 672, 489 (2008) 2. Donati, J.-F. et al.: MNRAS 380, 1297 (2007) 3. Donati, J.-F. et al.: MNRAS 386, 1234 (2008) 4. Johns-Krull, C. & Gafford, A.: ApJ 573, 685 (2002) 5. Ostriker, E. & Shu, F.H.: ApJ 447, 813 (1995) [OS95] 6. Mohanty, S. & Shu, F.H.: ApJ 687, 1323 (2008) 7. Shu, F.H., Najita, J., Ruden, S.P. & Lizano, S.: ApJ 429, 797 (1994)

S. Mohanty and F.H. Shu

Instabilities in Accretion Disks James M. Stone

Abstract In order to understand the production of jets and winds from disks, or from star-disk interaction, it is important to understand the internal dynamics of the disk itself. This requires studying the nonlinear regime of various hydrodynamic and magnetohydrodynamic (MHD) instabilities that operate in disks, using nonideal MHD and including the appropriate microphysics. Our current understanding of some of the most important instabilities in disks is reviewed, with focus on the magneto-rotational instability (MRI).

1 Introduction Numerical methods have played an important role in the study of the dynamics of accretion disks, including protoplanetary disks. In fact they are virtually the only way to study the time-dependent magnetohydrodynamics (MHD) in multidimensions. Early calculations [33] demonstrated how dynamic the evolution of magnetized

J.M. Stone () Princeton University, Department of Astrophysical Sciences, Princeton NJ 08544 e-mail: [email protected]

K. Tsinganos et al. (eds.), Protostellar Jets in Context, Astrophysics and Space Science Proceedings, DOI 10.1007/978-3-642-00576-3 7, c Springer-Verlag Berlin Heidelberg 2009 

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Fig. 1 Contours of the density and velocity vectors (top row), and representative magnetic field lines (bottom row) from the 2D evolution of a weakly magnetized Keplerian disk. The panels in each row are at times of 1, 2.5, and 5 orbits of the inner edge

disks can be. No matter whether the calculations begin from a rotationally supported equilibrium, or a sub-Keplerian rotation profile, the first two-dimensional (2D) simulations showed the disk would collapse on an orbital time scale. Figure 1 taken from Stone and Norman [29] demonstrates the collapse that occurs for a weakly magnetized, initially Keplerian disk in 2D. For strongly magnetized disks, the collapse could be driven by magnetic braking, a process identified as the “sweeping magnetic twist” by Shibata and Uchida. However, for weak fields, the process driving collapse was more puzzling. It was later realized it could be driven by axisymmetric modes of the magnetorotational instability [2]. Today, it seems clear that to understand the formation of jets and outflows from disks, and to understand the interaction of disks with the magnetosphere of the central star, we must understand the internal dynamics of magnetized disks. That is because the internal dynamics determines properties such as (1) the angular momentum transport and mass accretion rate, (2) the inward advection of vertical magnetic flux, and (3) the conditions in the corona of the disk, which serves as

Instabilities in Accretion Disks

59

the launching region of the wind. However, to understand the internal dynamics requires understanding the nonlinear saturation of a wide variety of hydrodynamic and MHD instabilities, with the appropriate physics for protoplanetary disks (nonideal MHD, radiation transport, ionization and recombination, etc.) included. This paper is intended as a brief review of recent progress in studies of instabilities in disks, with a particular emphasis on how they may affect the production of winds and jets from protoplanetary disks.

2 A Zoo of Instabilities Which instabilities are thought to be important in disks? It seems all of the most powerful are MHD, and a list of selected instabilities includes the following: 1. Magneto-rotational instability (MRI). Fundamental to weakly magnetized flow in orbital (Keplerian) rotation. Important for disks in every astrophysical context. 2. Parker instability. Produces vertical flux of magnetic energy in stratified disks. 3. Rayleigh-Taylor instability (RTI). Could be important in the star-disk interaction region as a process which controls the accretion of material to the stellar surface, and loading of plasma onto stellar field lines. Other MHD instabilities that have been identified and might be important in AGN disks, or disks around stellar-mass black holes, include the magneto-viscous instability (MVI) [16], the magneto-thermal instability (MTI) [1], and the photon-bubble instability [12, 5]. If the disk is very weakly ionized (e.g., ionization fraction below 1013 at 1 AU for a protoplanetary disk;[4]), then it may behave more like a hydrodynamic rather than a MHD flow. In this case, hydrodynamic instabilities might be important, such as: 1. Non-linear shear instability. Long proposed as a mechanism for angular momentum transport, however recent experiments suggest it is either not present or irrelevant in astrophysical disks [17]. 2. Baroclinic instability. Driven by radial entropy gradients in the disk [22, 20]. 3. Gravitational instabilities. Will be important for sufficiently massive disks, roughly Mdisk > 0:1Mstar . 4. Kelvin-Helmholtz instability (KHI). Can be driven in the shear layer formed by settling of dust grains to the center of the disk [38, 6]. 5. Dust streaming instability. Recent work has shown the radial drift of dust grains in a protoplanetary disk can drive turbulence and produce clumping of dust [39, 19]. Clearly there are far more processes that can be reviewed in this short paper. Most of the focus in the following sections will be on the MRI, and the application of the MRI to protoplanetary disks (the source of HH jets and protostellar outflows), since there is general agreement the MRI is of fundamental importance in disks.

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3 The MRI Figure 2 demonstrates the mechanism of the MRI. The image on the left is a volumetric rendering of the density taken from a global, 3D simulation of a black hole accretion disk [13]. A small pie wedge has been cut from the disk in order to show fluctuations in the interior. Imagine a 2D poloidal plane in the disk with a vertical magnetic field line. This plane is shown as the red square in the image, which is expanded in the upper right panel. Now imagine radial perturbations to the vertical field line. The top view of this plane (shown in the lower right) demonstrates that because of the orbital dynamics, the radial perturbations produce azimuthal motion (the flow moves faster at smaller radii, so the inward perturbation moves ahead, and outward perturbation falls behind, of the red square). Therefore there is an azimuthal component to the tension in the field line, represented as a spring in the lower right panel. This tension produces angular momentum transport, from (to) the inner (outer) fluid parcel as it tries to slow down (speed up) the azimuthal motion. However, this transport is de-stabilizing. As angular momentum is removed from the inner parcel, it falls to an even lower orbit, speeds up even more, and increases the azimuthal shear, and conversely the outer parcel moves to a larger radii as it gains angular momentum, slows down, and falls farther behind. The result is to increase the perturbation, which increases the tension, which further increases the perturbation, and the process runs away. A tremendous amount of effort has gone into the study of the nonlinear regime of the MRI through MHD simulations, and there are far too many results to review properly here. Generally it is found the MRI saturates as MHD turbulence in which Maxwell stress provides a significant transport of angular momentum. An example of the result from simulations that adopt the local, shearing box approximation (in which only a small radial patch of the disk is computed) is shown in Fig. 3. Most previous simulations using the shearing box have adopted rather narrow domains in the radial direction, no more than one scale height H . Calculations in wider boxes are very expensive, because the orbital flow becomes supersonic across

Radial perturbations to vertical B Side view (r-z plane)

Instability when (k · VA) ⱕ

Top view (r-f plane)

3W

Fig. 2 Mechanism of the MRI. See the text for details. A color version of this plot can be found in the electronic version of this book and in Appendix A (Fig. A.5)

Instabilities in Accretion Disks

61 Pseudocolor Var: dVy 0.4000 0.2000 0.0000 –0.2000 –0.4000

Max: 0.6242 Min: –0.5794

Z

Y X

Fig. 3 Images of the azimuthal velocity fluctuations in shearing box simulations of the MRI in a standard box (radial extent of H , left), and a wide (radial extent of 32H , right). Different color tables are used in the two images; the fluctuation amplitudes are similar. A color version of this plot can be found in the electronic version of this book and in Appendix A (Fig. A.6)

scales larger than H , and this places a severe constraint on the time step for explicit methods. Recently, orbital advection methods [23] have been extended to MHD [18,30], and this has allowed simulations in very wide domains (for example, Fig. 3). A more quantitative result from studies of the MRI in wide boxes is shown in Fig. 4, which plots the time evolution of the volume-averaged total stress (Maxwell plus Reynolds) in simulations in domains of size H  4H  H and 8H  8H  H . The domains contain a net vertical field of strength ˇ D 1600, and the resolution is the same in both calculations, 64 grid points per scale height. These results can be compared to Bodo et al. [3], see also Mignone (these proceedings). Although we find significant reduction in the fluctuations in the stress in wide boxes, in agreement with Bodo et al., note that the time-averaged stress is virtually unchanged in both cases. Thus, the radial extent of the domain does not affect the gross properties of the saturated state of the MRI. These, and many other results, will be presented in a forthcoming paper. Currently, a variety of groups are continuing studies of the MRI in shearing boxes to explore, for example, the effect of finite resistivity and viscosity on the saturation amplitude [10, 21], the energetics and dissipation in MRI-driven turbulence [28], and the saturation in radiation dominated disks [14]. One entirely new direction of research is studies of the MRI in very weakly collisional plasmas [27]. A new instability is possible in this regime, driven by anisotropic viscosity. Since in a weakly collisional plasma the mean free path of particles is much larger than their gyroradius, ions are confined to field lines, and viscous transport (which is mediated by ion collisions) can only occur along field lines. In this regime a new instability is possible in weakly magnetized Keplerian flows, the magneto-viscous instability (MVI;[16]). The mechanism of the MVI is

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Fig. 4 Volume average of the total stress in a local shearing box with net vertical flux in a domain with radial extent of H (dashed line), and 8H (solid line)

0.4

stress/P0

0.3

0.2

0.1

0

0

20

40

60

80

100

8

10

t (orbits)

–4

stress / P0

Fig. 5 Reynolds stress in 3D shearing box simulations of the MVI, with Reynolds number Re D Cs H= D 103 , and magnetic Prandtl numbers Pm D = of 10 (dotted line), 5 (dashed line), and 2 (long dashed line). The solid line has  D 0. The line segment shows the expected growth rate from a linear analysis. Saturation is independent of Pm , at ˛  104

–6

–8

–10 0

2

4

6

t (orbits)

identical to the MRI (see Fig. 1), except viscosity rather than magnetic tension transports angular momentum along field lines. The MVI will be important in very diffuse, and very weakly magnetized plasmas: one obvious application is to magnetic field amplification in proto-galaxies. Even if the primordial magnetic field is too weak to drive the MRI, the MVI may play a role in the dynamics, and may even be partially responsible for the dynamo action leading to G fields observed in galactic disks today. We have been studying the MVI in the shearing box, an early result is shown in Fig. 5.

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4 The MRI in Protoplanetary Disks Since HH jets and protostellar outflows arise from weakly magnetized protoplanetary disks, it is of interest to study the MRI in non-ideal MHD. To study the effects of interest, we first write out the generalized Ohm’s Law E D v  B 

.J  B/  B 4J J  B  C c ne e c i

(1)

where  is the Ohmic resistivity, ne the electron number density,  the ion-neutral collision rate, and i the mass density of ions. There are four terms, the first is the usual induction term of ideal MHD, the second represents Ohmic dissipation, the third the Hall effect, and the fourth is ambipolar diffusion (AD). Which term dominates in a protoplanetary disk? The answer depends on the radius and vertical location in the disk. The Ohmic term dominates in the densest inner regions of the disk, the AD term in the very diffuse regions, and the Hall everywhere else. Comprehensive studies of the MRI in realistic models of protoplanetary disks have been given by Salmeron and Wardle [25], Wardle [37], and Terquem [34]; see also Salmeron (these proceedings). These calculations typically adopt the minimum mass solar nebula as a background state, include non-thermal sources of ionization (especially x-rays from the central star), and use a multi-species chemistry model including the effect of grains. Generally the MRI is present in protoplanetary disks, but the linear modes are strongly affected, and thus the saturation may be different. With grains, the MRI is suppressed at the midplane, and column density in the disk unstable to the MRI is reduced to 1% or less of the total. Such layered disks were first discussed by Gammie [11]. The nonlinear regime of the MRI has been studied with each of these nonideal processes, for example Ohmic dissipation was studied in Fleming and Stone [8], ambipolar diffusion in Stone and Hawley [30], and the Hall effect in Sano and Stone [26]. The structure of layered disks was studied using a fixed vertical profile of Ohmic resistivity by Fleming and Stone [9]. More recently, these calculations have been repeated [36, 35, 15] using a non-equilibrium ionization and recombination method, in which several ion species are followed self-consistently along with the MHD flow. A space-time plot of a typical result in shown in Fig. 6, note how the dead zone shrinks in this case.

5 Global Simulations of the MRI The production of jets and winds, the subject of this meeting, is best studied in global simulations. There has been much work (e.g. Romanova, these proceedings) on the production of outflows in the star-disk interaction region. However, these calculations do not yet capture the MRI, the process which drives accretion in the disk, although such calculations should be available very soon. Fully global

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–4 0

50

100

4 –4

100

150

200

4

–4

Maxwell Stress (dyn / cm2)

Height / H

4

10–3 10–6 0 –10–6 –10–3 –1

200

250

300 Time (years)

Fig. 6 Time evolution of the vertical profile of the horizontally-averaged Maxwell stress in a 3D shearing box simulation of the MRI in a layered disk, taken from Turner and Sano [35]. A color version of this plot can be found in the electronic version of this book and in Appendix A (Fig. A.7)

models of MRI unstable protoplanetary disks have been reported by Nelson and Papaloizou [24], as part of a continuing study of planet migration in turbulent disks. Perhaps the most thorough study of the MRI in global disks is for black hole accretion flows. Although such studies are not immediately relevant to the protostellar case, there are common lessons to be learned. Interestingly, powerful jets are produced in global simulations of accretion onto rotating black holes [7]. Much of the thrust for the jet comes from the Blandford-Znajek effect. There is only a weak outflow produced by magneto-centrifugal processes in the disk. Since the Blandford-Znajek effect cannot play a role in protostellar disks, not all jets can be powered by exactly the same physics. Nevertheless, the fact that global simulations of black hole disks can resolve the internal dynamics of the disk, and generate outflows, is promising for the future of similar calculations for protoplanetary disks.

6 Summary The time-dependent MHD of protoplanetary disks is extremely complex and much work is required before it will be fully understood. There are various MHD instabilities that control the dynamics, such as the MRI, Parker instability, RTI, and KHI. Realistic simulations that study the nonlinear regime of these instabilities, including the appropriate non-ideal MHD and microphysics, are warranted. Global simulations of accretion flows around black holes have revealed that powerful jets and outflows are produced. However, similar simulations of MRI unstable disks around protostars have yet to show jets and outflows. Perhaps a different initial field geometry is needed (very few simulations use vertical fields with net flux), or perhaps better thermodynamics is needed (heating in the corona may be necessary to launch the wind – most models assume an isothermal gas), or perhaps the

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non-ideal MHD itself is not yet realistic (global models including the Hall effect – the most important non-ideal MHD process, have yet to be studied). At the moment, there is much effort in studying the global evolution of protostellar disks in which the angular momentum transport and turbulent viscosity driven by the MRI are parametrized by an anomalous viscosity and resistivity. However, it is not clear that 3D MHD turbulence can be characterized in this way. Fully global, 3D simulations of an MRI unstable non-ideal MHD disk, including interaction with the central star, would be an enormous step forward. Acknowledgements I am grateful to the organizers for the invitation to speak at a wonderful conference in spectacular settings. I thank my recent collaborators S. Balbus, T. Gardiner, J. Hawley, T. Sano, Y. Shen, and N. Turner for their contributions. This work was supported by grants from the DOE, NASA ATP, and computational facilities supported by NSF grant AST-0722479.

References 1. Balbus, S.A. 2001. ApJ 562, 909. 2. Balbus, S.A., & Hawley, J.F. 1991. ApJ 376, 214. 3. Bodo, G., Mignone, A., Cattaneo, F., Rossi, P., & Ferrari, A. 2008. A&A 487, 1. 4. Blaes, O., & Balbus, S.A. 1994. ApJ 421, 163. 5. Blaes, O., & Socrates, A. 2001. ApJ 553, 987. 6. Chiang, E. 2008. ApJ 675, 1549. 7. De Villiers, J.P., Hawley, J.F., Krolik, J.H., & Hirose, S. 2005. ApJ 620, 878. 8. Fleming, T.P., & Stone, J.M. 2000. ApJ 530, 464. 9. Fleming, T.P., & Stone, J.M. 2003. ApJ 585, 908. 10. Fromang, S., Papaloizou, J., Lesur, G., & Heinemann, T. 2007. A&A 476, 1123. 11. Gammie, C.F. 1996. ApJ 457, 355. 12. Gammie, C.F. 1998. MNRAS 297, 929. 13. Hawley, J.F., Balbus, S.A., & Stone, J.M. 2001. ApJ 554, L49. 14. Hirose, S., Krolik, J., & Blaes, O. 2008. arXiv0809.1708. 15. Ilgner, M., & Nelson, R.P. 2008. A&A 483, 815. 16. Islam, T., & Balbus, S.A. 2005. ApJ 633, 328. 17. Ji, H., Burin, M., Schartman, E., & Goodman, J. 2006. Nature 444, 343. 18. Johnson, B.M., Guan, X., & Gammie, C.F. 2008. ApJS 174, 145. 19. Johansen, A., & Youdin, A. 2007. ApJ 662, 627. 20. Klahr, H.H., & Bodenheimer, P. 2003. ApJ 582, 869. 21. Lesur, G., & Longaretti, P.-Y. 2007. A&A 378, 1471. 22. Lovelace, R.V.E., Li, H., Colgate, S.A., & Nelson, A.F. 1999. ApJ 513, 805. 23. Masset, F. 2000. A&A Suppl. Ser. 141, 165. 24. Nelson, R.P., & Papaloizou, J.C.B. 2004. MNRAS 350, 849. 25. Salmeron, R., & Wardle, M. 2005. MNRAS 361, 45. 26. Sano, T., & Stone, J.M. 2002. ApJ 577, 534. 27. Sharma, P., Hammett, G.W., Quataert, E., & Stone, J.M. 2006. ApJ 637, 952. 28. Simon, J.A., & Hawley, J.F. 2008. ApJ in press. 29. Stone, J.M., & Norman, M.L. 1994. ApJ 433, 746. 30. Stone, J.M., & Hawley, J.F. 1998. ApJ 501, 758. 31. Stone, J.M., & Gardiner, T.A. 2008. In preparation. 32. Stone, J.M., Gardiner, T.A., Teuben, P., Hawley, J.F., & Simon, J. 2008. ApJS 178, 137. 33. Shibata, K., & Uchida, Y. 1986. PASJ 38, 631. 34. Terquem, C. 2008. arXiv:0808.3897.

66 35. Turner, N.J., & Sano, T. 2008. ApJ 679, L131. 36. Turner, N.J., Sano, T., & Dziourkevitch, N. 2007. ApJ 659, 729. 37. Wardle, M. 2007. Ap&SS 311, 35. 38. Youdin, A., & Shu, F.H. 2002. ApJ 580, 494. 39. Youdin, A., & Johansen, A. 2007. ApJ 662, 613.

J.M. Stone

Theory of Wind-Driving Protostellar Disks Arieh K¨onigl

Abstract Observations reveal a strong link between disk accretion and energetic bipolar outflows in protostars. This connection could reflect the high efficiency of angular momentum transport by centrifugally driven winds launched from the disk surfaces along a large-scale, ordered magnetic field that threads the disk. Vertical angular momentum transport can also be achieved through torsional Alfv´en waves that propagate along the magnetic field lines into the ambient medium. There is observational evidence for a dynamically significant “open” magnetic field in molecular cloud cores, and one can show that when a core collapses to form the central star and disk the magnetic field lines are advected inward by the inflowing gas. It is possible to construct semi-analytic equilibrium models of wind-driving disks with a realistic ionization and conductivity structure and investigate their stability properties. These studies are being complemented by non-ideal–MHD numerical simulations.

A. K¨onigl () University of Chicago, Chicago, IL 60637, U.S.A. e-mail: [email protected]

K. Tsinganos et al. (eds.), Protostellar Jets in Context, Astrophysics and Space Science Proceedings, DOI 10.1007/978-3-642-00576-3 8, c Springer-Verlag Berlin Heidelberg 2009 

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1 Introduction Rotationally supported circumstellar disks evidently originate in the collapse of self-gravitating, rotating, molecular cloud cores [45]. The cores can transfer angular momentum to the ambient gas by magnetic braking (the launching of torsional Alfv´en waves along interstellar magnetic field lines that thread the cloud), and this mechanism also acts to align their angular momentum vectors with the local largescale magnetic field [32]. For a sufficiently strong magnetic field this alignment occurs on a dynamical time scale and hence can be achieved even in cores whose lifetimes are not much longer than that (as in certain models of the turbulent ISM [13]). The dynamical collapse might occur as a result of mass rearrangement in the core on the ambipolar diffusion time [33] or sooner if the core is close to the critical mass for collapse (the effective Jeans mass) from the start [14]. The presence of a large-scale, ordered, and dynamically significant magnetic field has been indicated by the hourglass field morphology detected by polarization measurements in several molecular cloud cores on sub-parsec scales [43, 15], as well as by Zeeman measurements [51] and (indirectly) by the predominantly oblate inferred intrinsic shapes of prestellar cores [48]. The interstellar field lines that thread the core are expected to be advected inward by the infalling gas when gravitational collapse commences (see Sect. 3), and they could naturally give rise to a large-scale, “open” field in the protostellar disk that forms around the protostar; alternative possibilities for a field of this morphology in the disk are a disk dynamo [50,53,4,52] or, in the vicinity of the protostar, stellar field lines that have penetrated the disk (as in the X-wind scenario [47]). It is believed that most of the mass assembled in protostars reaches them through a disk [7, 3]. In order for the disk to accrete, it must have an angularmomentum transport mechanism. Classical disk models have considered radial transport through the plane of the disk: either by a gravitational torque associated with nonaxisymmetric density perturbations [12], which operates in sufficiently < 1 massive disk regions in which the Toomre Q parameter Q  C =G˙ is  (where C is the isothermal sound speed, is the epicyclic frequency, G is the gravitational constant, and ˙ is the disk surface density), or through turbulence induced by the magnetorotational instability (MRI [1]) and associated with a small-scale, disordered magnetic field. However, when the disk is threaded by a large-scale, ordered magnetic field, vertical angular momentum transport through the disk surfaces could potentially be very efficient if the field is strong enough. Besides the magnetic braking mechanism mentioned above, in which an external mass load exerts a back torque on the disk, angular momentum could also be transported vertically by means of a centrifugally driven wind (CDW [5]), in which the requisite inertia is provided by mass uplifted from the disk itself. There are strong indications that the ubiquitous outflows in protostellar systems are intimately linked to the accretion process. This has motivated the construction of models in which vertical angular momentum transport plays a key role in the formation of protostellar disks and in their early evolution.

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2 The Disk–Wind Connection A common feature of accreting protostellar systems is their association with energetic bipolar outflows that propagate along the rotation axis of the source [2]. In T Tauri stars (similar to the early Sun), strong correlations have been found between the presence of outflow signatures (P-Cygni line profiles, forbidden line emission, thermal radio radiation, well-developed molecular lobes) and accretion diagnostics (UV, IR, and millimeter emission excesses, inverse P-Cygni line profiles) [16]. Such correlations evidently extend smoothly to protostars with masses of 10 Mˇ [10]. A related finding is that the apparent decline in outflow activity with stellar age follows a similar trend exhibited by disk frequency and inferred mass accretion rate [7, 6]. In addition, correlations of the type MP / Lqbol (with q  0:6–0.7) have been found in both low-Lbol and high-Lbol protostars for mass accretion rates and for mass outflow rates in ionized jets as well in bipolar molecular lobes [24]. Furthermore, T Tauri-like accretion and outflow phenomena have now been detected also in very-low-mass stars and brown dwarfs [31]. These findings suggest that outflows are powered by accretion and that the same basic physical mechanism operates in both low- (down to nearly the planetary mass limit) and intermediate-mass protostars, and possibly in some higher-mass objects as well [44]. Furthermore, the common occurrence of these outflows makes it seem plausible that they are a key ingredient of the accretion process in these systems, providing an efficient means of transporting away the excess angular momentum and of tapping the liberated gravitational potential energy of the accreted matter [20, 37]. In this connection it is worth recalling the strong indications of a disk origin of the outflows that accompany FU Orionis outbursts in rapidly accreting young protostars [18] and the suggestion [17] that most of the mass accumulation and ejection in low-mass protostars occurs during recurrent outbursts of this type. Interestingly, estimates of the ratio of the mass outflow and mass accretion rate during the outbursts, MP w =MP a  0:1, are comparable to the values estimated during the subsequent, quiescent T Tauri phase of such systems [23, 36, 38]. This value is consistent with disk models in which a CDW transports a significant fraction of the angular momentum of the accreted matter at its base. These models imply MP w =MP a  .r0 =rA /2 and often yield rA =r0  3–10, where rA is the (cylindrical) radius of the Alfv´en critical surface of the wind (which constitutes the effective lever arm for the back torque acting on the disk) and r0 is the launching radius. It is also noteworthy in this context that there are observational indications of a strong magnetic field in the inner disk regions of FU Ori [11], which are consistent with theoretical expectations that only magnetic stresses would be powerful enough to launch the massive disk outflows inferred in these outbursts (MP w  105 Mˇ yr1 ).

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3 Formation of Magnetically Threaded Disks As noted in Sect. 1, the interstellar magnetic field threading the natal cloud core is a likely source of a large-scale “open” field in the protostellar disk that forms when the core collapses. However, if the effective disk viscosity and magnetic diffusivity have a similar origin (turbulence) and hence magnitude, the diffusivity will be able to resist the advection by the inflowing matter and the field will not be carried inward [29]. This potential caveat can be eliminated if the field diffusivity derives from a different physical process than the angular momentum transport.1 This was demonstrated in [22] using semi-analytic r  t similarity solutions that attributed the diffusivity to ambipolar diffusion (AD, the drift of ions, electrons, and magnetic field relative to neutrals in a weakly ionized gas) and the angular momentum transport to magnetic braking. The fiducial solutions presented in [22] correspond to relatively slowly rotating cloud cores that are subject to moderately strong braking and have comparatively high magnetic diffusivities. These configurations assumed the following structure during the collapse, moving from larger to smaller radii:  An outer region of ideal-MHD infall.  An AD shock (resolved as a continuous transition).  An infall region dominated by the protostar’s gravity, within which AD is “revi-

talized” and the field largely decouples from the matter.  A centrifugal shock at the outer boundary of the rotationally supported disk.  A Keplerian accretion disk whose structure is well described by an asymptotic

analytic solution. In particular, the disk is characterized by MP in .r/ D const and, at any given time, by a magnetic field amplitude B.r/ / r 5=4 and by a surface field inclination br;s  Br;s =Bz  4=3 (where the subscript “s” denotes the disk surface). It is interesting to compare the value of br;s in the derived disk p solution with the launching condition for a CDW from a Keplerian disk, br;s > = 3 [5]. This shows that a disk formed from the collapse of a magnetized cloud core could launch a CDW. It is also noteworthy that the scaling B / r 5=4 inferred in the r  t similarity solution of [22] at any given time t is the same as the one obtained in the steady-state, r  z self-similar solution of [5]. This suggests that a CDW could be naturally incorporated into the collapse solution and that, in reality, it might complement magnetic braking in regulating the angular momentum transfer from the collapsing core. While the conclusions of the above semi-analytic model have been supported by numerical simulations employing 3D, nested-grid, resistive-MHD codes [49, 26], some doubts have been expressed in the literature about the validity of this picture. In particular, it was argued that magnetic braking in collapsing cores (modeled as singular isothermal toroids) could be so strong that no disk would form at all – a

1

For other suggestions on how to get around this conundrum, see [46, 39].

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conclusion previously reached by similar calculations carried out within the framework of ideal MHD – and that AD would likely be too weak to mitigate this effect [28]. These arguments have, however, been based on simulations in which the core was taken to be magnetically supercritical from the start, in contrast with the calculations reported in [22], which considered a scenario in which the core had evolved into a state of gravitational collapse from an earlier subcritical phase. In the latter case, the core at first contracts along the large-scale magnetic field lines, and by the time the collapse starts there is little mass left at high latitudes to exert a significant back torque on the equatorially concentrated inflow. In contrast, a core that starts out being supercritical is less affected by magnetic stresses, resulting in a significant mass accumulation near the center even as a comparatively massive envelope is still present and provides the inertia for efficient magnetic braking. While the ADdominated models presented in [22] were also able to produce diskless protostars, this was attained in cases in which the ratio of the sound speed in the disk to the Alfv´en speed in the external medium (a measure of the effective inertia acting to brake the disk) was unusually high or when jbs j  jB;s j=Bz (related to the magnitude of the magnetic torque at the disk surface) was allowed to become  1 (which is probably physically unlikely), but not for typical parameter values. Although it is still not established which of these two pictures better describes the situation in real systems, recent observational [48, 51] and numerical [34] studies appear to be consistent with the notion that the core settles into an oblate mass configuration before the start of dynamical collapse, in which case magnetic braking would be unlikely to prevent the formation of protostellar disks in the regime where AD is the main field diffusivity mechanism.

4 Equilibrium Structure and Stability of Disk/Wind Systems A simplified equilibrium model of weakly ionized protostellar disks that transport angular momentum by means of a CDW was presented in [55]. In this model, the disk is in the AD diffusivity regime (although the effects of the Hall current were also considered) and all the angular momentum of the accreted gas is carried away by the wind. The disk is further assumed to be geometrically thin, in near-Keplerian rotation, threaded by an “open” magnetic field, and to have a vertically uniform temperature and ion density. As a further simplification, only a radially localized region (r r) of the diffusive disk is considered, although it is matched to a global (self-similar) ideal-MHD wind. Despite the various approximations, this model captures the essential physical properties of such systems. In particular, the qualitative characteristics remain unchanged when the radially localized disk solution is generalized to a global self-similar configuration encompassing both the disk and the wind [25]. Furthermore, solutions with similar properties are obtained when the disk is in the Hall or Ohm diffusivity regimes [42] and when the full conductivity tensor is used in conjunction with a realistic ionization profile [40].

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The relevant parameters of the disk configuration are:  B  VB =C : normalized drift velocity of the poloidal magnetic flux surfaces

(labeled by the magnitude  of the flux function), which in the AD regime is equal to the midplane ion/electron radial velocity;    ni =˝K : ratio of the neutral–ion collision frequency and the Keplerian frequency, a measure of the strength of the neutral–ion (and hence, in the ADregime, field–matter) coupling;2  a  VA0 =C : ratio of the midplane Alfv´en speed and the isothermal sound speed, a measure of the magnetic field strength, whereas the model wind parameters are:    l=.˝K r02 /: normalized specific angular momentum (including both kinetic

and magnetic contributions); P w =d  : normalized mass/magnetic flux ratio, a measure of   .r0 ˝K =B0 /d M the degree of mass loading of the wind. Three model parameters are determined by conditions outside the disk:  by the sonic critical-point constraint;  jbs j D .  1/ by the Alfv´en critical-point constraint;  brs (which is related to the parameter a) by the magnetic flux distribution along

the disk [35].3 Figure 1 shows the loci of self-consistent disk/wind equilibrium solutions in the wind parameter space. To better understand the physical basis for the parameter constraints on viable solutions, it is useful to employ the hydrostatic approximation (i.e., the limit Vz ! 0) to the radially localized disk model [55, 42]. In the AD regime, one finds p <   < VK =2C;
, disk remains sub-Keplerian throughout; p , wind launching condition (br;s > 1= 3); , top of disk (zs ) > density scale height (h); , Joule heating < gravitational energy release.

> 1, whereas 2 Ineqaulities 1 and 2 imply that the coupling parameter  must be  and 3 imply that magnetic squeezing dominates the vertical confinement of the disk (h= hT  a= < 1, where hT  C =˝K is the tidal scale height). These inequalities can also be used to demonstrate that the minimum wavelength of the most unstable linear MRI mode exceeds the disk scale height (so that, by and large, a wind-driving disk is immune to MRI) and to identify the region in the disk that is susceptible to

2

Outside the pure-AD regime, this parameter is generalized by the Elsasser number  [42]. So far, none of the disk/wind models that appeared in the literature has incorporated all of the relevant constraints outside the disk; for example, [55] treated B as a free parameter rather than using the condition on brs to fix it. 3

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MRI (essentially the region where inequality 1 is violated), which in turn makes it possible to construct “hybrid” disk models in which both radial and vertical angular momentum transport mechanisms operate (possibly dominating in different vertical segments at the same radial location [41]). Analogous results to (1) can be derived also for the Hall and Ohm conductivity regimes [42]. The equilibrium disk/wind configurations considered above naturally lie in a stability “window” in which the magnetic field is strong enough to largely suppress the MRI (as already noted) but not so strong as to be subject to the radial interchange instability [21]. It was nevertheless suggested that such disks might be inherently unstable because of the expected increase in MP w =MP in with increasing brs [30, 8]. One can, however, argue that the generic turning point in the equilibrium curves indicates that they must also possess a stable branch, which can be identified p as the lower branch in the curves plotted in Fig. 1 [19]. Along this branch, as brs  2=a increases when a goes down,  increases and hence MP w =MP in  1=Œ4.  1/ goes down: the excess angular momentum is removed by an increase in the effective lever arm (the Alfv´en radius; rA =r0 D 1=2 ) rather than by a higher MP w =MP in . Non-ideal MHD simulations, which have already confirmed the basic equilibrium structure implied by the semi-analytic models [9, 27, 56], would be the best tool for testing these inferences about the stability properties of wind-driving disks. In the equilibrium solutions considered above, the field–matter coupling parameter  was >1 throughout the vertical extent of the disk. Such disks can be termed strongly coupled. In contrast, in weakly coupled disks the coupling parameter

Fig. 1 Mapping of the wind-driving disk solutions (for several values of the field–matter coupling parameter ) onto the self-similar wind solution space (defined by the values of the mass-loading parameter and angular-momentum parameter ). The field-strength parameter a increases on moving counterclockwise along a given curve

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(which, more generally, is given by the Elsasser number ) is <1 near the midplane >1 near the disk surface [25, 54]. Whereas in strongly coupled and increases to  < disks VA;0  C , j < Vr > j  C (where <> denotes a mass-weighted vertical average), and Br;s > jB;s j (with Br increasing already at z D 0), weakly coupled configurations are characterized by VA;0 C , j < Vr > j C , and Br;s < jB;s j (with Br taking off only when  increases above 1; .dBr =dB /0 D 2 ). An interesting implication of these results is that even in weakly coupled regions of the disk, where neither MRI nor wind-driving occurs, angular momentum can still be transported vertically because the magnetic torque on the disk is / Bz dB =d z and jB j=Bz can be comparatively large even if Br =Bz is still small. This could have implications to the issue of “dead zones” in protostellar disks.

5 Conclusion The main results reviewed in this contribution can be summarized as follows:  There is strong observational evidence for a disk–wind connection in protostars.









Large-scale, ordered magnetic fields have been implicated theoretically as the most likely driving mechanism of the observed winds and jets. Centrifugally driven winds (CDWs) are an efficient means of angular momentum transport in disks and could dominate the transport in regions where MP w =MP in  0:1 (or even less): the wind launching region in FU Orionis objects could be a case in point. This efficiency may be responsible, at least in part, for the ubiquity of collimated, energetic outflows in these sources. A natural origin for a large-scale, ordered magnetic field in a protostellar disk is the interstellar field that threads the parent molecular cloud core. Vertical angular momentum transport along the field (by torsional Alfv´en waves or a CDW) typically leads to inward advection of the field by the accretion flow despite the high diffusivity of the gas. The structure of wind-driving disks can be modeled under realistic ionization and conductivity conditions using semi-analytic techniques. Further progress is being made by numerical simulations that employ non-ideal MHD codes. The relevant parameter ranges in the different conductivity regimes can be delineated with the help of algebraic expressions obtained in the hydrostatic approximation. The derived disk configurations can be divided into strongly coupled ones, in which j < Vr > j  C , and weakly coupled ones, in which j < Vr > j C . One can also construct models that incorporate both vertical angular momentum transport by a large-scale mean field and (in regions where the field is comparatively weak) radial transport by MRI-induced turbulence. Wind-driving protostellar disks are inferred to have a stable equilibrium branch and, on the whole, to be immune to MRI and to magnetic interchange. The results of time-dependent simulations are consistent with basic stability, but further studies with MHD codes that have non-ideal, 3-dimensional, and heating/cooling capabilities are needed.

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In closing, it is worth emphasizing that wind-driving disks have distinct properties with potentially important observational consequences, some of which might be useful for testing the models. These have not been reviewed here for lack of space.

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Aspect Ratio Dependence in Magnetorotational Instability Shearing Box Simulations Andrea Mignone, Attilio Ferrari, Gianluigi Bodo, Paola Rossi, and Fausto Cattaneo

Abstract Three-dimensional numerical simulations of the magnetorotational instability in the shearing box approximation with a nonzero net flux are presented. By changing the size of the computational domain in the radial direction relative to the vertical box height, we find, in agreement with previous studies, that transport of angular momentum (associated with the so-called “channel solution”) is strongly intermittent and maximized for boxes of unit aspect ratio. On the other hand, in boxes with larger aspect ratio the intermittent behavior disappears and angular momentum transport is inhibited.

A. Mignone () and A. Ferrari DFG, University of Turin, via P. Giuria 1, 10125 Torino, Italy e-mail: [email protected] G. Bodo and P. Rossi INAF/Osservatorio Astronomico di Torino, via Osservatorio 20, 10025 Pino Torinese, Italy F. Cattaneo Department of Astronomy and Astrophysics, The University of Chicago, 5640 S. Ellis ave., Chicago IL 60637, USA

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1 Motivations It is commonly accepted that the magnetorotational instability (MRI, [2]) may provide a viable driving mechanism for turbulence in accretion disks and that it may account for the required rate of angular momentum transport [13]. Much of what is presently known about MRI driven turbulence relies on a local model of a differentially rotating system, the so-called “shearing-box”, where a small periodic patch of the disk is considered (see, e.g., [5, 11, 12, 1, 14, 7]). The advantage of a local approach (as opposed to a global disk simulation) is the possibility of reaching higher resolutions at the same computational cost. The model is based on a local expansion of the equations in a reference frame corotating with the disk at some fiducial radius R0 [6]. The validity of the model should be verified a posteriori, by checking, for example, that the properties of the solutions do not depend on the size and geometrical properties of the computational domain. A critical discussion of the validity of the shearing box approximation can be found in [10]. With the present study we investigate one aspect of the self consistency of the shearing box results, namely, the dependence on the box aspect ratio. In the MRI, the characteristic length scale is set by the vertical wavelength of the mode with maximum growth rate which in turn depends on the magnetic field strength. Typically, the vertical size of the computational domain is chosen to contain some multiple of that length whereas there is no general guidance in choosing the width in the radial and azimuthal directions. Practitioners in the field commonly adopt boxes with aspect ratio of unity, although this issue has not received careful consideration and there has been no systematic study of the dependence of the solutions on the aspect ratio. Here and in [3], we present a systematic study of shearing box results as a function of the aspect ratio by performing a series of 3D, compressible, isothermal numerical simulations in the shearing box approximation. In particular, we want to see if the results observed in boxes of unit aspect ratio are representative of more extended systems, or if they display peculiarities induced by an overly constrained geometry.

2 Results Simulations with unit aspect ratio show that the transport of angular momentum has an intermittent behavior with episodes of enhanced transport. During these states During these states, loosely referred to as “channel” solutions, velocity and magnetic perturbations are highly spatially correlated. It is now believed that the formation of near channel solutions and their subsequent disruption by the parasitic instabilities of [4] are the basis for the observed intermittent behavior [11]. Note that, although the channel solutions are, strictly speaking, exponentially growing exact solutions of the incompressible shearing box equations [4], the same terminology

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is used to describe behaviors with similar properties to that observed in the exact solutions. For the present purpose, we consider a Cartesian box corotating with the disk at the local angular velocity ˝. The box size is L41 in the x (radial), y (azimuthal) and z (vertical) directions, respectively, with L varying from 1 to 8. We take 2=˝ as the unit of time and we fix the sound speed cs and the plasma ˇ D pgas =pmag to the values of 4.56 and 104 , respectively. With these choices the fastest growing MRI mode has a vertical wavelength close to 1/3. Computations are carried out at both low and high resolutions (32 and 128 zones per unit length, respectively) on equally spaced grids. The interested reader is referred to [3] for more details. The MHD equations are solved in conservative form using the isothermal MHD module available in the PLUTO code [8, 9]. Figure 1 shows the volume averaged Maxwell stresses hwxy i D hBx By i as a function of time at low resolution. With the chosen normalization, the values shown in the figure correspond to the ˛ parameter of [13]. A striking difference appears if we compare the L D 1 case with the other two having larger aspect ratios in which the spikes are absent. This behavior is quantified by the respective probability distribution functions shown in Fig. 2. The case L D 1 is quite distinct from the other two with a long tail corresponding to the episodes of enhanced transport. However only a small tail is present in the case with L D 4 and absent for L D 8. Due to the presence of larger peaks, the time-averaged value of ˛ for L D 1 exceeds the other two cases by

Fig. 1 Time histories of the Maxwell stresses averaged over the computational box and normalized to the pressure for the low resolution simulations. The three panels correspond to cases with L D 1, 4, and 8. Note that time is in units of the rotation time

Fig. 2 Probability distribution functions of the Maxwell stresses for the same cases as in Fig. 1. The three curves correspond to the three values of the aspect ratio: L D 1—solid curve, L D 4— dashed curve, and L D 8—dot-dashed

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approximately a factor of 2. It is natural to assume that the distinctive behavior of the former case is due to the presence of channel solutions that are otherwise absent in the simulations with larger aspect ratios. This assumption can be verified by looking for typical imprints of the channel solutions, e.g., a high correlation coefficient between directional components of the magnetic field, or a specific relationship between the vertical wavenumber and the angle between the magnetic field and the azimuthal direction. Figure 3 introduces a scatter diagram whereby for each gridpoint the value of By is plotted as a function of Bx (or similarly for the velocities). The two panels on the left show examples of such distributions for L D 1 corresponding, respectively, to a maximum and a minimum of hwxy i. In the transition from the minimum to the maximum the magnetic field fluctuations increase their intensity and the x and y components tend to become more correlated. The velocity fluctuations show a similar behavior, with almost no correlation in the minimum state. In contrast, for L D 4, there are no significant variations in the distributions during the evolution, see Fig. 3.

Fig. 3 Scatter diagrams of horizontal magnetic fluctuations (dark tones) and velocities (lighter tones). The first panel corresponds to the case with L D 1 at an instant near a maximum in the stress. The second panel corresponds to the same case near a minimum. The third panel corresponds to a representative instant for the case with L D 4

Fig. 4 Probability distribution functions for the azimuthal component of magnetic field By . The left panel refers to the case with L D 1. Solid and dashed lines correspond to the first two panels in Fig. 3 while the dot-dashed line corresponds to the exact channel solution. The right panel corresponds to the third panel in Fig. 3

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Another property of the distributions is that, while for L D 4 and at the minimum of the L D 1 case the distributions peak at zero value, at the maximum of the L D 1 case there are two peaks around the largest values of the fluctuations. This is shown in Fig. 4 where we plot the probability distribution function of the azimuthal magnetic field component By . The panels show the behavior at an instant near a maximum in the stress and a minimum for two different aspect ratios. In the left panel, we also show the corresponding distribution for the channel solution. The presence of the double-peaked distribution in the case L D 1 can be considered again as an indication of the presence of the channel solution around the maxima in the stress, as it can be seen by comparing the solid and the dot-dashed lines.

3 Conclusions In this paper we have investigated the effects of the box aspect ratio in compressible 3D numerical simulations of the MRI in the shearing box approximation. For aspect ratios close to unity the solutions are strongly affected by the channel solutions, with frequent episodes of high correlation and efficient transport. For larger aspect ratios the channel solutions disappear, the system remains in the state with lower correlation and the average transport of angular momentum is correspondingly reduced. Further increase in aspect ratio does not lead to any significant changes. Thus we conclude that the less correlated state is more likely to be representative of the extended system. Whatever the reason for the differences, clearly, the obvious conclusion is that studies of turbulence driven by the MRI with net flux should be conducted in shearing boxes sufficiently large to allow the solution to develop naturally. The exact domain size depends on the particular problem and the quantities under consideration. However, it seems clear that boxes with aspect ratios close to unity over-emphasize the role of the channel solutions and may lead to misleading results. These results do not put in doubt the fact that MRI may be adequate to produce the turbulence necessary to support the required “viscosity” to transport angular momentum. However, the uncertainties discussed by many authors about the applicability of the shearing box approximation for the astrophysical problem of accretion disks suggest that the extension of the shearing box results to the full disk has to be tested on global simulations with sufficient resolution.

References 1. Balbus, S. A.: 2003, Ann. Rev. Astron. Astrophys., 41, 555 2. Balbus, S. A., & Hawley, J. F.: 1991, ApJ, 376, 214 3. Bodo, G., Mignone, A., Cattaneo, F., Rossi, P., & Ferrari, A.: 2008, Astronomy & Astrophysics, 487, 1

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4. Goodman, J., & Xu, G.: 1994, ApJ, 432, 213 5. Hawley, J. F., Gammie, C. F., & Balbus, S. A. 1995, ApJ, 440, 742 6. Hill, G. W. 1878, Amer. J. Math., 1, 5 7. Lesur, G. & Longaretti, P.-Y., 2007, MNRAS, 378, 1471L 8. Mignone, A., Bodo, G., Massaglia, S., Matsakos, T., Tesileanu, O., Zanni, C., & Ferrari, A. 2007, ApJs, 170, 228 9. Mignone, A. 2007, J. Comp. Phys., 225, 1427 10. Regev, O., Umurhan, O.M. 2007, Astronomy & Astrophysics in press (ArXiv Astrophysics e-prints, arXiv:0711.0794) 11. Sano, T. & Inutsuka, S. 2001, ApJ, 561, L179 12. Sano, T. & Stone, J. M. 2002, ApJ, 570, 314 13. Shakura, N. I., & Sunyaev, R. A.: 1973, Astronomy & Astrophysics, 24, 337 14. Turner, N. J., Stone, J. M, Krolik, J. H. & Sano, T. 2003, ApJ, 593, 992

Advection/Diffusion of Large Scale Magnetic Field in Accretion Disks Richard V.E. Lovelace, David M. Rothstein, and Gennady S. Bisnovatyi-Kogan

Abstract Winds and jets of proto-stellar systems are thought to arise from disk accretion involving (1) a small-scale turbulent magnetic field in the disk (due to the magneto-rotational instability or MRI) and (2) a large-scale magnetic field which gives rise to the winds and/or jets. An important problem with this picture is that the enhanced viscosity is accompanied by an enhanced magnetic diffusivity which acts to prevent the build up of a significant large-scale field. Recent work has pointed out that the surface layers of the disk are non-turbulent and thus highly conducting (or non-diffusive). This is because the MRI is suppressed in the surface layers where the magnetic and radiation pressures are larger than the thermal pressure. Here, we calculate the vertical (z) profiles of the stationary accretion flows (with

R.V.E. Lovelace () and D.M. Rothstein Department of Astronomy, Cornell University, Ithaca, NY 14853, USA e-mail: [email protected]; [email protected] G.S. Bisnovatyi-Kogan Space Research Institute, Russian Academy of Sciences, Moscow, Russia e-mail: [email protected]

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radial and azimuthal components), and the profiles of the large-scale, magnetic field taking into account the turbulent viscosity and diffusivity due to the MRI and the fact that the turbulence vanishes at the surface of the disk. We derive a sixth-order differential equation for the radial flow velocity vr .z/ which depends mainly on the midplane thermal to magnetic pressure ratio ˇ > 1 and the magnetic Prandtl number of the turbulence P D viscosity/diffusivity. Boundary conditions at the disk surfaces take into account possible magnetic winds or jets and allow for a surface current flow in the highly conducting surface layers. The stationary solutions we find indicate that a weak (ˇ > 1) large-scale field does not diffuse away as suggested by earlier work.

1 Introduction Early work on disk accretion to a black hole argued that a large-scale poloidal magnetic field originating from say the interstellar medium, would be dragged inward and greatly compressed near the black hole by the accreting plasma [2] and that this would be important for the formation of jets [6]. Later, the importance of a weak small-scale magnetic field within the disk was recognized as the source of the turbulent viscosity of disk owing to the magneto-rotational instability (MRI; [1]). Analysis of the diffusion and advection of a large-scale field in a disk with a turbulent viscosity comparable to the turbulent magnetic diffusivity (as suggested by MRI simulations) indicated that a weak large-scale field would diffuse outward rapidly [14, 9, 7]. This has led to the suggestion that special conditions (nonaxisymmetry) are required for the field to be advected inward [13]. Recently, Bisnovatyi-Kogan and Lovelace [3] pointed out that the disk’s surface layers are highly conducting (or non-diffusive) because the MRI is suppressed in this region where the magnetic energy-density is larger than the thermal energy-density. Rothstein and Lovelace [11] analyzed this problem in further detail and discussed the connections with global and shearing box magnetohydrodynamic (MHD) simulations of the MRI. In this work we calculate the profiles through the disk of stationary accretion flows (with radial and azimuthal components), and the profiles of a large-scale, weak magnetic field taking into account the turbulent viscosity and diffusivity due to the MRI and the fact that the turbulence vanishes at the surface of the disk. A full explanation of this work will be given elsewhere [8]. Related calculations of the disk structure were done earlier by K¨onigl [4], Li [5], Ogilvie and Livio [10] but without taking into account the absence of turbulence at the disk’s surface.

2 Theory We consider the non-ideal magnetohydrodynamics of a thin axisymmetric, viscous, resistive disk threaded by a large-scale dipole-symmetry magnetic field B. We use a cylindrical .r; ; z/ inertial coordinate system in which the time-averaged magnetic

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field is B D Br rO C B φO C Bz zO , and the time-averaged flow velocity is v D vr rO C v φO C vz zO . The main equations are

dv 1 D rp C g C J  B C F ; dt c @B D r.v  B/  r  .r  B/: @t

(1) (2)

These equations are supplemented by the continuity equation, by r  B D 4J=c, and by r B D 0. Here,  is the magnetic diffusivity, F D r T  is the viscous force with Tjk D  .@vj =@xk C @vk =@xj  .2=3/ıj k r v/ (in Cartesian coordinates), and  is the kinematic viscosity. We assume that both the viscosity and the diffusivity are due to magnetorotational (MRI) turbulence in the disk so that  D P D ˛

2 cs0 g.z/; ˝K

(3)

where P is the magnetic Prandtl number of the turbulence assumed a constant of order unity, ˛  1 is the dimensionless Shakura-Sunyaev [12] parameter, cs0 is the midplane isothermal sound speed, ˝K  .GM=r 3 /1=2 is the Keplerian angular velocity of the disk, and M is the mass of the central object. The function g.z/ acconds for the absence of turbulence in the surface layer of the disk [3, 11]. In the body of the disk g D 1, whereas at the surface of the disk, at say zS , g tends over a short distance to a very small value  108 , effectively zero, which is the ratio of the Spitzer diffusivity of the disk’s surface layer to the turbulent diffusivity of the body of the disk. We consider stationary solutions of (1) and (2) for a weak large-scale magnetic field. These can be greatly simplified for thin disks where the disk half-thickness, of the order of h  cs0 =˝K , is much less than r. Thus we have the small parameter " D h=r D cs0 =vK 1: In the following we use the dimensionless height   z= h. The three magnetic field components are assumed to be of comparable magnitude on the disk’s surface, but Br D 0 D B on the midplane. On the other hand the axial magnetic field changes by only a small amount going from the midplane to the surface, Bz  "Br Bz (from r B D 0) so that Bz  const inside the disk. As a consequence, the @Bj =@r terms in the magnetic force in (1) can all be dropped in favor of the @Bj =@z terms (with j D r; ). The velocity components are 2 assumed to satisfy v2z cs0 and v2r v2 . Consequently, v .r; z/ is close in value to the Keplerian value vK .r/  .GM=r/1=2 . Thus, @v =@r D .1=2/.v =r/ to a good approximation. With these assumptions, the radial component of (1) gives  @br ˇ Q  @ D 1  kp "2  u2 C ˛ 2 ˇ @ " @

  @ur

g Q ; @

(4)

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where Q  .r; z/= 0 with 0  .r; z D 0/. The midplane plasma beta is 2 =B02 ; where kp  @ ln p=@ ln r is assumed of order unity and ˇ  4 0 cs0 2 2 =v2A0 , where vA0 D B0 =.4 0 /1=2 is the midplane p D cs . Note that ˇ D cs0 Alfv´en velocity. The rough condition for the MRI instability and the associated turbulence in the disk is ˇ > 1 [1]. In the following we assume ˇ > 1, which we refer to as a weak magnetic field. We normalize the field components by B0 D Bz .r; z D 0/, with br D Br =B0 , b D B =B0 , and bz D Bz =B0  1. Also, u  v .r; z/=vK .r/ and the accretion speed ur   vr =.˛cs0 /. For the assumed dipole field symmetry, br and b are odd functions of  whereas ur and u are even functions. In a similar way one can derive an equation for @b =@ from the toroidal component of (1). The zcomponent of (1) corresponds to hydrostatic equilibrium for ˇ > 1. Equations for @u =@ and @br =@ follow from (2). Integration of the @b =@ equation from  D 0 (where b D 0 and @u =@ D 0) to the exterior of the disk (SC where g D 0) gives the average accretion speed, ur D u0 

2bSC ; ˛ˇ ˙Q

(5)

2 2 r = h/ which is the sum of a viscous contribution, u0  3"k (with k  @ ln. cs0 =@ ln.r/ > 0 of order unity), and a magnetic contribution (/ bSC ) due to the loss of angular momentum from the surface of the disk where necessarily bSC  0 R R [7]. Here ur  0 S d  u Q r =˙Q , ˙Q  0 S d  , Q and the S C subscript indicates evaluation outside the disk. A similar integration of the @br =@ eqaution implies R that brS D PS hur i, where h::i D 0 S d .::/=S . The equations for ur ; u ; br ; & b can be combined to give a single equation for ur ./,

     @ 1 @ @ur @

g Q

g Q g @ @ Q @ @    2  @ ur @ @

g Q  ˛ 2 ˇP 2 g @ @ @ g Q   2  @ur 1 @ @

g Q  ˛ 2 ˇP 2 @

Q @ @   2  2    ur 2 2 @ 2 @

g Q ur  gu0 C P C˛ ˇ @ 2 @ 2 g Q u r D0: C 3ˇP 2 g

˛4ˇ2

@2 @ 2

(6)

This equation is integrated from  D 0 out to the surface of the disk S where the boundary conditions apply. Because ur is an even function of , only ur .0/, u00r .0/, and uiv r .0/ need to be adjusted in order to satisfy the boundary conditions at the disk surface [8].

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disk surface

1

ζ = z /h

Bp vr

0

–1 vr –6

–4

–2

0

2

4

r/h Fig. 1 Radial flow speed vr D ur (normalized to ˛cs0 ) as a function of  D z= h and a sample poloidal .Br ; Bz / magnetic field line for ˇ D 102 and P D 1

Figure 1 shows a sample solution of (6) for " D 0:05, ˛ D 0:1, ˇ D 100, and P D 1 where we find ur =u0 D 1:30, bSC D 0:321, and brS D 0:276 (discussed more fully in [8]). Note that the midplane region of the disk flows radially outward while the regions closer to the disk’s surfaces flows radially inward. The stationary solutions we find indicate that a weak (ˇ > 1) large-scale field does not diffuse away as suggested by earlier work. Acknowledgements R.V.E.L. was supported in part by NASA grant NNX08AH25G and by NSF grants AST-0607135 and AST-0807129. D. M. R. is supported by an NSF Astronomy and Astrophysics Postdoctoral Fellowship under award AST-0602259. G. B.-K. was partially supported by RBFR grants 08-02-00491 and 08-02-90106, and RAN program P-04.

References 1. Balbus, S.A., & Hawley, J.F. 1991, ApJ, 376, 214 2. Bisnovatyi-Kogan, G.S., & Ruzmaikin, A.A. 1974, Ap&SS, 28, 45 – 1976, Ap&SS, 42, 401 3. Bisnovatyi-Kogan, G.S., & Lovelace, R.V.E. 2007, ApJ, 667, L167 4. K¨onigl, A. 1989, ApJ, 342, 208 5. Li, Z.-Y., 1995, ApJ, 444, 848 6. Lovelace, R.V.E. 1976, Nature, 262, 649 7. Lovelace, R.V.E., Romanova, M.M., & Newman, W.I. 1994, ApJ, 437, 136 – 1997, ApJ, 484, 628 8. Lovelace, R.V.E., Rothstein, D.M., & Bisnovatyi-Kogan, G.S. 2009, ApJ, 701, 885 9. Lubow, S.H., Papaloizou, J.C.B., & Pringle, J.E. 1994, MNRAS, 267, 235 10. Ogilvie, G.I., & Livio, M. 2001, ApJ, 553, 158 11. Rothstein, D.M., & Lovelace, R.V.E. 2008, ApJ, 677, 1221

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12. Shakura, N.I., & Sunyaev, R.A. 1973, A&A, 24, 337 13. Spruit, H.C., & Uzdensky, D.A. 2005, ApJ, 629, 960 14. van Ballegooijen, A.A. 1989, in Accretion Disks and Magnetic Fields in Astrophysics, ed. G. Belvedere (Dordrecht: Kluwer), 99

Part III

Jet Launching

Magnetic Reconnection in Accretion Disk Systems: From BHs to YSOs Elisabete M. de Gouveia Dal Pino, Pamela Piovezan, Grzegorz Kowal, and Alex Lazarian

Abstract We investigate the role of violent accretion and magnetic reconnection in different jet/disk accretion astrophysical systems, namely young stellar objects (YSOs), microquasars, and active galactic nuclei (AGNs). In the case of microquasars and AGNs, we find that violent reconnection episodes between the magnetic field lines of the inner disk region (which are established by a turbulent dynamo) and those that are anchored into the black hole are able to heat the coronal/disk gas and accelerate particles to relativistic velocities through a diffusive first-order Fermi-like process within the reconnection site that will produce relativistic blobs or plasmons. The heating of the coronal/disk gas is able to produce a steep X-ray spectrum with

E.M. de Gouveia Dal Pino () IAG-USP, Brazil e-mail: [email protected] P. Piovezan MPA, Germany e-mail: [email protected] G. Kowal and A. Lazarian University of Wisconsin-Madison, USA

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a luminosity that is consistent with the observations and we argue that it is being produced mainly at the foot of the reconnection zone in the disk, while the Fermilike acceleration process within the reconnection site results a synchrotron radio power-law spectrum with a spectral index (˛ D 0:75) that is compatible with that observed during the radio flares in microquasars. In the case of the YSOs, a similar magnetic configuration can be reached. The amount of magnetic energy that can be extracted from the inner disk region can heat the coronal gas to temperatures of the order of 108 K and could explain the observed X-ray flaring emission in some COUP sources.

1 Introduction Jets are observed to emerge from a wide variety of astrophysical objects like young stellar objects (YSOs), galactic black holes (or microquasars), and active galactic nuclei (AGNs). Despite their different physical scales (in size, velocity, and amount of energy transported), they have strong morphological similarities suggesting a common mechanism for their origin [4]. The most accepted model for jet production is the magneto-centrifugal acceleration out off a magnetized accretion disk that surrounds the central source [1]. However, within this scenario, there is a number of open questions, such as the quasi-periodic ejection phenomena which are often associated to these sources [3]. de Gouveia Dal Pino and Lazarian [5] (hereafter GDPL) have proposed, in the context of the MHD scenarios, that these large scale superluminal ejections observed during radio flare events are produced by violent magnetic reconnection (MR) episodes.

2 MR in Microquasars MR in Microquasars to analyze the effect of MR events in the inner zone of microquasars we have considered a rotating stellar mass black hole (BH) surrounded by a magnetized accretion disk. The geometry of the problem is schematized in Fig. 1. A magnetosphere around the central BH may be formed from the drag of magnetic field lines by the accretion disk [15,19]. The disk’s large scale magnetic field can be established by the action of a turbulent dynamo inside the accretion disk [9]. According to the magneto-centrifugal scenario [1], this poloidal magnetic flux summed to the disk differential rotation may generate a wind which removes angular momentum from the system increasing the accretion rate. This will increase the ram pressure of the accreting material that will then push the magnetic lines of the inner disk region towards the lines which are anchored in the BH ergosphere, allowing a violent MR event to occur (see the zone labeled as Helmet Streamer in Fig. 1). To simplify the analysis, we have considered that the inner zone of the accretion disk and the BH are nearly co-rotating in such a way that there is no neat angular

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Fig. 1 Schematic representation of the magnetic field geometry at the inner zone of the source/accretion disk system. From GDPL [5]

momentum and energy transfer between them. The accretion disk is modeled as a geometrically thin and optically thick disk with the opacity parameterized by the alpha-prescription [18]. The corona above and below the disk is modeled according to Liu, Mineshige and Shibata’s description [11]. The accretion rate and the magnetic field strength are related to each other by the equilibrium in the inner disk radius between the magnetic pressure and the ram pressure of the gas. Furthermore, the magnetic field strength can be parameterized by ˇ, the ratio between the magnetic and the total disk pressure. With the assumptions above, the amount of magnetic power released by MR is 19=16 9=16 19=32 25=32 11=16 ˇ0:8 M14 RX;7 l100 erg=s WP B Š 1:6  1035 ˛0:5

(1)

where ˛ D 0:5˛0:5 , ˇ D 0:8ˇ0:8 , M D 14Mˇ M14 is the BH mass, RX  3RS D 3GM=c 2 D 107 RX;7 cm is the inner disk radius, and the size of the magnetic tube flux is l D 100RX l100 cm. As required by the coronal model, the energy released can heat the coronal gas to temperatures of 109 K. Furthermore, GDPL have proposed that this energy can also be used to accelerate particles to relativistic velocities within the reconnection zone, in a mechanism similar to a first-order Fermi process. This acceleration mechanism results in a power-law electron distribution of N.E/ / E 5=2 and the associated synchrotron emission has a spectral index of 0:75, which is in agreement with radio observations [3]. Kowal, de Gouveia Dal Pino and Lazarian are presently testing this acceleration model in reconnection sites with the help of 3D Godunov-MHD simulations combined with a particle in-cell technique [10]. The enhanced x-ray emission that often accompanies the violent flares in microquasars can be easily explained within our scenario as due to the increase in the accretion rate immediately before the violent MR events. The observed soft

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x-ray emission is expected to be a fraction of the accretion power. For an enhanced accretion rate near the Eddington limit this is given by GMBN MP 1 Š 1:87  1039 M14 MP 19 RX;7 erg=s WP ac Š RX

(2)

which is compatible with x-rays observations [17]. On the other hand, the hard xray component that is also often observed can be explained by inverse Compton scattering of the soft x-rays photons by the hot electrons of the corona/jet [14].

3 MR in AGNs/Quasars MR in AGNs/Quasars despite the huge difference in scales, AGNs/quasars and microquasars have similar morphologies [12]. In fact, some studies indicate that the similarity between these systems is more than morphological: they are possibly subjected to the same physical processes. If this is the case, then a generalization of the above scenario to the extragalactic sources is almost straightforward. Considering the same assumptions presented above, we obtain similar scaling relations for AGNs [14]. Figure 2 depicts a synthesis of the generalization of the MR scenario for relativistic sources including both microquasars and AGNs. The diagram shows the calculated magnetic power released in violent MR events as a

Fig. 2 The vertical lines correspond to the magnetic power released by MR for a suitable parameter space. For microquasars: 5Mˇ M 20Mˇ [17], 0:05 ˛ 0:5 [8], 0:1 ˇ 1, 1RX l 1000RX , and MP < MP Ed d . For AGNs: the mass range is: 106 Mod ot M 109 Mod ot . The stars correspond to observed jet radio luminosities

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function of the central source mass. It indicates that the MR could be responsible for the initial particle acceleration and quasi-periodic relativistic blob emission in the case of the microquasars as this process releases power that is more than one order of magnitude greater than the observed radio luminosities. In the case of the AGNs, we find that the most radio luminous quasars require reconnection events with super-Eddington accretion in order to explain the formation of relativistic blobs. This super-Eddington regime would manifest itself only in a short time interval until the partial destruction of the magnetic flux by MR relaxes this extreme configuration.

4 MR in YSOS MR in YSOs differ in many aspects from microquasars. for instance, they produce thermal rather than relativistic jets and exhibit emission lines from which their physical properties (such as density and temperature) are inferred. But they also often exhibit an intense magnetic activity that results in a strong and variable x-ray emission. Observed flares in x-rays are often attributed to magnetic activity at the stellar corona [7]. However, some COUP (Chandra Orion Ultra-deep Project) sources have revealed strong flares that were related to peculiar gigantic magnetic loops linking the magnetosphere of the central star with the inner region of the accretion disk. It has been argued that this x-ray emission could be due to magnetic reconnection in these gigantic loops [6]. In this section, we examine this issue in more detail investigating the role of MR events in the inner disk/corona of YSOs to explain the X-ray flares and check how MR may be related to the thermal jets of these sources. The geometrical configuration of the problem is the same as that of Fig. 1. However, while the inner disk regions of microquasars are dominated by radiation pressure, the inner disk region as of YSOs are dominated by gas pressure. Instead of modeling the disk/corona of these systems, we have used the mean values of coronal density and temperature inferred from observations [6] to parameterize our analysis. The stellar magnetic field is assumed to be dipolar and we also considered the equilibrium between the stellar magnetic pressure and the disk ram pressure in the inner disk region. With these assumptions, the maximum possible accretion rate is reached when the disk touches the stellar surface. Associated to this maximum accretion rate there is a maximum magnetic power released by MR events. Figure 3 shows a comparison between the predicted maximum magnetic power released under these circumstances in violent reconnection events and the observed x-ray luminosities (represented by stars) for a sample of COUP sources [6]. For most sources, the magnetic power is of the order of (or greater than) the observed luminosities if the maximum accretion rate is (0.5–1)104 Mˇ /yr (10 sources) or (2.5–5)104 (5 sources). Such accretion rates of 104 Mˇ /yr are 100–1000 times greater than the mean typical rates for YSOs. However, in the scenario above, such high accretion rates would be required only for very short time intervals during violent MR events [5].

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Fig. 3 Magnetic power released by violent MR events for the COUP sources (see text for details)

This implies that the detection of any spectral signature of these higher accretion rates would be probably very difficult. Finally, the energy released by violent MR processes could help to heat the gas at the base of the jets [16]. A rough estimate indicates that the large amount of magnetic power released by MR events can be thermally conducted up to distances 2 10 AU [13] in a timescale cond  108 n10 T85=2 l10 s that is comparable to the dynamical timescale of the DGTau jet [2], for example. Acknowledgements The authors acknowledge partial support from FAPESP.

References 1. Blandford, R. D., & Payne, D. G. 1982, MNRAS, 199, 883 2. Cerqueira, A. H., & de Gouveia Dal Pino, E. M. 2004, A&A, 426, L25 3. Dhawan, V., Mirabel, I. F., & Rodr´ıguez, L. F. 2000, ApJ, 543, 373 4. de Gouveia Dal Pino, E. M. 2005, Advances in Space Research, 35, 908 5. de Gouveia Dal Pino, E. M., & Lazarian, A. 2005, A&A, 441, 845 6. Favata, F., Flaccomio, E., Reale, F., Micela, G., Sciortino, S., Shang, H., Stassun, K. G., & Feigelson, E. D. 2005, ApJs, 160, 469 7. Feigelson, E. D., & Montmerle, T. 1999, ARA&A, 37, 363 8. King, A. R., Pringle, J. E., & Livio, M. 2007, MNRAS, 376, 1740 9. King, A. R., Pringle, J. E., West, R. G., & Livio, M. 2004, MNRAS, 348, 111 10. Kowal, de Gouveia Dal Pino & Lazarian 2009, in prep. 11. Liu, B. F., Mineshige, S., & Shibata, K. 2002, ApJl, 572, L173 12. Mirabel, I. F., & Rodr´ıguez, L. F. 1998, Nature, 392, 673 13. Pesenti, N., Dougados, C., Cabrit, S., O’Brien, D., Garcia, P., & Ferreira, J. 2003, A&A, 410, 155 14. Piovezan and de Gouveia Dal Pino 2008, in prep. 15. Price, R. H., & Thorne, K. S. 1986, Black Holes: The Membrane Paradigm, 1 16. Ray, T., Dougados, C., Bacciotti, F., Eisl¨offel, J., & Chrysostomou, A. 2007, Protostars and Planets V, 231

Magnetic Reconnection in Accretion Disk Systems: From BHs to YSOs 17. Remillard, R. A., & McClintock, J. E. 2006, ARA&A, 44, 49 18. Shakura, N. I., & Syunyaev, R. A. 1973, A&A, 24, 337 19. Wang, D. X., Xiao, K., & Lei, W. H. 2002, MNRAS, 335, 655

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Self-Collimated Jets from Accretion Discs and Star-disc Interaction Zones Jonathan Ferreira

Abstract I will first briefly address the physical conditions that must prevail within near Keplerian accretion discs in order to steadily launch cold self-collimated jets. The most important condition is the presence of a large scale magnetic field whose vertical component is smaller but close to equipartition with the disc thermal pressure. Depending upon their history, some astrophysical systems may never establish regions within their accretion discs with such a field and will harbor no jet. In a second part, I will present some results on two simple magnetic configurations possibly linking the stellar magnetosphere to the disc and also suspected to give rise to ejection: X-winds and Reconnection X-winds.

1 Introduction Collimated ejection of matter is widely observed in several astrophysical objects: inside our own galaxy from Young Stellar Objects (this conference) and X-ray binaries [31], but also from the core of active galaxies [46]. All these objects share J. Ferreira () Laboratoire d’Astrophysique de Grenoble, BP 53 F-38041 Grenoble Cedex 9, France e-mail: [email protected]

K. Tsinganos et al. (eds.), Protostellar Jets in Context, Astrophysics and Space Science Proceedings, DOI 10.1007/978-3-642-00576-3 12, c Springer-Verlag Berlin Heidelberg 2009 

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the following properties: jets are almost cylindrical in shape; the presence of jets is correlated with an underlying accretion disc surrounding the central mass; the total jet power is a sizeable fraction of the accretion power. Most of observed YSO images show jets that are extremely well collimated already close to the source, with an opening angle of only some degrees. On the other hand, the derived physical conditions show that jets are highly supersonic. Indeed, emission lines require a temperature of order 104 K, hence a sound speed cs  10 km/s while the typical jet velocity is vj  300 km/s. The opening angle  of a ballistic hydrodynamic flow being simply tan  D cs =vj , this provides   5ı for YSOs, nicely compatible with observations. Thus, observed jets could well be ballistic. Note that this qualitative argument is not changed by recent observations of rotation in jets, since the rotation velocity is of the same order than the sound speed. Therefore, the fundamental question is how does a physical system produce an unidirectional supersonic flow ? This implies that confinement must be intimately related to the acceleration process. To date, the only physical process proved to be capable of accelerating plasma along with a self-confinement relies on the action of a large scale magnetic field carried along by the jet, with the so-called “hoop stress”. This field is then anchored onto a rotating object from which rotational energy is tapped: the star or the surrounding accretion disc. As discussed in detail in [17], observed YSO jets are probably a combination of several components, namely a stellar wind surrounded by the disc self-confined wind. However, energetic requirements are such that it is hardly conceivable that stellar winds alone could provide the bulk of the outflowing mass. Instead, disc winds are probably the main mass loss agent. This would then fit into a broad astrophysical picture where jets would be launched whenever accretion discs are present regardless of the central object (black hole or star). However, not all accretion discs do provide jets, a well known and still puzzling fact. Besides, high resolution images as well as recent jet kinematic constraints show that jets arise only from the innermost disc regions. Therefore, there must be something that allows, at least in some systems, a transition from an outer [40] standard accretion disc (hereafter SAD) driving no jets, to an inner jet driving disc (hereafter JED). Ferreira and Petrucci [22] made the conjecture that the SAD-to-JED transition agent is the amplitude of the large scale Bz magnetic field threading the disc. Indeed, it has been shown that a field close to equipartition is necessary in order to launch jets from JEDs [20, 11]. In the outer SAD regions, the Bz component can be advected inwards so that the magnetization D Bz2 = o P , where P is the disc thermal pressure, increases towards the center (see [22] and contributions from Lovelace, Murphy et al, this conference). Depending on the system, a situation may arise where innermost regions develop values of close to unity. Therefore, not all accretion discs should drive jets. There is an observational evidence that tends to confirm this conjecture [29]. Indeed, all CTTS discs are randomly oriented with respect to the local large scale (cloud) Bz field, hinting to the fact that magnetic forces are irrelevant in defining the plane of the disc. However, sources with no detected jets have discs that lie on

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the same plane as the field, whereas sources with powerful jets are perpendicular to it. It is therefore tempting to attribute the jet launching phenomenon to the capability of the disc to build up a region with a high enough magnetization (  1), by dragging in and amplifying the interstellar Bz field.

2 Disc Winds The global accretion-ejection mechanism is characterized by four steps: 1. The rotation of the disc, embedded in a large scale Bz magnetic field, generates R an electromotive force e D dr˝rBz which drives a radial electric current Rh I D 0 d z2 rJr ; 2. This current flows amid the jet and makes a closed electric circuit: its leakage through the magnetic surfaces provides positive magnetic forces, azimuthal (F , giving rise to a centrifugal force) and parallel to the magnetic surfaces (Fk ); 3. Once the total current I vanishes or when it starts to flow parallel to the magnetic surfaces, no magnetic acceleration arises anymore and the ejected plasma reaches an asymptotic state; 4. At a given altitude, the radial distribution of the current I.r; z/ determines the collimation properties of the jet: collimation and acceleration are thus tightly related. Studies of jets generally assume boundary conditions at the surface of Keplerian discs, either with analytical (e.g. [3, 45]) or numerical (e.g. [35, 34, 1]) approaches. It turns out that jets are quite easily launched ! But to assess the physical conditions allowing this to occur, one needs to compute the disc as well and take into account in particular the feedback of the jets on the disc internal structure: mass loss and angular momentum transport. In a standard accretion disc (i.e. steady-state, no mass loss), the accretion rate is a constant. In a Jet Emitting Disc (JED) it cannot be anymore constant and one defines the ejection efficiency  such that MP a .r/ / r 

(1)

Mass conservation1 provides then the ratio of ejection to accretion rates, namely , from a JED of inner and outer radii ri n and rout . f D 2MP j =MP a '  ln rrout in Observations show that this ratio varies between 0.01 and 0.1, which implies similar values for the ejection efficiency . A complete theory must then provide the allowed values of  as a function of the disc properties. This was done only within the framework of self-similar models [11, 6, 7]. In practice, one must take into account the full 2D problem and not treat the disc as infinitely thin as in a standard disc theory.

A local definition is possible, namely  D d ln MP a =d ln r, but accretion disc theory relies on quantities varying slowly, so that a constant  is a good first order approximation.

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This requires to treat the physics of both accretion and ejection and address the mass loading issue. Given the lack of space very few details can be given here. An interested reader will find more informations in a number of reviews (see for instance [12, 37, 13] and references therein). The most important aspects are the following: 1. The inward advection of poloidal magnetic field due to accretion must be balanced by some anomalous magnetic diffusivity m in order to maintain a steady-state. An alpha prescription, m D ˛m VA h has been used, as a general approach valid for all astrophysical systems [18]. 2. The magnetic torque must spin up material at the disc surface for a jet to be launched. The torque must therefore be positive at these layers (and remain so throughout the jet). As the magnetic field is actually spinning down the disc material, a magnetic acceleration implies a change of sign of F D Jz Br  Jr Bz on a disc scale height. If this is not fulfilled, the large scale magnetic field will continue to brake down the disc material and no jet is allowed. In a thin accretion disc, this is only possible if the radial current density Jr decreases on the same scale [19, 20]. 3. At the disc surface, where the resistive disc-ideal MHD jet transition operates, the ejected material must become supersonic. In a cold ejection process, the magnetic field is the only agent making a transition from radial to vertical motion. The only possibility for the ejected material to become supersonic is to already have an almost sonic accretion speed. This is totally impossible in a SAD where the sonic Mach number ms D ur =Cs ' ˛v h=r 1. In a JED, this number becomes of order unity thanks to near equipartition (  1) fields [20]. These results were all found in self-similar studies of accretion-ejection systems but can be recovered using physical arguments (Ferreira & Casse 2008, submitted). As a result of equipartition fields, the magnetic torque dominates over the viscous torque and most of the released accretion power feeds the jets. The final velocity of the jets then depends on the fraction of the mass that p has been ejected, namely  [18]. Indeed, the asymptotic velocity is V1 D VK 2  3 where VK is the Keplerian speed at the anchoring radius and  ' 1 C 1=2 is the magnetic lever arm, a jet parameter first introduced by [3]. In quasi-Keplerian discs the enthalpy carried away by the outflowing mass at the disc surface is a priori negligible and magnetized jets would thus be “cold” flows. The use of isothermal magnetic surfaces [11] or adiabatic ones [6] introduces no significant change. All solutions display Alfv´en surfaces located quite far from the disc (at zA  rA ), with typical ejection efficiencies   0:01. They all recollimate towards the axis because of a dominant “hoop-stress” and terminate with a shock at the location of the FM point (which is uncrossable for these solutions). Note however that all solutions are super-FM in the conventional sense (i.e. up > VF M , see [11] for more details). Note that meridional self-similar jets with a different asymptotic behavior were obtained but for magnetic flux distributions incompatible with a Keplerian accretion disc [9, 33].

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Warmer and denser jets can be produced from quasi-Keplerian discs if the outflowing gas temperature undergoes a sudden rise above the disc. This can be done in astrophysical systems by two ways: (a) stellar X-rays or UVs from the accretion shock will illuminate and possibly photo-ionize the surface layers of the disc; (b) local dissipation of accretion power itself in the highly turbulent magnetized corona (or more correctly chromosphere) expected to be present above the turbulent disc. This last possibility is actually suggested by both observations [24] and numerical simulations [30]. Including the energy equation and using a prescription for additional heat deposition, [7] showed that enhancing the temperature at the disc surface layers has dramatic effects. Figure (1) shows one such solution. For example, the disc vertical equilibrium can be changed so that a balance can now be achieved with magnetic configurations much more bent. As a result, much smaller values of , down to 0.001, can be obtained. On the other extreme, providing a large enthalpy allows more mass to be loaded onto the field lines: these thermally and magnetically driven jets can accelerate up to   0:1 (jet parameter  1). Jets 3 to 5 times slower but denser than in the “cold” case can be obtained. These solutions have been successfully tested against observations [36].

Fig. 1 Typical super-FM disc wind with  D 0:03;  D 0:03 (h D r). Density, pressure and temperature are normalized to their value at the disc midplane, the magnetic field components to Bz .z D 0/ and the velocities to the Keplerian speed at the anchoring radius ro . All magnetic field components remain comparable from the disk surface to the Alfv´en point. Note that the density profile inside the disc, where both ur and uz are negative, is very different from a gaussian. Recollimation takes place at z ' 3 103 ro . In such a solution 84% of the released accretion power feeds the jets, the remaining (16%) being radiated at the disc surfaces [14]

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3 Winds from the Star-disc Interaction Zone It has been argued that YSO jets could be driven by the rotational energy of the rotating star [42, 41, 39, 27]. This is a very tempting idea since it would also explain why T Tauri stars rotate at about 10% of their breakup speed [4, 38]. In these scenarios, there is no need for a magnetic field in the disc and, if there is some, it plays no strong role in driving the jets. X-winds [41] or Accretion-Powered Stellar Winds [27] invoke a Y-type magnetic interaction as pictured in Fig. 2. A X-type magnetic neutral line is formed only if the two fields cancel each other at some radius, which is a situation depicted in Reconnection X-winds [21]. But, contrary to the previous models, these last winds require the presence of an outer disc wind. The funny thing is that both X-type and Y-type configurations give birth to a magnetic neutral line at the equator, which is good for chondrule formation [44, 23]. To date, there is no complete model for each of these three cases: X-winds do not incorporate yet the disc and probably suffer from a major deficiency in mass loss capability (see below); Stellar winds must rely on some unknown but highly efficient energy transfer mechanism (see [17, 28]); Dynamical calculations remain to be done for Reconnection X-winds.

Fig. 2 Two simple magnetic star-disc interactions: Y-type (left) where the stellar magnetic moment is anti-parallel to the disc magnetic field and X-type (right), where they are parallel to each other (adapted from [17])

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3.1 A Y-type Configuration: X-winds The theory of X-winds is a mixture of cold super-Alfv´enic wind calculations plus a global picture or scenario (see [5] and references therein). This scenario assumes that about f D 1=3 of the accreted mass is actually ejected by the X-winds. Until now this is a working hypothesis as no computation nor analytical argument proved that such a high ejection efficiency could be achieved. Another assumption is the locus of the X-wind: it is assumed that the wind is launched from a tiny region around a radius rX of extent rX rX , with rX close to the corotation radius rco . Both assumptions (locus and mass loss) are actually required if one desires to construct an angular momentum balance between accretion and ejection onto the star [41]. Let us examine the consequences of a huge mass loss ratio f D 1=3 on the underlying discR physics. In steady state, the total power carried by the wind is 2Pjet;X D 2 E. / up dS where E. / is the Bernoulli invariant and the integration is made over the disc surface at the base of the wind. It can also ki n / where Pki n is the flux of thermal be written 2Pjet;X D 2PMHD .1 C PPMHD and mechanical (kinetic and gravitational) energy advected by the matter, whereas PMHD is the R MHD Poynting flux emitted from the Bdisc surface. This flux writes PMHD D SMHD dS where SMHD D ˝F rB po is the Poynting vector and ˝F the angular velocity of the field line. The total jet power is then

2Pjet;X D 2PMHD

HQ  1 1C 2.  1/

!

  2PMHD ' Pacc

2q ms



rX  rX ' Pacc rX 1 C  rX (2)

where   J. /  3 in X-wind terminology and HQ is the enthapy normalized 2 2 rX =2 (taking the same notations as [43]). In the cold wind approximation to ˝X HQ 1 and the dominant term is the MHD Poynting flux, as assumed in all Xwind calculations done so far. The two rhs expressions appearing in (2) are valid whatever the origin of this flux of MHD power (accretion and/or stellar rotational B C Bo

energy). They were obtained making use of 2PMHD ' 2˝X rX o 2 rX rX and Pacc D GM MP a =2rX is the released accretion power. In this expression q ' BC =Bo is the shear of the magnetic field at the disc surface and the sonic Mach number ms D ur =Cs is provided by the disc angular momentum equation. In a near Keplerian steady-state accretion disc it writes ms D 2q .1 C 1 /, where  is the ratio of the magnetic (due to the X-wind) to the viscous (turbulent) torque [6]. Thus, it appears that the power carried initially by a cold X-wind is only a tiny fraction of the released accretion power Pacc . In the most favorableqcase with

  1, the average asymptotic jet speed would be only Vjet;X ' VX rXf=rX , 2 =2. X-wind models are expected to arise from where we used Pjet;X D MP X Vjet;X a region of length rX =rX  h=r  0:05 [41, 5]. The resulting speeds are then far too low in order to explain observed jets, unless f is substantially smaller than

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the assumed canonical value of 1/3. This result might be surprising at first sight. Indeed, X-wind models published in the literature do obtain fast flows! There is therefore a discrepancy between published X-wind calculations and the global scenario. The reason of this discrepancy is subtle and lies in the value of the toroidal field at the disc surface. While it has been freely specified in X-wind flow calculations2 , arguments based on magnetic flux conservation put a much severe constraint on it (Ferreira & Casse 2009, submitted to ApJ). And the larger the B of course, the larger the power feeding the jets. This critical question deserves a thorough analyses as it conditions the capability of X-winds to describe YSO jets.

3.2 A X-type Configuration: Reconnection X-winds To my knowledge, the Reconnection X-wind model [21] is the only one that addresses the issue of stellar angular momentum removal from embedded sources. In this model, it is assumed that the interstellar magnetic field is advected with the infalling material in such a way that a significant magnetic flux ˚ is now threading the protostellar core and the inner disc regions (as simulations show, [2, 25]). This self-gravitating core will develop a dynamo of some kind but whose outcome is assumed to be the generation of a dipole field with a magnetic moment parallel to the disc magnetic field (Fig. 3). This is clearly an assumption as there is no theory of such a constrained dynamo that takes into account both the presence of an initial strong fossil field and the outer disc (see however [32]). The coexistence of this dipolar stellar field with the outer disc field generates an X-type magnetic neutral line where both fields cancel each other at a radius rX . From the point of view of the disc, nothing is changed beyond rX : a disc wind is taking place in the JED and disc material accretes by loosing its angular momentum in the jets. At rX however, magnetic reconnection converts closed stellar field lines and open disc field lines into open stellar field lines. Accreting material that

1 Ω Altitude Z (AU)

0.5

0 rm rx –0.5

–1

–1

0 Cylindrical radius R (AU)

1

Fig. 3 Star-disc interaction where the stellar magnetic moment is parallel to the disc magnetic field. There are three distinct types of ejection: a stellar wind on the axis, a disc wind (shown in colors) and intermittent bullets launched at the interface (Reconnection X-wind), braking down the protostar and channeled by the outer disc wind. A color version of this plot can be found in the electronic version of this book and in Appendix A (Fig. A.8)

2

X-wind models cross only the Alfv´en point, thus B remains free at the base of the wind.

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was already at the disc surface at rX is now loaded into these newly opened field lines (there is a strong upward Lorentz force above rX ). Since these lines are now rotating at the stellar rotation rate, they exert a strong azimuthal force that drives ejection. This is a new type of wind: although material is ejected along field lines anchored onto the star, this is not a stellar wind since material did not reach the stellar surface and thus did not loose its rotational energy. Moreover, they are fed with disc material but mostly powered by the stellar rotational energy. As a consequence, Reconnection X-winds exert a negative torque on the protostar which leads to a stellar spin down. On the other hand, an increase of the stellar angular velocity ˝ is expected from both accretion and contraction, with a typical Kelvin-Helmoltz time scale of several 105 yrs. Because of the huge stellar inertia, the evolution of ˝ with time must be followed on these long time scales. One assumption used to compute the angular momentum history of the protostar on those scales is that rX ' rco . Such an assumption relies on the possibility for the protostellar magnetosphere to evacuate angular momentum through violent ejection events (Reconnection X-winds) whenever rX > rco , while quasi steady accretion columns form when rX < rco . Consistently with rX ' rco , a constant fraction f D MP X =MP a is assumed on these long time scales, where MP X is the ejected mass flux in Reconnection X-winds, as well as a constant magnetic lever arm parameter . These winds are therefore best seen as violent outbursts carrying disc material (blobs?) and stellar angular momentum from the star-disc interaction, channeled and confined by the outer disc wind. Note that a conventional stellar wind would of course take place and fill in the inner field lines with mass, but its effect on the stellar spin evolution has been neglected in this work. As the protostar is being spun down, the co-rotation radius rco increases and so must rX . The stellar dipole field is assumed to follow Bstar D B .r=R /n where the index n describes a deviation from a pure dipole in vacuum. Now, rX is defined by the cancellation of the stellar and disc field, whose scaling is imposed 5=4 by the conditions prevailing in the outer JED, namely Bd isc / M1=4 MP a1=2 rX . The only way to ensure rX ' rco on these long time scales is then to decrease MP a in time as well. Note that this is not a surprise as the accretion rate onto the star is controlled by the star-disc interaction. Thus, while computing the stellar spin evolution in time ˝ .t /, starting from conditions prevailing in Class 0 objects and using f ,  and n as free parameters, one gets also R .t /, M .t / and MP a .t /. It was found that all low-mass Class 0 objects can indeed be spun down, from the break-up speed to about 10% of it, on a time scale consistent with the duration of the embedded phase for very reasonable values of the free parameters (n D 3 or 4, f  > 0:1, see Fig. 4). Stellar period, mass, radius and disc accretion rates were found consistent with values for T Tauri stars with a dipole field smaller than 1 kG.

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Fig. 4 Time evolution of disc accretion rate, protostellar period, mass and radius as function of n for f D 0:1 and  D 3:n D 3 (solid), n D 3:41 (dashed), n D 3:87 (dotted), n D 4:4 (dashdotted) and n D 5 (long-dashed). For these reasonable values of the parameters, a very significant braking is obtained in only a few 105 yrs [21]

4 Conclusion The theory of self-confined steady jet production from Keplerian accretion discs has been completed in the framework of “alpha” discs. The physical conditions required to thermo-magnetically drive jets have been constrained. The investigation of the observational appearance of JEDs in the innermost regions of circumstellar accretion discs has just began but provides promising results [10, 26, 8]. The role of large scale magnetic fields in discs has gradually emerged and it seems now an unavoidable ingredient of star formation theory as a whole. The amount of magnetic flux ˚ in the disc is an unknown parameter but it is reasonable to assume that it scales with the total mass M . If this is verified then two important aspects could be naturally explained:

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1. Reconnection X-winds can brake down a protostar during the embedded phase, explaining why T Tauri stars rotate at about 10% of the break-up speed. Remarkably, the mystery of the low dispersion in angular velocities would be naturally accounted by a low dispersion in the ratio ˚=M [21]. These winds are actually time dependent massive bullets channeled by the outer steady disc wind, hence good candidates for explaining HH objects. 2. The transition from Classes 0, I and maybe II (with both disc winds and Reconnection X-winds) to Classes II and III (with stellar winds) would follow the evolution of the disc magnetic flux ˚, with a transition radius between the outer SAD and the inner JED decreasing in time.

References 1. Anderson, J. M., Li, Z.-Y., Krasnopolsky, R., & Blandford, R. D. 2005, ApJ, 630, 945 2. Banerjee, R. & Pudritz, R. E. 2006, ApJ, 641, 949 3. Blandford, R. D. & Payne, D. G. 1982, MNRAS, 199, 883 4. Bouvier, J., Wichmann, R., Grankin, K., et al. 1997, A&A, 318, 495 5. Cai, M. J., Shang, H., Lin, H.-H., & Shu, F. H. 2008, ApJ, 672, 489 6. Casse, F. & Ferreira, J. 2000a, A&A, 353, 1115 7. Casse, F. & Ferreira, J. 2000b, A&A, 361, 1178 8. Combet, C. & Ferreira, J. 2008, A&A, 479, 481 9. Contopoulos, J. & Lovelace, R. V. E. 1994, ApJ, 429, 139 10. Donati, J.-F., Paletou, F., Bouvier, J., & Ferreira, J. 2005, Nature, 438, 466 11. Ferreira, J. 1997, A&A, 319, 340 12. Ferreira, J. 2002, in “Star Formation and the Physics of Young Stars”, J. Bouvier and J.-P. Zahn (eds), EAS Publications Series, astro-ph/0311621, 3, 229 13. Ferreira, J. 2008, New Astronomy Review, 52, 42 14. Ferreira, J. & Casse, F. 2004, ApJ, 601, L139 15. Ferreira, J. & Casse, F. 2008, submitted 16. Ferreira, J. & Casse, F. 2009, submitted 17. Ferreira, J., Dougados, C., & Cabrit, S. 2006a, A&A, 453, 785 18. Ferreira, J. & Pelletier, G. 1993a, A&A, 276, 625 19. Ferreira, J. & Pelletier, G. 1993b, A&A, 276, 637 20. Ferreira, J. & Pelletier, G. 1995, A&A, 295, 807 21. Ferreira, J., Pelletier, G., & Appl, S. 2000, MNRAS, 312, 387 22. Ferreira, J., Petrucci, P.-O., Henri, G., Saug´e, L., & Pelletier, G. 2006b, A&A, 447, 813 23. Gounelle, M., Shu, F. H., Shang, H., et al. 2006, ApJ, 640, 1163 24. Kwan, J. 1997, ApJ, 489, 284 25. Machida, M. N., Inutsuka, S.-i., & Matsumoto, T. 2006, ApJ, 647, L151 26. Masset, F. S., Morbidelli, A., Crida, A., & Ferreira, J. 2006, ApJ, 642, 478 27. Matt, S. & Pudritz, R. E. 2005, ApJ, 632, L135 28. Matt, S. & Pudritz, R. E. 2008, ArXiv e-prints, 801 29. M´enard, F. & Duchˆene, G. 2004, A&A, 425, 973 30. Miller, K. A. & Stone, J. M. 2000, ApJ, 534, 398 31. Mirabel, I. F. & Rodr´ıguez, L. F. 1999, ARA&A, 37, 409 32. Moss, D. 2004, A&A, 414, 1065 33. Ostriker, E. C. 1997, ApJ, 486, 291 34. Ouyed, R., Clarke, D. A., & Pudritz, R. E. 2003, ApJ, 582, 292 35. Ouyed, R. & Pudritz, R. E. 1997, ApJ, 482, 712 36. Pesenti, N., Dougados, C., Cabrit, S., et al. 2004, A&A, 416, L9

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Large-Scale 3D Simulations of Protostellar Jets Jan Staff, Kai Cai, Brian Niebergal, Rachid Ouyed, and Ralph Pudritz

J. Staff () Department of Physics, Purdue University 525 Northwestern Avenue West Lafayette, IN 479072036, USA, and Department of Physics and Astronomy, Louisiana State University, 202 Nicholson Hall, Tower Dr., Baton Rouge, LA 70803-4001, USA e-mail: [email protected] K. Cai Department of Physics, McMaster University & Astronomy, ABB-241, 1280 Main St. W, Hamilton, ON, L8S 4M1, Canada and Department of Physics, University of Notre Dame, 225 Nieuwland Science Hall, Notre Dame, IN 46556-5670 USA e-mail: [email protected] B. Niebergal and R. Ouyed Department of Physics and Astronomy, University of Calgary, SB 605, 2500 University Drive NW, Calgary, Alberta, Canada, T2N 1N4, e-mail: [email protected]; [email protected] R. Pudritz Department of Physics & Astronomy, McMaster University, ABB-241, McMaster University, 1280 Main St. W, Hamilton, ON, L8S 4M1, Canada e-mail: [email protected]

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Abstract Results of three-dimensional time-dependent magneto-hydrodynamic simulations of T Tauri jets, launched magneto-centrifugally from the surface of Keplerian disk is presented. We extend the calculations to scales probed by HST ( 50 AU) thereby allowing a direct comparison between simulations and observations. We explore the effects of different initial magnetic field configurations on evolution and stability of these jets and focus on comparing the generated (poloidal and azimuthal) velocity profile maps with observed velocity structures.

1 Introduction Astrophysical jets are observed to be associated with young stellar objects surrounded by an accretion disk. Accretion disks threaded by large-scale magnetic field lines will have their surface layer stripped away by centrifugal stresses acting along the field lines. The outflow is then magneto-centrifugally accelerated and is collimated by hoop stresses that arises from toroidal magnetic fields developed within the jet itself (see [6] for a recent review). In this paper we present three-dimensional MHD simulations of protostellar jets launched from accretion disks. The goal of this work is to use the HST observations of jets observed in forbidden lines to confront simulations of disk winds in order to deduce information about the hydromagnetic drive. The simulation box is 60 AU along the jet axis and 30 AU wide (in the other two direction). Jets are observed down to 14 AU scales [1, 2], explaining the size of our simulation box. The work presented here is work in progress, and a more detailed paper is on the way. The layout of this paper is as follows: In Sect. 2 we describe the simulation setup, followed by our results in Sect. 3 and a summary in Sect. 4.

2 Simulation Setup A version of the ZEUS-MP code (a parallel version of the Zeus code [7]) is adapted to simulate MHD jets launched from the surface of a Keplerian disk. The setup closely follows that in [4] (hereafter OCP) for which we refer the reader for details. We do however use a different initial magnetic field configuration and the simulation box is much bigger than in OCP. The code solves the ideal time dependent MHD equations. As in OCP, we assume a polytropic equation of state with  D 5=3. The grid is in Cartesian coordinates (x, y, z). The z-axis is taken to be the axis of the disk and the jet when it is launched. The disk is a Keplerian disk (not resolved in the z direction in the simulations) that lies along the x-y plane, at z D 0. The z D 0 boundary is an inflow boundary, all other boundaries are outflow boundaries. The initial magnetic field configuration is current-free and is set following (10) in [3] with D 0:01 to mimic the ’potential’ configuration used in [5] (see the paper by Cai et al. in these proceedings for a discussion of the case D 0:25).

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This setup is in perfect hydrostatic balance. The magnetic field in the corona continues into the disk. Hence there is no toroidal magnetic field in the disk. Our goal with the simulations is to reach observational HST scales. The smallest scale that can be observed with HST is 14 AU [1] which explains the size of our simulation box (60 AU30 AU30 AU). In order to achieve this, we use a uniform grid of 300300800 zones, corresponding to 75ri (about 2 AU; 1ri is the radius of the inner edge of the disk) on either side of the axis and 400ri (12 AU) along the axis. Outside of this uniform grid is a ratioed grid. The total grid consist of 5005001536 zones, covering 250ri (15 AU) on either side of the axis and 2000ri (60 AU) along the axis.

3 Results The simulation has currently reached 1,500 rotations of the inner disk, and we here report on this work in progress. Figure 1 shows the density structure and the poloidal magnetic field in a cut straight through the middle of the grid after 1,500 rotations of the inner disk. The density inside the bow shock fluctuates within three orders of magnitude. The density in the bow shock remains roughly constant about 21017 g/cm3 . The higher densities are found in the jet, it ranges up to 51016 g/cm3 . The poloidal component of the magnetic field shows a straight “backbone” along the center of the jet out to about 300ri . Further away from the star some circular features are seen on either side of the axis, indicating a spiral structure spiraling in and out of the plane of view. The units of the magnetic field is in Gauss, and the field has been normalized assuming a field of 10 G at the inner edge of the accretion disk. The poloidal field varies from 0.1 to 1 G in the straight part of the backbone and drops somewhat when it enters the spiral phase.

3.1 Emission Line Maps ˚ We have made a first attempt on constructing emission line maps of OI 6,300 A from the simulation. By assuming a polytropic equation of state, we assume that temperature is proportional to density to the power of   1. The origin of the emission of these forbidden lines is a matter of debate. In our 3D simulations, it is the dense network of internal shocks that is responsible for the emission seen in Fig. 2. If the density is above quenching density, emission along the line of sight from behind this dense region is suppressed. We show in Fig. 2 how these emission line maps for the OI line would depend on exposure time of the observing telescope (assuming that the telescope can observe with a resolution similar to that in the simulation). The shortest exposure only sees the jet. In Fig. 3 the radial velocity (the radial velocity is the velocity along the jet) as measured when looking at the OI emission line is shown. This radial velocity is

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Fig. 1 top: Snapshot of log. / from the simulation at t D 1500. The entire bow shock is still well within the simulation box. Notice the rapid variability in density inside the bow shock. A jet (higher density) can be traced out along the horizontal line at y D 0. The units on the colorbar is in logŒ =g=cm3 , the axes of the plot is in units of ri . bottom: Snapshot of log jBpoloidal j at t D 1500. This plot is a lot “cleaner” than the density plot. A backbone is clearly seen along the axis out to about 200 ri from the star, further away some circular structures are seen on either side of the axis, indicative of a spiral structure spiraling in and out of the plane of view. The units are log.Bp =G/

measured a distance of 100ri from the disk. The figures show that the radial velocity is highest in the center of the jet, falling off on either side and reaching zero around 200ri from the axis. This coincides with the edge of the bow shock as seen in Fig. 1. The maximum velocity found is about 40–50 km/s. We note that this is similar to what we find from looking at the velocity component in the simulations itself, also indicating that we are doing this consistently.

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˚ at t D 1; 500. These figures are with different floor level Fig. 2 Emission line maps of OI 6,300 A for the emission, basically this is how it could look with different exposure time. The left figure has the shortest exposure. The jet itself stands out as the brightest part, being the only part seen in the left figure. In the figure to the right (long exposure) a large part of the bow shock is also visible. The intensity is in arbitrary units

˚ The velocity is in km/s, and the distance is in units Fig. 3 Radial velocity plots of OI 6,300 A. of ri

4 Summary We have here presented three-dimensional time-dependent magneto-hydrodynamic simulations of T Tauri jets launched magneto centrifugally from the surface of an accretion disk. After 1,500 rotations of the inner Kepler disk, the jet head has reached about 20 AU. The simulation box is sufficiently large that the entire bow shock is captured in the simulations, something that has shown to be important in producing the density structure inside the bow shock. Assuming a magnetic field of 10 G at the inner edge of the accretion disk, the magnetic field in the jet is of the order 1 G. The poloidal component of the magnetic field traces out a spiraling magnetic “backbone” of the jet. We have here presented a first attempt to create forbidden emission line maps of ˚ In our simulation, this emission comes from internal shocks. OI 6,300 A. The smallest scale observed in jets are 14 AU, so we are now entering a regime where it is possible to do direct comparisons between simulations and observations. Acknowledgements This work was made possible by the facilities of the Shared Hierarchical Academic Research Computing Network (SHARCNET:www.sharcnet.ca).

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References 1. Bacciotti, F., Ray, T., Mundt, R., Eisl¨offel, J., & Solf, J. 2002, ApJ, 576, 222 2. Dougados, C., Cabrit, S., & Lavalley-Fouquet, C. 2002, RMxAC, 13, 43 3. Jørgensen, M, Ouyed, R., & Christensen, M, 2001, A&A 379, 1170 4. Ouyed, R., Clarke, D. A., Pudritz, R. E. 2003, ApJ, 582, 292 5. Ouyed, R., Pudritz, R. E., 1997, ApJ, 482, 712 6. Pudritz, R. E., Ouyed, R., Fendt, Ch., & Brandenburg, A., 2007 in Protostars and Planets V, (Tucson: Univ. Arizona Press), 277 7. Stone, J. M. & Norman, M. L. 1992, ApJS, 80, 753

Magnetic Field Advection in Weakly Magnetised Viscous Resistive Accretion Disks: Numerical Simulations Gareth C. Murphy, Claudio Zanni, and Jonathan Ferreira

Abstract Observations of jets from accretion disks have constrained the launching radius to be confined to a zone close to the centre of the disk. The time evolution of a viscous, resistive accretion disk with a radially varying magnetisation distribution is studied. The magnetic field is advected into the centre of the disk over two accretion timescales. The magnetic field distribution evolves towards a power law. The field lines bend sharply near the disk surface. A super-fast magnetosonic wind is launched from the innermost region of the disk due to numerical effects.

G.C. Murphy () Dublin Institute for Advanced Studies, 31 Fitzwilliam Place, Dublin 2, Ireland e-mail: [email protected] C. Zanni INAF Osservatorio Astronomico di Torino, Via dell’Osservatorio 20, 10025 Pino Torinese, Italy e-mail: [email protected] J. Ferreira Laboratoire d’Astrophysique de Grenoble, Grenoble, France e-mail: [email protected]

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1 Motivation In order to allow material to accrete onto a central object, it is necessary to lose some angular momentum in an efficient way. This is possible in a disk system in one of two ways, either by vertical transport upwards out of the disk in a jet or outflow, or radial outward transport in a disk due to some anomalous viscosity, possibly engendered by magnetohydrodynamic turbulence. Two possible types of disk are then expected. The Jet Emitting Disk (JED) will have an equipartition magnetic field with a large braking lever arm, defined by a lengthscale equivalent to the Alfven radius. The dominant torque in the JED is magnetic. Thus the inner regions of the disk from whence the jet is observed to be emitted are expected to be JED-like. Extrapolation of slitless images of Class II jets from accretion disks around protostars have constrained the launching region to be confined to a zone of radial extent  a few AU close to the centre of the disk [1]. The outer regions are expected to behave more like the standard accretion disk (SAD). The standard accretion disk is well studied in the literature [2]. Here the braking arm is shorter and is of the order of ˛H [3]. The magnetisation is expected to be well below equipartition thus the dominant braking torque is due to the gradient of the viscous stress tensor. The SAD/JED structure has been put forward in several papers, e.g. [4]. The question then arises: how to get the large scale magnetic field into the inner region of the disk to launch a jet? Dynamo action [5, 6] may have difficulties producing a well-ordered equipartition magnetic field. Another possibility is magnetic field advection by the gas in a similar vein to the dragging of the field by gas in the interstellar medium [7]. The effective magnetic Prandtl number Pm D =v is a key parameter defining the relationship between the turbulent viscosity, v and the turbulent magnetic resistivity, . When these processes are equivalent (Pm D 1), the outward diffusion speed was found to be greater than the accretion speed, preventing the build up of a large scale field [8]. Efficient advection was found only for Pm < 0:1 [9]. However these results were originally derived for infinitely thin disks, purely poloidal potential fields and fixed boundary conditions. Considering a stratified disk advection may occur since the upper layers have much lower diffusivity [10], [11], see also Lovelace, these proceedings). In this proceedings numerical simulations over accretion timescales are performed to study this scenario. Previous numerical simulations have found indications of advection for very thick disks [12].

2 Method Numerical simulations are performed using the 2.5D resistive MHD code, PLUTO [13], using a constrained transport method to preserve r B D 0. The code is modified to include viscosity. The condition E D 0 is imposed the inner boundary in the induction equation.

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We define a reference timescale, K , as the time taken for one complete orbit ˇ at ˇ the inner disk radius, r0. We define a reference accretion timescale, acc D urr ˇ . rD6 This corresponds to 288 K . For the standard accretion disk, the initial conditions defined in [14] are used. We start with a standard accretion disk, with Prandtl number of order unity, pervaded with a magnetic field such that magnetisation, D B 2 =p, decreases radially and is at all points smaller than 102 . We let the simulation evolve in time until two characteristic accretion timescales have elapsed and inspect the magnetic flux.

3 Results 3.1 Field Advection in the Accretion Disk The entire weakly magnetised disk behaves mostly like a standard accretion disk. The accretion Mach number is always close to the analytical estimate for an isothermal disk in [15]. The viscous torque dominates over the magnetic torque for most of the disk heightscale, as expected since the magnetic field is weak. The vertical distribution of the magnetic, viscous and the components of the viscous torque are shown in Fig. 1. The magnetic flux threading the disk midplane is plotted in Fig. 2. The total flux is conserved and clearly redistributed toward the disk centre. The magnetisation parameter, D B 2 =p at z D 0, plotted logarithmically as a function

Fig. 1 Torques, magnetic, viscous and the main components of the viscous torque

Fig. 2 The magnetic flux plotted logarithmically as a function of radial extent after 0 and 2 accretion timescales have elapsed. The total flux remains constant but clearly redistributed toward the centre

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Fig. 3 The magnetisation plotted logarithmically as a function of radial extent after 0,1 and 2 acc . A net increase in the magnetisation is evident

Fig. 4 Plot of accretion disk simulation after 692 Keplerian rotations of the inner disk. The fast and Alfven surfaces are indicated. The jet has a low power and is emitted from the inner region of the disk where the magnetisation is high

of radial extent after 2 accretion timescales have elapsed, is shown in Fig. 3. The magnetisation of the inner regions of the disk is enhanced over time. The outer regions are correspondingly depleted. This may be interpreted as a redistribution in magnetic field, the field is dragged in to the inner disk region resulting in an increase in net flux threading the inner disk and in magnetisation, leading eventually to the creation of a zone at smaller radii where it is possible to reach equipartition, hence allowing the transition to a JED. The magnetisation evolves towards a power law distribution of index 2 in the inner disk.

3.2 Numerical Jet Launching Surprisingly, a superfast magnetosonic outflow is launched from the disk during the simulation. In Fig. 4 the accretion disk simulation after 692 Keplerian rotations of the inner disk is shown. The jet has a low power compared to the accretion power. The enthalpy is negligible compared to the magnetic energy. The dominant force acting is thermal pressure, as may be seen in Fig. 5. The initial lift is prompted by numerical diffusion, which allows thermal pressure from the disk to leak upwards and drive the jet up to the slow magneto sonic surface, while simultaneously numerical diffusion of density allows mass-loading onto the field lines [16].

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Fig. 5 Poloidal forces projected along the poloidal projection of a magnetic surface with anchoring radius at r D 2:4r0 plotted against altitude above the disk midplane. Fp is the pressure gradient, Fm the Lorentz force, Feff is the net gravitational+centrifugal forces and Fv is the gradient of the viscous stress tensor. The sum of all is also plotted

4 Summary For the weakly magnetised disk the magnetic field is advected inwards, contradicting the results found by [8, 9] but in agreement with the conclusions of [11]. This suggests that any large scale field present is likely to be advected inwards increasing the magnetisation to the point where jet launching becomes possible. Additionally a super fast magnetosonic outflow is launched from the inner regions which is an effect of numerical resolution. Acknowledgements This work has been supported by the ANR-05-JC42835 project funded by the “Agence National de la Recherche”. The authors wish to acknowledge the SFI/HEA Irish Centre for High-End Computing (ICHEC) for the provision of computational facilities and support. The authors acknowledge support through the Marie Curie Research Training Network JETSET (Jet Simulations, Experiments and Theory) under contract MRTN-CT-2004-005592.

References 1. P. Hartigan, S. Edwards, R. Pierson, ApJ, 609, 261 (2004) 2. N.I. Shakura, R.A. Sunyaev, A&A, 24, 337 (1973) 3. G. Pelletier, R.E. Pudritz, ApJ, 394, 117 (1992) 4. J. Ferreira, New Astronomy Review, 52, 42 (2008) 5. R.E. Pudritz, MNRAS, 195, 897 (1981) 6. A. Brandenburg, B. von Rekowski, MmSAI, 78, 374 (2007) 7. T.H. Troland, C. Heiles, ApJ, 301, 339 (1986) 8. A.A. van Ballegooijen, in Accretion Disks and Magnetic Fields in Astrophysics, ed. by G. Belvedere (1989), ASSL, vol. 156, pp. 99–106 9. S.H. Lubow, J.C.B. Papaloizou, J.E. Pringle, MNRAS, 267, 235 (1994) 10. G.S. Bisnovatyi-Kogan, R.V.E. Lovelace, ApJ, 667, L167 (2007)

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11. D.M. Rothstein, R.V.E. Lovelace, ApJ, 677, 1221 (2008) 12. I.V. Igumenshchev, R. Narayan, M.A. Abramowicz, ApJ, 592, 1042 (2003) 13. A. Mignone, G. Bodo, S. Massaglia, T. Matsakos, O. Tesileanu, C. Zanni, A. Ferrari, ApJS, 170, 228 (2007) 14. W. Kluzniak, D. Kita, astro-ph/0006266 (2000) 15. M. R´oz˙ yczka, P. Bodenheimer, K.R. Bell, ApJ, 423, 736 (1994) 16. G. Murphy, C. Zanni, J. Ferreira, Magnetic field advection in a weakly magnetised accretion disk (2008). In preparation

Extending Analytical MHD Jet Formation Models with a Finite Disk Radius Matthias Stute, Kanaris Tsinganos, Nektarios Vlahakis, Titos Matsakos, and Jos´e Gracia

Abstract The available analytical MHD models for jets, characterized by the symmetries of radial self-similarity (ADO, Analytical Disk Outflow solutions) in general have two geometrical shortcomings, a singularity at the jet axis and the nonexistence of an intrinsic scale, i.e., the jets formally extend to radial infinity. The present study focuses on imposing an outer ejecting radius of the underlying accreting disk and thus providing a finite width disk-wind. The simulations are carried out using the PLUTO code. We study the time evolution of these modified analytical models and we investigate the rich parameter space and compare the results directly with observations.

M. Stute (), K. Tsinganos, and N. Vlahakis IASA and Section of Astrophysics, Astronomy and Mechanics, Department of Physics, University of Athens, Panepistimiopolis, 157 84 Zografos, Athens, Greece e-mail: [email protected]; [email protected]; [email protected] T. Matsakos Dipartimento di Fisica Generale, Universit`a degli Studi di Torino, via Pietro Giuria 1, 10125 Torino, Italy e-mail: [email protected] J. Gracia School of Cosmic Physics, Dublin Institute of Advanced Studies, 31 Fitzwilliam Place, Dublin 2, Ireland e-mail: [email protected]

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1 Introduction Astrophysical jets and disks [11] seem to be inter-related, notably in the case of young stellar objects (YSOs), wherein jet signatures are well correlated with the infrared excess and accretion rate of the circumstellar disk [3, 9]. Disks provide the plasma to feed the jets, while jets in turn provide the disk with the needed angular momentum removal in order that accretion in the protostellar object takes place [10]. On the theoretical front, the most widely accepted description of this accretionejection phenomenon is based on the interaction of a large scale magnetic field with an accretion disk around the central object. Then, plasma is channeled and magnetocentrifugally accelerated along the open magnetic field lines threading the accretion disk, as firstly described in [2]. Several works have extended this study either by semi-analytic models using radially self-similar solutions of the full magnetohydrodynamics (MHD) equations with the disk treated as a boundary condition [17], or, by selfconsistently treating numerically the disk-jet system, e.g. [19]. The original model of [2], however, has serious limitations for a needed meaningful comparison of its predictions with observations. The outflow is not asymptotically superfast, singularities exist at the jet axis, and most importantly, an intrinsic scale in the disk is lacking with the result that the jet formally extends to radial infinity. The outflow speed at large distances may be tuned to cross the corresponding limiting characteristic [18]. The other problems have to be overcome in numerical simulations extending the analytical solutions. Gracia et al. [7] took care of the singularity at the axis, [12] showed that the solution of [18] is structurally stable. The aim of this work is to investigate numerically how imposing an outer radius of the jet, i.e. cutting off the analytical solution at arbitrary radii, affects the topology, structure and stability of a particular radially self-similar analytical solution and hence its ability to explain observations.

2 Numerical Models We solved the MHD equations with the PLUTO code1 [13] starting from initial conditions set according to a steady, radially self-similar solution as described in [18] which crosses all three critical surfaces (ADO, analytical disk outflow solution). At the symmetry axis, the analytical solution was modified as described in [7] and [12]. To study the influence of the truncation of the analytical solution, we divide our computational domain in a jet region and an external region, separated by a truncation field line ˛trunc . For lower values of the normalized magnetic flux function, i.e. ˛ < ˛trunc – or conversely smaller cylindrical radii – our initial conditions are fully determined by the solution of [18] and the modification of [7] and [12] close to the axis. In the outer region, we modify all quantities and initialize them with another

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http://plutocode.to.astro.it/

Extending Analytical MHD Jet Formation Models with a Finite Disk Radius Table 1 List of numerical science models Name Description model SC1a ˛trunc D 0:4, external analytical solution 1 D 103 , 2 D 103 model SC1b ˛trunc D 0:2, external analytical solution 1 D 103 , 2 D 103 model SC1c ˛trunc D 0:1, external analytical solution 1 D 103 , 2 D 103 model SC1d ˛trunc D 0:01, external analytical solution 1 D 103 , 2 D 103 model SC1e ˛trunc D 0:001, external analytical solution 1 D 103 , 2 D 103 model SC2 ˛trunc D 0:4, external analytical solution 1 D 100, 2 D 0:1 model SC3 same as model SC2, but solutions are swapped model SC4 ˛trunc D 0:4, external analytical solution 1 D 1, 2 D 0:1 model SC5 same as model SC4, but solutions are swapped model SC1f ˛trunc D 0:0005, external analytical solution 1 D 103 , 2 D 103 model SC1g ˛trunc D 1  105 , external analytical solution 1 D 103 , 2 D 103

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Rtrunc ŒR0  5.375 5.125 4.875 3.625 2.625 5.375 5.375 5.375 5.375 2.375 0.575

analytical solution but with modified parameters. From [18], one can show that if we start with an arbitrary MHD solution with the variables , p, v, B, then one can easily construct a second solution by using two free parameters 1 and 2 , with

0 D 1 22 , p 0 D 22 p, B0 D 2 B and v0 D 1=2 v. Thus some or all quantities 1 are scaled down in the external region depending on our choice of parameters. In Table 1, we give the parameters of the models used in this study. As expected a priori, the flows behave qualitatively very differently, depending on whether the scaled-down solution is inside or outside the original one. In the cases where the quantities are scaled down in the exterior region, the opening angle of the flow increases (see Fig. 1). The jet expands outwards due to high thermal and magnetic pressure gradients across the truncation field line. If all quantities are reduced in the internal region (models SC3 and SC5), then the opening angle decreases. In both cases, a new equilibrium is established within a few orbital periods. We investigated the integrals of motion, along the truncation field line at the interface of both regions, along an inner field line which is anchored at half of the radius of the truncation field line, and along an outer field line anchored at twice the radius. All integrals of motion converge smoothly to an asymptotic value in the inner and also in the outer region. In conclusion, in most of our models the external region also reaches a steady state [15].

3 Comparison with Observations We compare the width of eight jets measured from HST and AO observations, [6, 14, 5], to synthetic emission maps, calculated from our MHD models with a set of tools described in [8]. In order to determine the width of the jets in our models, we use a method which is as close as possible to that applied by the observers. We created convolved synthetic maps for the emission in the [SII] 6731 and [OI] 6300 lines for each numerical model and determine the jet width from the maps full-width half-maximum (FWHM) as a function of distance along the axis.

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Fig. 1 Structure of the flow (logarithmic density plots) for models SC1a – SC1e (from left to right) at timesteps t D 0 (top), t D 25 (middle) and t D 50 t0 (bottom). Also plotted is the magnetic field line anchored in the lower boundary where ˛ D ˛trunc (white line). At t D 50 t0 , we also plot the two field lines with half and twice the radius of that of the truncation field line used for investigating the intregrals of motion. A color version of this plot can be found in the electronic version of this book and in Appendix A (Fig. A:9)

We scale the PLUTO units to physical units by assuming the mass of the central object, velocity scale v0 and density scale 0 accordingly in order to match the average values in our simulations at a reference point to observed quantities in YSOs (vjet D 300 km s1 , 2 Mˇ , njet D 500 cm3 ) and calculate the remaining scaling factors from these. We run the pipeline for a grid of these parameters centered on the values given above. We vary 0 in runs 1–4 and 8 and v0 in runs 1 and 5–7. The extracted jet widths for model ADO are plotted in Fig. 2, top left. Runs 1–4 and 8 are very close together as expected. Since we extract the jet width by using the FWHM, the factor 2 in the emissivities cancels out. Much larger changes are present when we compare runs 1 and 5–7. The jet widths decrease monotonically with increasing v0 , i.e. temperature. Even the smallest jet width of all our runs, run 5, is too large by a factor of two up to seven with respect to the observations. Then we extract the jet width from emission maps for the truncated models. The resulting width are also shown in Fig. 2. As in the untruncated model ADO, the

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influence of our choice of 0 still seems to be of less importance. The models 1 and 5–7 show the same behavior as described in the previous paragraph for model ADO. The observed jet widths lie between those extracted from models SC1f and SC1g [16].

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4 Summary We have studied the effects of imposing an outer radius of the underlying accreting disk, and thus also of the outflow, on the topology, structure and time-dependence of a well studied radially self-similar analytical solution [18] and also on the observable structure of the jet. We matched an unchanged and a scaled-down analytical solution of [18] and found in all these cases steady two-component solutions. We showed that the boundary between both solutions is always shifted towards the solution with reduced quantities. Our truncated disk-wind solutions are stable for several orbital periods at the truncation radius. We created synthetic images based on our simulations of untruncated and truncated disk winds [15]. We found that the untruncated model ADO of [18] cannot account for the small jet widths found in recent optical images taken with HST and AO. With the highest degree of truncation, we can explain the observational data. The observed jet widths lie between those extracted from models SC1f and SC1g. This results can be used to infer the “real” value of the truncation radius in the observed sample of jets and interpret it as the transition radius of the jet-emitting disk to the standard accretion disk, e.g. [4]. For model SC1f, Rtrunc jzD0 D 0:26 AU, and for model SC1g, Rtrunc jzD0 D 0:033 AU. These values are consistent with other studies of this transition, e.g. [1]. Acknowledgements The present work was supported by the European Community’s Marie Curie Actions - Human Resource and Mobility within the JETSET (Jet Simulations, Experiments and Theory) network under contract MRTN-CT-2004 005592.

References 1. Anderson, J. M., Li, Z.-Y., Krasnopolsky, R., Blandford, R. D., ApJ, 590, L107 (2003) 2. Blandford, R. D., Payne, D. G., MNRAS, 199, 883 (1982) 3. Cabrit, S., Edwards, S., Strom, S. E., Strom, K. M., ApJ, 354, 687 (1990) 4. Combet, C., Ferreira, J., A & A, 479, 481 (2008) 5. Dougados, C., in Jets from Young Stars II: Clues from High Angular Resolution Observations, Lecture Notes in Physics, Vol. 742, ed. by F. Bacciotti, E. Whelan, L. Testi, (Springer, Heidelberg, 2008) 6. Dougados, C., Cabrit, S., Lavalley, C., Menard, F., A & A, 357, L61 (2000) 7. Gracia, J., Vlahakis, N., Tsinganos, K., MNRAS, 367, 201 (2006) 8. Gracia, J., et al., in preparation 9. Hartigan, P., Edwards, S., Pierson, R., ApJ, 609, 261 (2004) 10. Hartmann, L., in Protostellar Jets in Context, ed. by K. Tsinganos, T. Ray & M. Stute. (Springer, Heidelberg, 2009) 11. Livio, M., in Protostellar Jets in Context, ed. by K. Tsinganos, T. Ray & M. Stute. (Springer, Heidelberg, 2009) 12. Matsakos, T., Tsinganos, K., Vlahakis, N., Massaglia, S., Mignone, A., Trussoni, E., A & A, 477, 521 (2008) 13. Mignone, A., Bodo, G., Massaglia, S., et al., ApJS, 170, 228 (2007) 14. Ray, T. P., Dougados, C., Bacciotti, F., et al., in Protostars and Planets V, ed. by B. Reipurth, D. Jewitt, K. Keil, (University of Arizona Press, Tucson ,2007), pp. 231–244

Extending Analytical MHD Jet Formation Models with a Finite Disk Radius 15. Stute, M., Tsinganos, K., Vlahakis, N., Matsakos, T., Gracia, J., A & A, 491, 339 (2008) 16. Stute, M., Gracia, J., Tsinganos, K., Vlahakis, N., Matsakos, T., A & A, submitted 17. Vlahakis, N., Tsinganos, K., MNRAS, 298, 777 (1998) 18. Vlahakis, N., Tsinganos, K., Sauty, C., Trussoni, E., MNRAS, 318, 417 (2000) 19. Zanni, C., Ferrari, A., Rosner, R., Bodo, G., Massaglia, S., A & A, 469, 811 (2007)

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Magnetohydrodynamic Jets from Different Magnetic Field Configurations Christian Fendt

Abstract Using axisymmetric MHD simulations we investigate how the overall jet formation is affected by a variation in the disk magnetic flux profile and/or the existence of a central stellar magnetosphere. Our simulations evolve from an initial, hydrostatic equilibrium state in a force-free magnetic field configuration. We find a unique relation between the collimation degree and the disk wind magnetization power law exponent. The collimation degree decreases for steeper disk magnetic field profiles. Highly collimated outflows resulting from a flat profile tend to be unsteady. We further consider a magnetic field superposed of a stellar dipole and a disk field in parallel or anti-parallel alignment. Both stellar and disk wind may evolve in a pair of outflows, however, a reasonably strong disk wind component is essential for jet collimation. Strong flares may lead to a sudden change in mass flux by a factor two. We hypothesize that such flares may eventually trigger jet knots.

C. Fendt () Max Planck Institute for Astronomy, K¨onigstuhl 17, D-69117 Heidelberg, Germany e-mail: [email protected]

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1 Jets as Collimated MHD Flows Astrophysical jets are launched by magnetohydrodynamic (MHD) processes in the close vicinity of the central object – an accretion disk surrounding a protostar or a compact object [1, 2, 9, 20, 21, 24]. Numerical simulations of MHD jet formation are essential for our understanding of the physical processes involved. In general, simulations may be distinguished in those taking into account the evolution of the disk structure and others considering the disk surface as a fixed-in-time boundary condition for the jet. The first approach allows to directly investigate the mechanism lifting matter from the disk into the outflow [3, 10, 11, 14, 17, 18, 22, 24]. This approach is computationally expensive and still somewhat limited by spatial and time resolution. In order to study the acceleration and collimation of a disk/stellar wind it is essential to follow the dynamical evolution for (a) very long time (b) on a sufficiently large grid with (c) appropriate resolution. For such a goal, the second approach is better suited [4, 5, 6, 7, 12, 13, 15, 19, 25], allowing as well for parameter studies. The case of superposed stellar/disk magnetic field is rarely treated in simulations, still, the first model was discussed already in [23]. Simulations of a dipole with aligned vertical disk field are presented by [16,18]. The stellar field has important impact on the jet formation process as enhancing the magnetic flux, adding a central pressure, and providing excess angular momentum for the launching region.

2 Model Setup We use the ZEUS-3D MHD code extended for physical magnetic diffusivity (see description in [6]). The set of MHD equations considered is the following,   4 @B 4 @

Cr . v/ D 0; r B D 0; j D r B; r  v  B  j D 0 (1) @t c @t c @u jB C .v r/ v C r.p C pA / C r˚  D 0;

@t c 

(2)

with the usual notation [4,5,6,7,19]. We do not solve the energy equation, but apply an internal energy e D p=.  1/ of a polytropic gas ( D 5=3). Turbulent Alfv´enic pressure pA allows for a “cool” corona. The turbulent magnetic diffusivity .r; zI t / can be related to pA applying our toy model [6]. The  ' 0:01 was chosen low and does not affect collimation. Diffusivity is, however, essential for reconnection processes. We distinguish setup DW (pure disk wind) and SDW (stellar wind plus disk wind) by choice of boundary and initial conditions. Model DW investigates different disk magnetic field and mass flux profiles [4]. Model SDW investigates the interrelation of the stellar magnetosphere with the surrounding disk jet [5].

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2.1 Boundary Conditions In setup DW we distinguish along the equatorial plane the gap region r < 1:0 and disk region r > 1:0. The magnetic field is fixed in time and is determined by the initial condition. We have chosen a power law, Bp .r; 0/  r  , and investigate different . In setup SDW we further distinguish the star from r D 0  0:5, and the gap from r D 0:5  1:0. Co-rotation radius and inner disk radius coincide. A Keplerian disk is the boundary condition for the mass inflow from the disk surface into the corona. Matter is “injected” from the disk (and the star) with low velocity vinj .r; 0/ D i vK .r/BP =BP and density inj .r; 0/ D i .r; 0/. Typically, i ' 103 and i ' 100 for stellar and disk wind, but could be chosen differently.

2.2 Initial Conditions As initial state we prescribe a force-free magnetic field and a hydrostatic equilibrium

.r; z; t D 0/ D .r 2 C z2 /3=4 . For model DW we calculate the initial field distribution from the disk magnetic field profile using our finite element code (see [4, 8]). For model SDW the initial field is a superposed dipole plus disk field. For the disk component we apply the potential field of R [6, 19]. We prescribe the initial field by the magnetic flux distribution  .r; z/  Bp d A,

 .r; z/ D 0;d

 1 p 2 r2 r C .zd C z/2  .zd C z/ C 0;? : (3) r .r 2 C .zd C z/2 /3=2

Certain field combinations are investigated, parameterized by the disk 0;d and stellar magnetic flux 0;? (Fig. 1).

Fig. 1 Example initial magnetic field distributions (poloidal magnetic field lines). Full and dashed lines indicate the direction of magnetic flux. Magnetic field parameters: 0;d D 0:01; 0:01; 0:1; resp. 0;? D 5:0; 5:0; 3:0 (from left to right). From [5]

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3 Disk Wind Magnetization and Jet Collimation Simulations of setup DW were run for different disk magnetic field profiles Bp .r; z D 0/  r  and density profiles inj .r; z D 0/  r  . In general, we find an increasing degree of collimation with decreasing slope of the disk magnetic field profile (see [4]). This seems to rule out launching models for collimated jets from a concentrated magnetic flux such as e.g. the X-wind scenario. A steep density profile leads to a higher collimation degree, which is not surprising as the mass flux is more concentrated just by definition of the boundary condition. A physically meaningful classification taking into account both density and magnetic field can be achieved by comparing the degree of collimation degree versus 1 1 vinj .r/˝K .r/2 , thus, disk wind magnetization profile,  .r; z D 0/  Bp2 .r/r 4 inj .2 C1=2/   .r; z D 0/  r  r . The resulting diagram Fig. 2 shows a convincing correlation between the magnetization power law index  and the average degree of collimation <  >. The width of the .   <  >/-correlation is due to further differences in the parameter space.

4 Jet Mass Flux Triggered by Star-disk Magnetospheric Flares Simulations of setup SDW were run for aligned and anti-aligned orientation of dipole versus disk field and for different strength of both field contributions [5]. Independent of the alignment, the central dipole does not survive on the large scale. A two-component outflow emerges as stellar wind plus disk wind. For a reasonably strong disk magnetic flux a collimated jet emerges. If the overall outflow

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Fig. 3 Poloidal magnetic field evolution during one example flare around t D 180. Solid and dashed lines indicate the direction of total magnetic flux of the superposed dipolar and disk magnetic field components. Shown are time steps: 1760, 1790, 1810 (from left to right)

Fig. 4 Axial mass flux integrated along the upper z-boundary versus time. Note the change of mass flux of 10–50% during the flare events. High mass fluxes for t < 500 indicate sweeping off of the initial corona

is dominated by a strong stellar outflow a low mass flux disk wind remains uncollimated. The best setup to launch a collimated jet from a star-disk magnetosphere is that of a relatively heavy disk wind and high disk magnetic flux. Stellar wind dominated simulations may give a high degree of collimation, however they collimate to too small radii. Stellar magnetic flux dominated outflows tend to stay un-collimated. In some simulations we observe reconnection flares, similar to coronal mass ejections, typically expanding and reconnecting within 70 orbital periods of the inner disk. This is similar to [10], however, in their case reconnection is are triggered by time-variation of the accretion rate. In our case the reconnection/flares seem to be triggered by the evolution of the outer disk wind. Even for our very long time-scales the outer disk outflow is still dynamically evolving, thus changing the cross-jet force equilibrium and forcing the inner structure to adjust accordingly. The flare events are accompanied by a temporal change in outflow mass flux and momentum. Figure 4 shows the mass loss rate in axial direction integrated across the jet. We see two flares with a 10%-increase in the mass flux followed by a sudden decrease of mass flux by a factor of two. This behavior is also mirrored in the poloidal velocity profile. Considering the ejection of large-scale flares and the follow-up re-configuration of outflow dynamics, we hypothesize that the origin of jet knots is triggered by such

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flaring events. Our time-scale for flare generation is of 1,000 rotational periods and longer than the typical dynamical time at the jet base, but similar to the observed knots. The flare itself for about 30–40 inner disk rotation times.

References 1. Blandford, R., Payne, D., 1982, MNRAS 199, 883 2. Camenzind, M.: Magnetized disk-winds and the origin of bipolar outflows. In: Klare, G. (ed) Rev. Mod. Astron. 3, p.234, Springer, Heidelberg (1990) 3. Casse, F., Keppens, R. 2002, ApJ, 581, 988 4. Fendt, C. 2006, ApJ, 651, 272 5. Fendt, C. 2008, ApJ, in press, arXiv:0810.4154v1 [astro-ph] 6. Fendt, C., Cemeljic, M. 2002, A&A, 395, 1045 7. Fendt, C., Elstner, D. 2000, A&A, 363, 208 8. Fendt, C., Camenzind, M., Appl, S. 1995, A&A, 300, 791 9. Ferreira, J., Dougados, C., Cabrit, S. 2006, A&A, 453, 785 10. Goodson, A., Winglee, R., B¨ohm, K. 1997, ApJ, 489, 199 11. Hayashi, M., Shibata, K., Matsumoto, R. 1996, ApJ, 468, L37 12. Kigure, H., Shibata, K. 2005, ApJ, 634, 879 13. Krasnopolsky, R., Li, Z.-Y., Blandford, R. 1999, ApJ, 526, 631 14. Kudoh, T., Matsumoto, R., Shibata, K. 1998, ApJ, 508, 186 15. Matsakos, T., Tsinganos, K., Vlahakis, N., Massaglia, S., Mignone, A., Trussoni, E. 2008, A&A, 477, 521 16. Matt, S., Goodson, A., Winglee, R., B¨ohm, K.-H. 2002, ApJ, 574, 232 17. Meliani, Z., Casse, F., Sauty, C. 2007, A&A, 460, 1 18. Miller, K., Stone, J. 1997, ApJ, 489, 890 19. Ouyed, R., Pudritz, R., 1997, ApJ, 482, 712 20. Pudritz, R., Norman, C. 1983, ApJ, 274, 677 21. Pudritz, R., Ouyed, R., Fendt, C., Brandenburg, A.: Disk Winds, Jets, and Outflows: Theoretical and Computational Foundations. In: Reipurth, B., Jewitt, D., Keil, K. (ed) Protostars & Planets V, p.277, University of Arizona Press, Tucson (2007) 22. Romanova, M., Ustyugova, G., Koldoba, A., Lovelace, R. 2002, ApJ, 578, 420 23. Uchida, Y., Low, B. 1981, Journal of Astroph. and Astron., 2, 405 24. Uchida, Y., Shibata, K. 1984, PASJ, 36, 105 25. Ustyugova, G., Koldoba, A., Romanova, M., Chechetkin, V., Lovelace, R. 1995, ApJ, 439, L39

Resistive MHD Jet Simulations with Large Resistivity ˇ Miljenko Cemelji´ c, Jos´e Gracia, Nektarios Vlahakis, and Kanaris Tsinganos

Abstract Axisymmetric resistive MHD simulations for radially self-similar initial conditions are performed, using the NIRVANA code. The magnetic diffusivity could occur in outflows above an accretion disk, being transferred from the underlying disk into the disk corona by MHD turbulence (anomalous turbulent diffusivity), or as a result of ambipolar diffusion in partially ionized flows. We introduce, in addition to the classical magnetic Reynolds number Rm, which measures the importance of resistive effects in the induction equation, a new number Rb, which

ˇ M. Cemelji´ c () TIARA, Academia Sinica, National Tsing Hua University, No. 101, Sec. 2, Kuang Fu Rd., Hsinchu 30013, Taiwan e-mail: [email protected] J. Gracia School of Cosmic Physics, Dublin Institute for Advanced Studies, 31 Fitzwilliam Place, Dublin 2, Ireland e-mail: [email protected] N. Vlahakis and K. Tsinganos IASA and Section of Astrophysics, Astronomy and Mechanics, Dpt. of Physics, Univ. of Athens, Panepistemiopolis 15784 Zografos, Athens, Greece e-mail: [email protected]; [email protected]

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measures the importance of the resistive effects in the energy equation. We find two distinct regimes of solutions in our simulations. One is the low-resistivity regime, in which results do not differ much from ideal-MHD solutions. In the high-resistivity regime, results seem to show some periodicity in time-evolution, and depart significantly from the ideal-MHD case. Whether this departure is caused by numerical or physical reasons is of considerable interest for numerical simulations and theory of astrophysical outflows and is currently investigated.

1 Introduction In Vlahakis and Tsinganos [7] general classes of self-consistent ideal-MHD solutions have been constructed. In Vlahakis et al. [8] Blandford and Payne [1] model was analysed, and the problem with the terminal wind solution (which was

Fig. 1 The initial setup, which is slightly modified analytical solution, is shown in the Left panel. The solid lines represent logarithmically spaced isocontours of density. It is also shown in colour grading, in red to violet colour, for the logarithm of density 1 to 4; respectively. In the Right panel shown is, in the same grading, the solution with large magnetic diffusivity. It does not reach stationary state, and shows some periodicity in time evolution. The dashed lines depict poloidal magnetic field lines, and the dotted lines depict the fast magnetosonic, Alfven and slow magnetosonic critical surfaces, top to bottom, respectively. A color version of this plot can be found in the electronic version of this book and in Appendix A (Fig. A.11)

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not causally disconnected from the disk) has been solved. The common deficiency of all radially self-similar models, a cut-off of the solution at small cylindrical radii and also at some finite height above the disk because of a strong Lorentz force close to the system’s axis has been corrected numerically. A search in the numerical simulations for solutions at larger distances from the disk has been performed in Gracia et al. [5] with NIRVANA code (version 2.0, [9]), and similar results were obtained also using the PLUTO code in Matsakos et al. [6]. Extension in the resistive-MHD ˇ has been investigated in Cemelji´ c et al. [3] using the NIRVANA code and some of results we present here. Results of these investigations could also have implications in the numerical simulations of magnetospheric interaction in vicinity of the young stellar objects, where the resistivity plays important role. Our numerical simulations are initiated by the slightly modified analytical solutions for radially self-similar flow from Vlahakis et al. [8], and then evolved in the resistive MHD simulations by NIRVANA code. Our initial setup is shown in the left panel of Fig. 1.

2 New Characteristic Number In addition to the magnetic Reynolds number Rm D VR/, which describes influence of the magnetic diffusivity  in the induction equation, we introduced a new number, which describes the influence in the energy transport equation-see Fig. 2. It can be written in terms of Rm and plasma beta as Rb D Rmˇ/2. It is the

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ratio of the pressure term over the Joule heating term in the energy equation. When Rb is smaller or close to unity, which can happen even when Rm is much larger than unity, the energy dissipation becomes important. It might define one additional mode of resistive-MHD solutions, indicated in our search for eventual onset of super-critical resistive regime.

3 Results The resistive MHD jets are similar to ideal-MHD solutions for a finite range of magnetic diffusivity, in which they reach a well defined stationary state. This state only slightly differs from the initial state, as expected, since the initial setup was slightly modified analytical stationary solution. Departure from the ideal-MHD

Fig. 3 Reconnection and re-shaping of the magnetic field in the vicinity of the young stellar object. Initially pure stellar dipole field reshapes into the stellar and disk open field during the time-evolution. The time is measured in the number of rotations at the inner disk radius, which is at Ri D 3:0 in these simulations. Without the substantial resistivity, reconnection does not occur and simulations terminate because of numerical reasons

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regime occurs for larger values of magnetic diffusivity, above some critical value. One such result is shown in the right panel of Fig. 1. We note possible existence of the distinct super-critical regime in magnetic diffusivity in our simulation for the outflows initialised with a self-similar analytical solutions. We also define the new characteristic number, which describes the influence of the resistivity on the energy transport equation. Physical parameters and the eventual periodicity of the supercritical resistive solutions are currently under investigation. Such solutions might be interesting for investigations of accretion flows in the vicinity of young stellar objects, where the magnetic resistivity seems to play important role. In the Fig. 3 we show one case in numerical simulations of magnetospheric interaction in the closest vicinity of young stellar object [2]. These simulations have been performed with code ZEUS347, which is our resistive version of Zeus3D code [4]. Magnetic reconnection shows to play essential role in re-shaping the initial stellar dipole, which enables the launching of outflows. In our numerical simulations, if the resistivity is too small, reconnection will not occur. Therefore, we need to investigate parameter space for resistivity, and we need to understand what are the effects of large, and not only negligible or very small resistivity. Acknowledgements This work was supported in part by EC’s Marie Curie Actions-Human Resource and Mobility within the JETSET network under contract MRTN-CT-2004005592. MCˇ expresses gratitude to TIARA/ASIAA in Taiwan for possibility to use their Linux clusters and JETSET for supporting this collaboration.

References 1. Blandford R. D., Payne D. G., 1982, MNRAS, 199, 883 ˇ 2. Cemelji´ c M., Hsien S., Chiang T.-Y., 2009, in preparation ˇ 3. Cemelji´ c M., Gracia J., Vlahakis N., Tsinganos K., 2008, MNRAS, 389, 1022, 1032 ˇ 4. Fendt Ch., Cemelji´ c M., 2002, A&A, 395, 1045 5. Gracia J., Vlahakis N., Tsinganos K., 2006, MNRAS, 367, 201 6. Matsakos T., Tsinganos K., Vlahakis N., Massaglia S., Trussoni E., 2008, A&A, 477, 521 7. Vlahakis N., Tsinganos K., 1998, MNRAS, 298, 777 8. Vlahakis N., Tsinganos K., Sauty C., Trussoni E., 2000, MNRAS, 318, 417 9. Ziegler U., 1998, Comput. Phys. Commun., 109, 111

The X-wind Model Mike J. Cai

Abstract Based on the principle of magneto-centrifugal acceleration, the interaction of an electrically conducting accretion disk with the magnetosphere of a young star will naturally give rise to X-winds and funnel flows. A distinguishing feature of the X-wind model is the trapped flux in a small region at the inner edge of the disk, where both the X-wind and the funnel flow originate. We will review some of the theoretical developments of the X-wind model and how some of its predictions can be tested by observation.

1 Introduction Jets and outflows are an integral part of the star formation process [22, 34]. Blandford & Payne [2] proposed the idea of magnetohydrodynamic (MHD) winds from accretion disks as a viable mechanism for driving jets. They demonstrated that if a magnetic field is anchored in the accretion disk rotating about a gravitating source and if the field lines bend outward by more than 30ı , gas in the disk can be accelerated centrifugally along flux tubes. Since then, two classes of theories have been developed based on this concept of magneto-centrifugal acceleration, known as the disk wind [18] and the X-wind [38]. They differ in attributing the origin of the magnetic field and boundary conditions in the disk. The disk wind follows more faithfully to the original model of Blandford and Payne [2]. It assumes the magnetic field is either advected in by the accretion flow or generated by the disk dynamo. In principle, depending on the distribution of the magnetic field lines in the disk, jets can be launched from a large range of radii. Extensive work has been done both analytically by utilizing self-similarity (e.g., [11]) and numerically by finite differencing (e.g., [28]).

M.J. Cai () Academia Sinica, Institute of Astronomy and Astrophysics, Taiwan e-mail: [email protected]

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The X-wind assumes that the magnetic field originates from the protostar, and the outflow is driven from a small region near the inner edge of the disk. In addition to outflows, the X-wind theory also includes a natural mechanism for magnetically channeled accretion flows which truncate the disk near the corotation radius. In the remainder of this article, we shall outline our current understanding of the Xwind theory, and review some observational studies which either support or help to extend the model. We will also offer a brief discussion on the implications of disk magnetization on MHD winds.

2 General Properties of the X-wind The X-wind model considers the interaction between the magnetosphere of a young stellar object and an unmagnetized accretion disk. For simplicity, we assume the unperturbed stellar magnetic field to be a dipole aligned with the spin axis of the disk. The inward drift of the accreting material also pushes in the stellar magnetic field C right above the disk surline in the midplane, p developing a radial component B$ C face. Once B$ > Bz = 3, where Bz is the vertical field component, the gas frozen to the flux tube becomes unstable to magneto-centrifugal acceleration [2]. The outgoing wind carries away angular momentum at the expense of the disk, which drives the footpoints of the magnetic field further inward. This process would continue unabated until there is a source of angular momentum that stops the inward migration of the field lines. In the original model of Shu et al [33], the steady state solution calls for crushing the magnetosphere all the way onto the stellar surface by an accretion disk. The star would then be rotating at breakup near the equator, where outflow occurs. To accommodate the slow rotators, such as the classical T Tauri stars (CTTSs) that typically only rotate at one tenth of the breakup [3, 12], the X-wind model was generalized to include a truncated disk and the subsequent accretion flow was funneled by the magnetic field connecting the star to the disk [35]. As matter moves from the disk to the star at nearly constant angular velocity, the excess angular momentum is deposited in the magnetic field in the form of Maxwell torque, which is then transported back to the disk. The footpoints of the funnel-flow field lines gain angular momentum and try to move outward. The outward press of the funnel-flow field lines and the inward press of the wind field lines pinch the magnetic field into an annulus centered at radius RX , which we call the X-region. For a protostar with mass M , magnetic dipole momentum  , this radius can be estimated to be

RX D

4=7 ˚dx

 2 GM MP D

!1=7 ;

(1)

where MP D is the mass accretion rate from the disk, and ˚dx is an order unity parameter that measures the amount of trapped magnetic flux [26]. For parameters appropriate for CTTSs, RX has a typical value of 5–10 times the stellar radius.

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The migration of the field lines stops only when the excess angular momentum from the funnel flow can diffuse across the field lines deep inside the disk, where non-ideal MHD effects become important, and be removed by the wind. In the absence of efficient mechanisms for angular momentum transport, the radial extent of the X-region can be written as $ D RX , where  is the ratio of the sound speed to the Keplerian speed at RX [35]. If we take the thermal sound speed to be a  5kms1 , and the Keplerian speed at RX to be vK  100kms1 [25], then   0:05. In the asymptotic analysis of a cold MHD wind when we take  ! 0, the trapped magnetic field lines all appear to originate from a single point in the meridional plane. In steady state, the trapped magnetic flux plays a dynamically important role in driving the funnel flows and outflows, as well as locking the spin of the star to the Keplerian frequency at RX ,  ˝ D ˝X D

GM 3 RX

1=2 :

(2)

Conservation of mass and angular momentum allows us to estimate the mass loss from the wind MP w D fw MP D , and the mass accretion rate in the funnel flow MP  D .1  fw /MP D , where 1  JN  : (3) fw D JNw  JN Here JNw and JN are the average specific angular momentum of the wind and the 2 ˝X , and is the negative of the funnel flow, respectively, measured in units of RX 2 ˝X . viscous torque acting on the inner edge of the disk measured in units of MP D RX Substituting realistic parameters for CTTSs, both JN and are negligible compared to unity, and we may approximate the mass loss fraction in the wind as fw  1=JNw .

3 Mathematical Formulation and Global Solutions Under the assumption of stationarity, axisymmetry, and field freezing, Shu et al. [33] and Shu et al. [36] wrote down the non-dimensionalized governing equations of an MHD flow in a frame that is rotating with the star [15, 31]. The continuity equation and the induction equation can be satisfied if we introduce a stream function

u$ D

1 @ ; $ @z

uz D 

1 @ ; $ @$

(4)

and write the magnetic field as (see, e.g., [23]) B D ˇ u:

(5)

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The assumed symmetries allow us to identify three conserved quantities along each streamline, corresponding to inverse mass loading ˇ. /; the specific angular momentum J. /, (6) J  $ 2 C $ .1  ˇ 2 /u' ; and the Bernoulli’s integral H 

1 2 juj C  2 ln C Veff ; 2

(7)

where the effective potential in its non-dimensional form is written as Veff D  p

1

1 3  $2 C : 2 $ 2 C z2 2

(8)

Formally, these conserved quantities are determined by requiring that the solution passes smoothly across three critical surfaces, corresponding to when the flow speed equals to the slow, Alfv´en, and fast speeds of MHD waves. In the asymptotic limit  ! 0, one can show that H ! 0 and the slow surface shrinks to a point. To the leading order, the momentum equation may be written in terms of the other conserve quantities as the Bernoulli’s equation jr j2 C

1 A2



2 J 2$ 2 Veff  1 C 2 D 0; 2 $ .ˇ  $ 2 A2 /2

(9)

which describes energy conservation along each streamline, and the Grad-Shafranov equation, 1 r .Ar / C A



 0 J J 2ˇˇ 0 Veff  1 C D 0; $2 $2 .ˇ 2  $ 2 A/2

(10)

which describes force balance across streamlines. Here A  .ˇ 2  1/=.$ 2 / is the Alfv´en discriminant, and it takes on positive (negative) values when the flow is sub-Alfv´enic (super-Alfv´enic). Combined with the Bernoulli’s equation (9), the Grad-Shafranov equation (10) is a partial differential equation of mixed type, which changes from elliptic to hyperbolic across an internal surface where the poloidal speed of the gas equals the Alfv´en speed [16, 29]. To complicate matters even more, the characteristics on the critical surface are singular, which renders a direct numerical integration of the governing equations formidable. To alleviate these problems, Cai et al. [6] adopted a different approach based on a variational principle. They noticed that the governing equations may be derived from an action Z " S  2

1 1 Ajr j2  2 2A



# 2 J Veff 1 C 2 $d$d z: $2 ˇ  $ 2A

(11)

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Fig. 1 A typical solution with ˇN D 3, adopted from [6]. The dotted curves represent the streamlines, and the solid curves are the isodensity contours, separated by logarithmic intervals. The dashed and dash-dotted curves mark the location of Alfv´en and fast surfaces, respectively

By interpolating between the X-point solution of Shu et al. [36] and the asymptotic solution of Shu et al. [37], the global solution may be represented by a set of trial functions with undetermined parameters. These parameters are then varied until an extremum of the action (11) is reached. In the absence of a detailed calculation describing how matter is loaded onto the field lines subsonically, the authors adopted N an ad hoc inverse loading function ˇ. / D .2ˇ=3/.1  /1=3 , for an order unity N N ˇ. Figure 1 shows a typical solution with ˇ D 3. It is noticed here that the density contours become very quickly cylindrically stratified even though the streamlines follow more or less radial trajectories at large distances. Shang et al. [32] argued that because of this density stratification, the jet-like appearances are merely optical illusions.

4 Observational Tests Other than the inverse loading function ˇ. /, the X-wind model has very few free parameters, which allows it to make specific and testable predictions [38]. By comparing with observations, we either gain confidence in the validity of the model or are motivated to extend the theory. Here we review some of the recent developments.

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4.1 Multipolar Funnel Flow The work by Valenti & Johns-Krull [42] showed that for most CTTSs, the overall level of polarization is low on the stellar surface. Furthermore, Johns-Krull & Gafford [17] indicated that the hot spots caused by accretion funnel typically only cover less than 1% of the total stellar surface. The two pieces of evidence suggest that the field geometry on the star is most likely more complicated than a simple dipole, as calculated by Ostriker & Shu [26]. Fortunately, this is not a crucial assumption of the X-wind model. The central pillar of the theory is the trapped flux in the X-region; changing the boundary condition on the star only modifies the detailed geometry of the accretion flow and leaves the X-wind virtually intact. Mohanty & Shu [24] redid the calculation of the funnel flow assuming a multipolar field geometry on the star. The readers may refer to the article by Mohanty in this volume for a more detailed discussion. It suffices to mention that there exists a large class of solutions within the frame work of X-wind that can reproduce the polarization and hot spot coverage observed.

4.2 Disk Locking Using spectroscopy of the CO fundamental emission from the near-infrared rovibrational transitions, Carr [7] computed the inner gas disk radius for a collection of CTTSs. If we exclude the two transitional disks (out of a total of 13 in the sample), all derived values of the inner disk radii Rdisk fall within a factor two of the corotation radii Rcorot , with a mean value of Rdisk D 0:7Rcorot . It was pointed out that the fundamental emission of CO is sensitive to very low column densities, which are typically four or more orders of magnitude lower than the disk. Thus part of the emission may come from the accretion funnel between the disk and the star, and it is perhaps better to interpret the value quoted above as a lower limit for the inner disk radius.

4.3 Jet Rotation and Launch Radius Bacciotti et al. [1] and Coffey et al. [9] reported velocity gradients across the jet axis in four T Tauri systems. If interpreted as rotation, the amount of angular momentum corresponds to jets launched from disks at radii 0:5  2AU from the star, which are too large to be compatible with X-winds. However, from radio measurements, Cabrit et al. [5] showed that in one of the systems, the disk is rotating in the opposite direction as the inferred jet rotation, which questions the validity of interpreting the velocity asymmetries as true rotation. If one were to measure jet rotation, the best candidates would be the edge-on systems, so that the confusion from intrinsic poloidal velocity asymmetries in the jet can be minimized. At millimeter

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wavelengths, Pety et al. [27] found no evidence of outflow rotation in HH 30, whose jet axis lies within 8ı to the plane of the sky [4]. This result was later confirmed by Coffey et al. [10] with optical and ultra-violet observations. Recently, Lee et al [19, 20, 21] mapped two other edge-on systems HH 211 and HH 212 in submillimeter wavelengths. Of all the SiO knots, only a small fraction exhibit velocity asymmetries, which corresponds to an inferred terminal specific angular momentum of jt D 45 AU km s1 for HH 211 and jt D 25 AU km s1 for HH 212. The authors argued that since these knots were observed in a bow shock cavity, the specific angular momentum is unlikely to be diluted by the entrained material. Variations in the shock speed may also contribute to the observed velocity asymmetry, and hence the numbers quoted above are to be interpreted as upper limits. Independent of specific models, MHD driven winds would reach a terminal velocity of vt D .2J  3/1=2 $b ˝b , and a specific angular momentum jt D J $b2 ˝b along any given streamline launched from a radius $b rotating at an angular velocity ˝b . Combining the two equations allows us to determine the launch radius jt $b D vt

p 2J  3 : J

(12)

Assuming a terminal velocity vt > 200 km s1 , and an average value of J D 4 appropriate for the X-winds, we deduce $b < 0:12 AU for HH 211 and $b < 0:07 AU for HH 212. Any larger choice of J will make the launch radius even smaller. In other words, even if these jets were disk winds (with typical values of J > 15), they must be driven from the disk at radii very close to the inner edge.

5 Magnetized Disks and Implication for Disk Winds and X-winds In the foregoing discussion, the accretion disk has been assumed to be unmagnetized. Since the ambipolar diffusion time scale is much longer than the dynamic collapse time scale, one may invoke approximate field freezing condition in which a fraction of the magnetic flux contained in the original molecular cloud core is advected inward during the gravitational collapse process. If non-ideal MHD conditions such as electric resistivity are taken into consideration, Shu et al [40] estimated a value of mass-to-flux ratio 2G 1=2 M =˚  4 by the time a protostar and an accretion disk are formed. Because the disk is much less massive but has much more cross section area than the star, it serves as a reservoir for the magnetic flux and angular momentum, while the star contains most of the mass. A direct consequence of disk magnetization is sub-Keplerian rotation [40]. With a net flux threading the disk, the vertically averaged radial force equation reads ˝2 D

C GM Bz B$ ;  $3 2$˙

(13)

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where ˙ is the surface density. The magnetic tension provides an extra support against gravity so that the gas at a given radius does not need to rotate at its full Keplerian rate. In order for MHD winds to carry enough mass flux, the gas in the disk would require an additional boost to climb out of the depth of the effective gravitational potential before magneto-centrifugal acceleration can operate. If we use fN to denote the vertically averaged fraction of Keplerian rotation, ˝ D fN.GM =$ 3 /1=2 , then thermal launching generally requires 1 fN2  O.A2 /, where A is the aspect ratio of the disk [41]. On the other hand, the condition of hydrostatic equilibrium in the z direction takes the form 1 fN2  O.A/, which creates an effective potential too deep to be compensated by thermal effects alone. Without coming to the edge, the inner disk may have an aspect ratio A . 1%, and even fN D 0:99 will hinder a thermally driven disk wind. Wardle & K¨onigl [43], Ferreira & Pelletier [14], Ferreira [13], Casse & Ferreira [8], and Salmeron et al. [30] invoked magnetic diffusivity to overcome the potential barrier. They were successful in producing fast jets with reasonable mass loading by assuming an order unity Prandtl’s number. Unfortunately, this choice yields supersonic accretion flows in the midplane, which is unfeasible for ordinary protostellar disks. If the Prandtl’s number is taken to be O.A1 /, as advocated by Shu et al. [40] and Shu et al. [41], then the accretion flow remains subsonic, but the jets are too lightly loaded to be compatible to observation. In general, if the Prandtl’s number remains large throughout the disk, there can be no significant mass loss until one reaches the inner edge (see below). For numerical evidences of this phenomenon, refer to the reviews by Romanova and Zanni in this volume. The X-wind escapes the fate of sub-Keplerian rotation by having a magnetic fan structure (see Fig. 3 of [39]). As one moves across the X-region toward the inner edge of the disk, the radial component of the field changes from positive to negative, which makes the rotation change from sub-Keplerian to super-Keplerian according to the force equation (13). Since the density changes on an  scale near the truncation radius, the sharp pressure gradient (ignored in previous analysis for the disk proper) will also help to maintain this configuration and to load field lines. The net effect of the disk field is to provide an external pressure for the expanding X-wind. Because the X-wind field lines are being pushed toward the axis, a smaller fraction of the magnetic field lines would bend outward by more than 30ı , and fw becomes smaller. The resulting jet is faster and better collimated.

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Disk-Magnetosphere Interaction and Outflows: Conical Winds and Axial Jets Marina M. Romanova, Galina V. Ustyugova, Alexander V. Koldoba, and Richard V.E. Lovelace

Abstract We investigate outflows from the disk-magnetosphere boundary of rotating magnetized stars in cases where the magnetic field of a star is bunched into an X-type configuration using axisymmetric and full 3D MHD simulations. Such configuration appears if viscosity in the disk is larger than diffusivity, or if the accretion rate in the disk is enhanced. Conical outflows flow from the inner edge of the disk to a narrow shell with an opening angle 30–45ı . Outflows carry 0.1–0.3 of the disk mass and part of the disk’s angular momentum outward. Conical outflows

M.M. Romanova () and R.V.E. Lovelace Department of Astronomy, Cornell University, Ithaca, NY 14853, USA e-mail: [email protected]; [email protected] G.V. Ustyugova Keldysh Institute of the Applied Mathematics RAS, Moscow 125047, Russia e-mail: [email protected] A.V. Koldoba Institute for Mathematical Modeling RAS, Moscow 125047, Russia e-mail: [email protected]

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appear around stars of different periods, however in case of stars in the “propeller” regime, an additional – much faster component appears: an axial jet, where matter is accelerated up to very high velocities at small distances from the star by magnetic pressure force above the surface of the star. Exploratory 3D simulations show that conical outflows are symmetric about rotational axis of the disk even if magnetic dipole is significantly misaligned. Conical outflows and axial jets may appear in different types of young stars including Class I young stars, classical T Tauri stars, and EXors.

1 Introduction Jets and winds are observed in young stars at different stages of their evolution from very young stars up to classical T Tauri stars (CTTSs) where smaller-scale jets and winds are observed (see review by [23]). A significant number of CTTS show signs of outflows in spectral lines, in particular in He I [6,14]. High-resolution observations show that outflows often have an “onion-skin” structure, with highvelocity outflows in the axial region, and lower-velocity outflow at larger distance from the axis [2]. High angular resolution spectra of [FeII] 1.644 m emission line taken along the jets from DG Tau, HL Tau and RW Auriga revealed two components: a high-velocity well-collimated extended component with v  200–400 km/s and a low-velocity  100 km/s uncollimated component which is close to a star [21, 22]. High-resolution observations of molecular hydrogen in HL Tau have shown that at small distances from the star the flow shows a conical structure with outflow velocity 50–80 km/s [31]. Different models have been proposed to explain outflows from CTTSs (see review by [8]), including models where the outflow originates from the inner regions of the accretion disk (e.g., [16, 12, 8]), and the X-wind type models [28, 29, 20, 4] where most of the matter flows from the disk-magnetosphere boundary. In this work we consider only the second type of models. We developed conditions favorable for X-type outflows and performed axisymmetric and exploratory 3D MHD simulations for both slowly and rapidly rotating stars including stars in the propeller regime.

2 Conical Winds 2.1 Axisymmetric (2.5D) Simulations To investigate outflows from the disk-magnetosphere boundary it was important that the magnetic field lines be bunched into an X-type configuration. Such bunching will occur if magnetic field lines threading the disk move inward to the star faster than they diffuse outward. This happens for example when the viscosity in the disk is larger than the diffusivity. In axisymmetric simulations we have both viscosity and

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Fig. 1 Snapshots from axisymmetric simulations of conical winds. The background shows the matter flux with light color corresponding to higher flux. The lines are magnetic field lines. A color version of this plot can be found in the electronic version of this book and in Appendix A (Fig. A.12)

diffusivity incorporated in the code, both in ˛ - prescription [30]. The coefficients ˛v and ˛d control these processes [26, 32]. We investigate a wide range of parameters: 0:01 < ˛v < 1 and 0:01 < ˛d < 1 and choose ˛v D 0:03 and ˛d D 0:1 as a main case. We assume that after period of low accretion rate the disk matter comes to the region from the boundary. Matter efficiently bunches field lines and in our case ˛v > ˛d this configuration exists for a long time. The disk matter comes close to the star, is stopped by the magnetosphere, and part of it moves into persistent conical outflows (see Fig. 1). Our simulations are dimensionless. As an example we chose parameters of the typical CTTS with mass M D 0:8 Mˇ , R D 2 Rˇ , magnetic field B D 1 kG, period P D 5:4 days. In Figs. 1–3 the inner boundary corresponds to two radii of the star. We accepted this choice of units so as to compare results with the propeller case (see Sect. 4) where the inner boundary is a factor of two smaller. Analysis of conical winds done by Romanova et al. [27] have shown that they are driven mainly by the magnetic pressure force (e.g., [16]) which is largest right above the disk and acts up to distances of about 12 stellar radii. Figure 2 shows typical parameters in a conical wind. Figure 2 shows that matter start to flow to a conical wind with very high azimuthal velocity, equal to Keplerian velocity at the base of the outflow (v  130 km/s in our main case). The poloidal velocity increases along the flow from few km/s right above the disk up to vp D 40–60 km/s at larger distances. Azimuthal velocity remains larger than poloidal velocity inside the simulation region. In the conical wind matter flows into a relatively narrow shell and the cone has an opening angle,  D 30ı –40ı . This may be explained by the fact that the magnetic pressure force acts almost vertically. This may also explain frequent events of reconnection of the inflating magnetic field lines in the outflow. We note that in addition to the main conical wind there is matter acceleration along magnetic field lines closer to the axis. The low-density matter is accelerated up to hundreds of km/s right near the

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Fig. 2 Typical flow in conical winds (at t D 380 days). The background shows matter flux, lines are selected field lines, arrows are proportional to velocity. The numbers show poloidal vp and total vt ot velocities and number density at sample places of the simulation region. A color version of this plot can be found in the electronic version of this book and in Appendix A (Fig. A.13)

Fig. 3 Two components of winds from slowly rotating star are labeled. A color version of this plot can be found in the electronic version of this book and in Appendix A (Fig. A.14)

star and may be important in explanation of some highly blue-shifted spectral lines which form near CTTSs. Matter which is accelerated in this region may come from the star, or may be partially captured from the main accretion flow. Figure 3 shows two components of the flow around a slowly rotating star.

2.2 The Fluxes of Matter and Angular Momentum The fluxes of matter and angular momentum flowing to or out from the star and fluxes flowing with conical winds through the surface with radius R D 0:1 AU were calculated. Figure 4 shows that matter flux to the wind is only several times

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Fig. 4 Left panel: matter flux to the star MP st ar and to conical wind MP wi nd (calculated at the radius R D 0:1 AU) as function of time. Right panel: same but for a shorter time-interval

smaller than that to the star, MP wi nd 0.2–0.3 MP star . The matter flux going to the wind varies, which is connected with frequent events of reconnection of the magnetic flux. It is often the case that matter is outbursted to the conical winds in an oscillatory regime, in particular if ˛v and ˛d are not very small, ˛v;d 0.1–0.3. If the diffusivity is small, ˛d D0.01–0.03, then outbursts to winds are sporadic and occur with a longer time-scale. Analysis of the angular momentum shows that in the case of a slowly rotating star the star spins-up by accreting matter (through magnetic torque at the surface of the star, e.g. [25]). Conical winds carry away part of the angular momentum of the disk (0.5 in this example), however a star may spin-up or spin-down depending on P . It spins-up in our example of a slowly rotating star. We also checked the case of very slow rotation, P D 11 days, and observed that persistent conical winds form in this case as well.

2.3 3D Simulations We performed exploratory 3D MHD simulations of conical winds in the case where the dipole magnetic field is misaligned relative to rotational axis by an angle  D 30ı . Compared with the axisymmetric simulations, the accretion disk is situated at r > 10R and the simulation region is much larger. Viscosity is incorporated in the code and we chose ˛vis D 0:3 while the diffusivity is not incorporated and is only numerical (small, at the level ˛d D 0:01–0.02 at the disk-magnetosphere boundary). We observed that the disk moved inward, bunched field lines and formed conical winds. Figure 5 shows that conical winds are approximately symmetric about rotation axis. There is however enhancement in the density distribution inside conical winds which is associated with a spiral wave generated by the misaligned dipole. Recent 3D modeling have shown that at a wide range of parameters matter penetrates through the magnetosphere due to interchange instability [24, 13]. 3D simulations of conical wind show that formation of conical winds occurs at larger distances from the star and are not influenced by instabilities.

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Fig. 5 Conical winds obtained in 3D MHD simulations for  D 30ı . Left panel: density distribution and sample field lines in the ˝-plane. Right panel: same but in the perpendicular plane. A color version of this plot can be found in the electronic version of this book and in Appendix A (Fig. A.15)

a

b Strong Wind

RT Outburst: . Macc~10–5 M. yr–1

c CO Wind

RT Approach to Equilibrium: . –7 Macc~10 M. yr–1

RT Equilibrium: . –7 Macc~10 M. yr–1

Fig. 6 Schematic model of an Exor V1647 Ori. During the outburst the accretion rate is enhanced so that the magnetospheric radius Rm decreases and the magnetic field lines were bunched (a). This results in a fast, hot outflow. As the accretion rate decreases, the disk moves outward and this results in a slower, cooler CO outflow (b). Further decrease in the accretion rate leads to a quiescence state where the production of warm outflows stops (c). From Brittain et al. [3]

3 Enhanced Accretion and Outflows CTTSs are strongly variable on different time-scales including a multi-year scale [11, 10]. This is connected with variation of the accretion rate through the disk which may lead to the enhancement of outflows (e.g., [5]). Simulations have shown that the bunching of field lines by the new matter after period of the low-density accretion may lead to quite long outburst of matter to the conical winds and may be the reason for formation of micro-jets in the CTTSs. If CTTS is in a binary system, then an accretion rate may be episodically enhanced due to interaction with the secondary star. Events of fast, implosive accretion are possible due to thermal instability or global magnetic instability, where the accretion rate is enhanced due to the formation of disk winds [17]. Enhanced accretion may lead to outbursts in EXors, where the accretion rate increases up to 105 Mˇ /yr and strong outflows are observed. Brittain et al. [3] reported on the outflow of warm gas from the inner disk around EXor V1647 observed in the blue absorption of the CO line during the decline of the EXor activity. He concluded that this outflow is a continuation of activity associated with early enhanced accretion and bunching of the field lines (see Fig. 6). In our main example of a CTTS the disk stops at Rm D 2:4R . In EXors, we take the radius of a star at the Figs. 1–3 equal to the inner boundary, so that the

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disk stops much closer to the star, Rm D 1:2R . Then all velocities are a factor 1.4 higher and densities a factor of 32 higher (compared to Figs. 2 and 7), and matter fluxes in Figs. 4 and 9 are a factor of 11 higher than in the main example relevant to CTTSs.

4 Outflows in the “Propeller” Regime In the propeller regime the magnetosphere rotates faster than inner region of the disk. This occurs if the co-rotation radius Rcr D .GM=˝2 /1=3 is smaller than magnetospheric radius Rm (e.g., [18]). Young stars are expected to be in the propeller regime in two situations: (1) At the early stages of evolution (say, at T < 106 years), when the star formed but did not have time to spin-down, and (2) at later stages of evolution, such as at CTTS stage, when the star is expected to be on average in the rotational equilibrium state (e.g., [15]) but variation of the accretion rate leads to variation of Rm around Rcr , where Rcr < Rm is possible. We performed axisymmetric simulations of accretion to a star in the propeller regime, taking a star with the same parameters as in case of conical winds, but with period P D 1 day [26,32]. We chose ˛v D 0:3 and ˛d D 0:1 and thus bunched the field lines to the Xtype configuration. We observed that in addition to conical wind there is a fast axial jet (see Fig. 7) so that the outflow has two components (Fig. 8). The conical wind in this case is much more powerful – it carries most of the disk matter away. The axial jet carries less mass, but it is accelerated to high velocities. Acceleration occurs due to the magnetic pressure of the “magnetic tower” which forms above the star

Fig. 7 Outflows in the propeller regime. The background shows matter flux, lines are selected field lines, arrows are proportional to velocity. Labels show total velocity and density at sample points. A color version of this plot can be found in the electronic version of this book and in Appendix A (Fig. A.16)

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Fig. 8 Two components of outflows in the propeller regime. A color version of this plot can be found in the electronic version of this book and in Appendix A (Fig. A.17)

Fig. 9 Left panel: matter fluxes to the star MP st ar and to the conical wind MP wi nd (calculated at R D 0:1 AU) as function of time. Right panel: same but for a shorter time-interval

as a result of winding of magnetic field lines of the star. Outbursts to conical winds occur sporadically with a long time-scale interval (see Fig. 9) which is connected with the long time-scale interval of accumulation and diffusion of the disk matter through the magnetosphere of the star (see also [9, 7]). These propeller outflows were obtained in conditions favorable for such a process: when the star rotated fast and an X-type configuration developed. Future simulations should be done for the case of propeller-driven outflows from slower rotating CTTS. Collimation of conical winds may occur at larger distances from the star for example, by disk winds (e.g., [12, 8, 19]).

5 Conclusions We discovered a new type of outflows – conical winds – in numerical simulations where magnetic field lines are bunched into an X-type configuration. In many respects these winds are similar to X-winds proposed by Shu and collaborators (e.g., [28]): (1) They both require bunching of the field lines; (2) They both have

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Fig. 10 Modeling of the Hˇ line in RW Aurigae led to the conclusion that a conical shaped wind with opening angle 30–40ı and a narrow annulus gives the best match to the observations of this line (from [1])

high rotation of the order of Keplerian rotation at the base of outflow, and gradual poloidal acceleration; (3) They both are driven by magnetic force. However, there are a number of important differences: (1) Conical winds flow in a thin shell, while X-winds flow at different angles below the “dead zone”; (2) Conical winds form around stars of any rotation rate including slow rotation, and do not require the fine tuning of angular velocity of the inner disk to that of magnetosphere; (3) Conical winds are non-stationary: the magnetic field constantly inflates and reconnects; (4) Conical winds carry away part of the angular momentum of the inner disk and are not responsible for spinning-down the star, while X-winds are predicted to take away angular momentum from the star; (5) In conical winds there is a fast component of the flow along field lines threading the star. Some of these differences, such as nonstationarity of conical winds is connected with natural restrictions of the stationary model of X-winds. Conical winds can explain conical shape of outflows near young stars of different type (CTTSs, EXors, Type I objects) which have been recently resolved. In another example, Alencar et al. [1] analyzed blue-shifted absorption of Hˇ line in RW Aurigae and concluded that conical shape wind with opening angle 30–40ı and narrow annulus gives best match to the observations of this line (see Fig. 10). In the propeller regime the flow has two components: (1) a rapidly rotating, relatively slow, dense conical wind, and (2) a fast, lower density axial jet where matter is accelerated by magnetic pressure up to hundreds of km/s very close to the star. Young stars of classes 0 and I may be in the propeller regime and can lose most of their angular momentum by this mechanism [26]. Or any slower rotating magnetized stars may enter the propeller regime if the accretion rate becomes sufficiently low and the magnetospheric radius becomes larger than the corotation radius. The last possibility requires additional numerical simulations and analysis. Acknowledgements The authors were supported in part by NASA grant NNX08AH25G and by NSF grants AST-0607135 and AST-0807129. MMR thanks NASA for use of the NASA High Performance Facilities. AVK and GVU were supported in part by grant RFBR 06-02016608, Program 4 of RAS. MMR and RVEL thank the organizers for a very interesting meeting and MMR is grateful to the organizers for the generous support.

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References 1. Alencar, S.H.P., Basri, G., Hartmann, L., Calvet, N. 2005 A&A, 440, 595 2. Bacciotti, F., Mundt, R., Ray, T.P., Eisl¨offel, J., Solf, J., Camezind, M, ApJ, 537, L49 3. Brittain, S., Simon, T., Rettig, T.W., et al. 2007, Star-Disk Interaction in Young Stars, IAU Symposium No. 243, ed. J. Bouvier & I. Appenzeller, p. 223 4. Cai, M.J., Shang, H., Lin, H.-H., & Shu, F.H. 2008, ApJ, 672, 489 5. Cabrit, S., Edwards, S., Strom, S.E., & Strom, K.M. 1990, ApJ, 354, 687 6. Edwards, S., Fischer, W., Hillenbrand, L., Kwan, J. 2006, ApJ, 646, 319 7. Fendt, C. 2009, ApJ, 692, 346 8. Ferreira, J., Dougados, C., & Cabrit, S. 2006, A&A, 453, 785 9. Goodson, A.P., Winglee, R.M., & B¨ohm, K.-H. 1997, ApJ, 489, 199 10. Grankin, K.N., Melnikov, S.Yu., Bouvier, J., Herbst, W., Shevchenko, V.S. 2007, A&A, 461, 183 11. Herbst, W., Herbst, D.K., Grossman, E.J., Weinstein, D. 2004, AJ, 108, 1906 12. Konigl, A., & Pudritz, R. E. 2000, Protostars and Planets IV, Mannings, V., Boss, A.P., Russell, S. S. (eds.), University of Arizona Press, Tucson, p. 759 13. Kulkarni, A., & Romanova, M.M. 2008, MNRAS, 386, 673 14. Kwan, J., Edwards, S., & Fischer, W. 2007, ApJ, 657, 897 15. Long, M., Romanova, M.M., & Lovelace, R.V.E. 2005, ApJ, 634, 1214 16. Lovelace, R.V.E., Berk, H.L., & Contopoulos, J. 1991, ApJ, 379, 696 17. Lovelace, R.V.E., Romanova, M.M., & Newman, W.I. 1994, ApJ, 437, 136 18. Lovelace, R.V.E., Romanova, M.M., & Bisnovatyi-Kogan, G.S. 1999, ApJ, 514, 368 19. Matsakos, T., Tsinganos, K., Vlahakis, N., Massaglia, S., Mignone, A., Trussoni, E. 2008, A&A, 477, 521 20. Najita, J.R., & Shu, F.H. 1994, ApJ, 429, 808 21. Pyo, T.-S., Hayashi, M., Kobayashi, N., et al. 2003, ApJ, 649, 836 22. Pyo, T.-S., Kobayashi, N., Hayashi, M., et al. 2003, ApJ, 590, 340 23. Ray, T., Dougados, C., Bacciotti, F., Eisl¨offel, J., & Chrysostomou, A. 2007, Protostars and Planets V, B. Reipurth, D. Jewitt, and K. Keil (eds.), University of Arizona Press, Tucson, p. 231 24. Romanova, M.M., Kulkarni, A.K., & Lovelace, R.V.E. 2008, ApJ, 673, L171 25. Romanova, M.M., Ustyugova, G.V., Koldoba, A.V., & Lovelace, R.V.E. 2002, ApJ, 578, 420 26. Romanova, M.M., Ustyugova, G.V., Koldoba, A.V., & Lovelace, R.V.E. 2005, ApJ, 635, 165L 27. Romanova, M.M., Ustyugova, G.V., Koldoba, A.V., & Lovelace, R.V.E. 2009, MNRAS, in press 28. Shu, F., Najita, J., Ostriker, E., Wilkin, F., Ruden, S., Lizano, S. 1994, ApJ, 429, 781 29. Shu, F.H, Galli, D., Lizano, S., Glassgold, A.E., & Diamond, P.H. 2007, ApJ, 665, 535 30. Shakura, N.I., & Sunyaev, R.A. 1973, A&A, 24, 337 31. Takami, M., Beck, T.L., Pyo, T.-S., McGregor, P., Davis, C. 2007, ApJ, 670, L33 32. Ustyugova, G.V., Koldoba, A.V., Romanova, M.M., & Lovelace, R.V.E. 2006, ApJ, 646, 304

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Simulating the Launching of YSO Jets Claudio Zanni

Abstract Different numerical models for the launching of jets from Young Stellar Objects (YSO) are presented. I will show numerical magnetohydrodynamic (MHD) simulations of outflows launched from the accretion disk (disk winds), from the stellar surface (stellar winds) and from the region of interaction between the stellar magnetosphere and the accretion disk (magnetospheric ejections). I will characterize the dynamical properties of each component and outline its impact on the star formation process.

1 Introduction The star formation process is often accompanied by the presence of outflows. For instance, accreting “classical” T Tauri stars (CTTS) often display collimated jets on scales of 10–100 AU propagating with a typical speed of the order of the

C. Zanni () Laboratoire d’Astrophysique de Grenoble, 414 Rue de la Piscine, BP 53, F-38041 Grenoble, France e-mail: [email protected]

K. Tsinganos et al. (eds.), Protostellar Jets in Context, Astrophysics and Space Science Proceedings, DOI 10.1007/978-3-642-00576-3 20, c Springer-Verlag Berlin Heidelberg 2009 

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escape velocity from the central star (200 km s1 ). The clear connection between ejection and accretion [3, 9] strongly supports the idea of accretion-driven magnetohydrodynamic (MHD) outflows. On the other hand, the location of the launching region is still a matter of debate: are the jets launched from the protostar, from an extended region of the accretion disk or from the interface between the stellar magnetosphere and the disk? In this short contribution I will briefly present numerical MHD models of these three different scenarios: the main features of extended disk-winds will be depicted in Sect. 2, stellar winds will be presented in Sect. 3 while magnetospheric ejections will be described in Sect. 4. In order to better understand the dynamical link between accretion and ejection, all of these models consider both the physics of the accretion disk and of the outflows.

2 Disk Winds The idea of magnetocentrifugal jet acceleration from Keplerian accretion disks was originally proposed by Blandford and Payne [1]. Analytical studies dealing with the connection between accretion and ejection [7] have shown that a powerful1 stationary launching requires a magnetic field around equipartition with the thermal energy of the disk (B 2 =4  P ) and a rather high level of turbulent resistivity to balance the field advection and twisting (˛m  1, parametrizing the disk resistivity as m D ˛m VA H , where VA is the Alfv´en speed and H the disk thickness). Numerical MHD simulations ([5, 18], see Tzeferacos’ contribution in this volume) confirm these conclusions showing in addition that non-stationary acceleration can still happen for lower values of ˛m  0:1 (see Fig. 1). For instance, the numerical solution shown in the central panel of Fig. 1 is characterized by a mass ejection/accretion ratio MP jet =MP acc  0:1, terminal speeds of a few hundreds km s1 and a magnetic lever arm around rA =r0  3. The launching region coincides with the equipartition zone and it can extend for more than one decade from the inner disk radius rin . Such an outflow can extract essentially all the accretion power Pacc D GM? MP acc =2rin , leaving less than 10% of it to be dissipated and radiated inside the disk. The jet torque represents also the dominant mechanism to extract angular momentum and drive the accretion process. It has been shown [8] that an outflow with these characteristics launched from some radial extension of the disk (from 0.1 to 1-2 AU) is the best candidate to reproduce the dynamical properties of observed T Tauri jets, in terms of terminal speed, collimation and (possible) rotation speed. If the launching zone coincides

1

Here I define a powerful disk wind as an outflow which extracts most of the accretion power and supplies the main torque to drive accretion. Weak ejection is still conceivable in weakly magnetized disks, where the main accretion torque is plausibly provided by the turbulent “viscous” transport which locally dissipates the accretion power. Nevertheless, these outflows require some additional mechanism acting at the disk surface (heating or diffusion) to provide the mass load to the wind.

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Fig. 1 Left panels: axisymmetric simulations of magnetized accretion-ejection structures [18] characterized by two different disk resistivities; ˛m D 0:1 (left) and ˛m D 1 (right). In both cases the accretion disk has B 2 =4 D 0:6 P . Right panel: time evolution of the two-sided ejection efficiency. Going from the solid to the dashed line the resistivity parameter decreases from ˛m D 1 to ˛m D 0:1. In these units, the Keplerian period at the inner disk radius is 2

with the equipartition region, then the study of the possibility of advecting this field from the outer “standard” accretion disk (see Lovelace and Murphy’s contributions in this volume) and of local dynamo processes becomes crucial. One caveat dealing with this kind of numerical simulations must be finally pointed out: due to problems of numerical diffusion at the disk surface, the simulations tend to overestimate the mass outflow rates. The ejection efficiencies measured normally in numerical models (10%) require the presence in exact analytical solutions of an extra heating term at the disk surface to increase the mass loading of the outflow [4].

3 Stellar Winds Outflows can be launched from the surface of the protostar along the opened field lines of the stellar magnetosphere. Since CTTS are “slow” rotators (their period of rotation, around 3–10 days, is less than 10% of their break-up speed [2]) their rotational energy cannot be used to launch centrifugally-driven stellar winds: the initial thrust must come from a pressure gradient which can be of thermal origin [15] or associated with turbulent Alfv´en waves [6]. Moreover, it has been suggested that this energy input is in some way related to the energy deposited by accretion on the stellar surface [12]. MHD simulations taking into account simultaneously the accretion onto the stellar surface and the emerging outflows must include in the computation the evolution of the stellar magnetosphere and its interaction with the surrounding accretion disk. One example of such a simulation is shown in Fig. 2: the numerical experiment, performed with the PLUTO code [14], models a protostar (M? D 0:5Mˇ , R? D 2Rˇ ) rotating at 10% of its breakup speed (P? D 4:5 days) with an (initially) dipolar magnetosphere (B? D 800 G). The star is surrounded by a viscous and resistive

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Fig. 2 Left panel: outcome of an axisymmetric simulation of the interaction of an accretion disk with a dipolar stellar magnetosphere. The snapshot is taken after 92 periods of rotation of the protostar. The dot-dashed line marks the magnetic surface anchored at the corotation radius, arrows indicate the velocity vectors and the Alfv´en surface of the stellar wind is marked by a dotted line. Right panel: time evolution of the ratio between the mass outflow rate of the wind and the mass accretion rate measured on the surface of the star (solid line, left scale); the average magnetic lever arm of the wind is also plotted (dot-dashed line, right scale). Time is given in units of the stellar period of rotation

(˛v D ˛m D 1) accretion disk (MP acc  108 Mˇ yr1). The stellar wind visible in Fig. 2 is thermally driven, having an enthalpy comparable to the gravitational potential energy at the surface of the star: this enthalpy corresponds to a temperature T D 106 K, which poses a serious cooling problem [13]. Its mass outflow rate is around 1% of the accretion rate measured at the surface of the star (the oscillations visible in Fig. 2 are due to a mismatch between the magnetic and viscous torques controlling the accretion in the inner part of the disk). The thermal power driving the outflow is about 2% of the accretion power and it can push the jet up to a terminal speed 250 km s1 . In spite of being thermally driven, most of the energy is transported along the stellar wind by the Poynting flux: as a matter of fact, the wind is characterized by a huge magnetic lever arm (rA =R?  18). This lever arm deter˝? ) which mines an enhanced torque braking the stellar rotation (JPwind D MP wind rA2p corresponds to about 20% of the accretion spin-up torque (JPacc D MP acc GM? rin ). This is a very important effect: the slow rotation period of CTTS requires in fact an efficient mechanism of angular momentum removal not only during the embedded phase, but also during the T Tauri stage of evolution, when the rotation period seems to stay constant [10], despite the fact that the protostar is still contracting and accreting. A more efficient spin-down torque would nevertheless require a higher mass ejection efficiency (10%) posing an even more serious energetic/cooling problem. It must be pointed out also that, at least in this simulation, the wind torque is more efficient than the torque exerted by the magnetosphere connecting the star and the disk beyond the corotation radius (Rco D 4:64R? ), which is usually invoked as an effective mechanism to brake the stellar rotation [11].

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4 Magnetospheric Ejections If the stellar magnetosphere is strongly coupled to the disk material, the differential rotation between the star and the disk creates a strong toroidal component of the magnetic field whose pressure can inflate and open the closed structure of the magnetosphere [17]. The opening of the magnetospheric field lines can lead to reconnection events and to the ejection of plasmoids. This phenomenon is clearly visible in the simulation shown in Fig. 3, which is characterized by the same parameters of the numerical experiment presented in Sect. 3 except for a smaller disk resistivity (˛m D 0:1), i.e. a stronger coupling between the disk and the magnetic field. Since they move ballistically at 45ı and they are not confined by any external agent, these episodic ejections are not a good candidate to explain the dynamical features of T Tauri jets. On the other hand the unsteady magnetospheric ejections can represent another mechanism to remove angular momentum from the central parts of the system, thus helping the spin-down of the star. The initial acceleration happens along field lines which are still connecting the star with the disk: the angular momentum is extracted therefore both from the star and the disk, accumulated on the tip of the highly deformed magnetospheric field lines and then released in the reconnection event, as in a huge magnetic slingshot. As it is shown in Fig. 3 even when these outflows are characterized by a small mass ejection efficiency the torque exerted on the accretion disk is stronger than a Keplerian torque, i.e. the torque needed to accrete from one Keplerian orbit to another. This type of outflows can be therefore an efficient mechanism to remove angular momentum from the disk before it is accreted onto the star.

Fig. 3 Left panel: outcome of an axisymmetric simulation of the interaction of an accretion disk with a dipolar stellar magnetosphere showing the episodic ejection of plasmoids. The snapshot is taken after 80 periods of rotation of the protostar. Right panel: time evolution of the mass ejection efficiency (dot-dashed line) and the ratio between the ejection torque and the accretion torque (solid line) characterizing the magnetospheric ejections. The “Keplerian” torque necessary to accrete from one Keplerian orbit to another is plotted with a dotted line. Time is given in units of the stellar period of rotation

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It can be finally noticed that matter can be also ejected along the opened magnetospheric field lines which are now threading the disk, as originally proposed for the “X-wind” scenario [16]: on the other hand, since the size of the equipartition region is very small, the amount of matter and energy extracted along these magnetic surfaces is negligible.

5 Summary and Conclusions I have presented numerical MHD simulations of different scenarios proposed to explain the launching of YSO jets. In agreement with [8], disk winds magnetocentrifugally launched from a sizeable fraction (0.1–1 AU) of the Keplerian accretion disk seem to be the best candidate to explain the main dynamical features of observed T Tauri jets, in terms of terminal speed, collimation and rotation. Pressuredriven stellar winds are not rotating fast enough and require an energy input of still unknown origin. On the other hand, they can contribute to the spin-down of the protostar. Episodic magnetospheric ejections move ballistically at 45ı and are not confined, but represent another mechanism to remove angular momentum from the central parts of the star-disk system.

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14. Mignone, A., Bodo, G., Massaglia, S. et al.: PLUTO: A Numerical Code for Computational Astrophysics. Astrophys. J. Suppl. Ser., 170, 228–242 (2007) 15. Sauty, C., Trussoni, E. & Tsinganos, K.: Nonradial and nonpolytropic astrophysical outflows. V. Acceleration and collimation of self-similar winds. Astron. Astophys., 389, 1068–1085 (2002) 16. Shu, F., Najita, J., Ostriker, E., Wilkin, F., Ruden, S. & Lizano, S.: Magnetocentrifugally driven flows from young stars and disks. 1: A generalized model. Astrophys. J., 429, 781–796 (1994) 17. Uzdensky, D. A., K¨onigl, A. & Litwin, C.: Magnetically Linked Star-Disk Systems. I. Forcefree Magnetospheres and Effects of Disk Resistivity. Astrophys. J., 565, 1191–1204 (2002) 18. Zanni, C., Ferrari, A., Rosner, R., Bodo, G. & Massaglia, S.: MHD simulations of jet acceleration from Keplerian accretion disks. The effects of disk resistivity. Astron. Astophys., 469, 811–828 (2007)

On the Effect of Stellar Wind Braking onto the Central Object Christophe Sauty, Noemie Globus, Zakaria Meliani, Kanaris Tsinganos, Nektarios Vlahakis, and Edo Trussoni

Abstract Stellar winds seem to be very efficient at removing angular momentum from stars. By means of analytical axisymmetric solutions of the ideal MHD equations for steady outflows, we show via a specific example how collimated stellar winds can brake Weak T Tauri stars in a reasonable time. This result can be generalized to Classical T Tauri stars provided that part of the accreted angular momentum is removed by the inner disk wind. We also extend briefly to Kerr metrics the self similar MHD solutions for relativistic flows and conjecture that relativistic outflows may efficiently slow down spinning black holes at the center of Active Galactic Nuclei or microquasars. C. Sauty (), N. Globus, and Z. Meliani Observatoire de Paris, LUTH, 92190 Meudon e-mail: [email protected] Z. Meliani Centre for Plasma Astrophysics, KU Leuven (Leuven Mathematical Modeling and Computational Science Center), Belgium K. Tsinganos, and N. Vlahakis IASA and Section of Astrophysics, Astronomy and Mechanics, Department of Physics, University of Athens, Panepistimiopolis, 15784 Zografos, Athens, Greece E. Trussoni INAF/Osservatorio Astronomico di Torino, via Osservatorio 20, 10025 Pino Torinese, Italy K. Tsinganos et al. (eds.), Protostellar Jets in Context, Astrophysics and Space Science Proceedings, DOI 10.1007/978-3-642-00576-3 21, c Springer-Verlag Berlin Heidelberg 2009 

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1 Introduction: Multicomponent Outflows Low mass stars in their formation process exhibit highly collimated outflows. This collimation seems to be induced by the “hoop stress” of their own magnetic field. Their association with accretion disks has put forward the existence of disk powered winds. However, the presence of a stellar wind in the central part of the system is also demanded by several other reasons. First, simulations show that disk winds alone are usually over collimated to be realistic. Including an inner stellar jet can open the jet channel sufficiently, provided that the stellar mass loss is comparable to the disk mass loss [5]. Second, the current in the disk closes through the axis, as shown by simulations and analytical disk wind models. Hence, this means that the global electric circuit of the jet closes inside the stellar component. Third, recent simulations ([3]; see also these proceedings) have shown that variable conditions at the base of the inner stellar wind can induce large scale shocks at large distances in the outflow, supporting the idea that knots in HH objects may be induced by the stellar variability itself. In fact it has been proposed (e.g. [7]) that during the formation sequence of low mass stars, the wind starts at the inner edge of the accretion disk during evolutionary phase 0. Then from phase I to III, the wind proceeds gradually towards the central star until the disk finally disappears and only the stellar wind remains. During this evolution the mass loss rate reduces and the jet speed increases. We plot in Fig. 1 two possible configurations of the magnetic field close to the central star. The first (Fig. 1a) corresponds to the configuration of the self-similar solution used in the next section while the second to the classical X-wind model (Fig. 1b). In the late stage of the evolution the disk wind and the stellar wind components play a comparable role. Both configurations should be surrounded by a Keplerian disk wind at larger distances. Analytical modeling of disk and stellar jets have been studied independently and shown to be numerically stable [2]. Only recently two component simulations using analytical solutions as initial conditions have been performed [3]. A combination of analytical solutions is stable and they can be used to illustrate the jet physics. Following this argument, we find useful to study stellar jets using meridional self similarity, assuming that the external part of the jet is modeled by a radially self similar disk wind solution.

a

b

Fig. 1 Sketch of the poloidal fieldlines of two rather similar magnetic topologies. In (a) is shown the magnetic configuration of meridionally self-similar MHD winds including the inner disk and in (b) the magnetic configuration of the X-wind models

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As a matter of fact, this approach may also be adapted to relativistic jets. In the case of AGN jets, it is usually assumed that a spine jet from the central black hole or its magnetosphere is surrounded by an external disk wind. If the spine jet is made of electron positron pairs, and in order for these pairs to survive, it is suggested that the inner jet (corresponding to the meridionally self similar solution) is not directly in contact with the outer part (the one modeled by radially self similar solutions) conversely to the case of YSO jets.

2 An Example of an Efficiently Braking Stellar Wind Solution We have proposed [7] to use the meridionally self similar solutions (see [8] for more details) to model the jets of late type T Tauri Stars like RY Tau. For this RY Tau system, we know the radius of the star (2 solar radii) and the rotation period of (24 days). We also know the asymptotic speed of the jet (200 km s1 ) and the density before and after the shock (103 cm3 and 104 cm3 respectively). Assuming that the jet starts underpressured and mostly thermally confined, it becomes overpressured and purely magnetically confined around a few tens of AU. Its effective temperature decreases continuously from 105 K at the stellar surface to 103 K asymptotically. We recall that by rescaling the various parameters, we were able to reproduce, with the same numerical solution, various T Tauri jets. For this solution, the total mass loss rate reaches 109 Mˇ if the solution is extended up to 3 stellar radii within the disk. The stellar jet itself has a mass loss rate of only 1010 Mˇ . We may have here a clue on the difference between jets from classical T Tauri Stars (noted CTTS hereafter) and weak T Tauri Stars (noted WTTS hereafter). In both cases we conjecture that the stellar jet is very similar. However, as WTTS are not connected to a disk, the total mass loss rate remains of the order of 1010 Mˇ , which is almost not detectable. Conversely, jets from CTTS like RY Tau would have contributions from the inner disk which give a mass loss rate 10 times higher. The time scale at which the star loses all its angular momentum is given by P is the rate of angular D LP where L is the angular momentum of the star and L L 7 momentum loss. For the stellar wind alone we get  10 yr which is indeed typical of the life time of WTTS before they reach the main sequence. This corresponds in the diagram of [1] (see Fig. 2) to the lower star exactly inside the domain of stellar winds. However if we include the inner disk wind, the time scale drops to  106 yr which indeed corresponds to the life time of the CTTS phase. In the same diagram this corresponds to the upper star which is close to the X-Wind. Thus we conjecture that during the CTTS phase the inner part of the disk can remove almost all angular momentum from the accreting plasma, before it reaches the star. The remaining angular momentum is lost by the star itself during the Weak T Tauri phase. We should add in conclusion that the solution we discuss here – as well as those published so far in the literature – are not the most powerful ones in terms of the braking efficiency. In [8], we made an asymptotic analysis and insisted on getting solutions with high asymptotic speeds compared to the Alfv´en speed in order to

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Fig. 2 Plot of the jet specific angular momentum versus its asymptotic speed, adapted from [1]. The lower star shows the location of our stellar jet solution. The upper star corresponds to the same solution but with the inclusion of the inner disk wind

finally get large jet radii in units of the Alfv´enic radius. Yet, if most of the acceleration and expansion of the wind occurs before the Alfv´en surface, this does not need to be the case. Meridionally self similar solutions exist with very large lever arms and small acceleration beyond the Alfv´en surface. In this case most of the collimation occurs before the Alfv´en surface. This makes the braking of the star by the stellar wind more efficient (see arrow in Fig. 2).

3 Extension to Relativistic Flows in Kerr Metrics Meridionally self-similar MHD models can be also used to study relativistic jets. Thus, models for non rotating black holes, using Schwarzschild metrics, have been constructed [4]. Such solutions can be applied to relativistic winds rotating at subrelativistic rotation frequency, emerging from a hot corona around the central black hole. They describe the inner spine jet component but not the overall jet which again needs to be associated with a disk wind. In fact, meridionally self similar models have given a simple energetic criterion for the collimation of jets [7]. According to this criterion, if there is an excess of volumetric energy along a non polar streamline with respect to the axis, the outflow collimates asymptotically into cylinders. There are however two contributions, one thermal measuring the efficiency of the pressure confinement and another one

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measuring the efficiency of the magnetic pinching ("). Using this criterion we can interpret the variation between radio quiet and radio loud galaxies. In Seyfert galaxies, we observe winds rather than jets. This may be associated to a lack of energy on the external lines, thus without collimation. In the frame of this interpretation, radio galaxies are instead efficient magnetic rotators. Pressure confinement seems to be at work in FRI jets because of the rich environment of the host galaxy. This thermal confinement could be also associated with a deceleration of the jet (see also [6]). Conversely, FRII jets are propagating in a poor environment and the strong collimation is self induced by the magnetic field. We have been able to generalize the same model to rotating black holes in Kerr metrics and deduce a criterion for collimation. The efficiency of the magnetic rotator measures the excess of magnetic energy on a non polar streamline that is not used to accelerate the wind. It can be estimated, at the base of the flow, as follows: "D

EPoynt:;o C EG C h2o ER;o C L.!0  !? / ; h2o L˝

(1)

where we have introduced the following notations: h2o L˝ is the energy of the magnetic rotator with L the angular momentum, ˝ the corotation frequency and h0 the Kerr radial line element, EPoynt:;o is the Poynting flux and ER;o is the rotational energy per particle. The term EG is similar to the corresponding one in the nonrelativistic case where it measures the excess or the deficit on a non polar streamline, compared to the polar one, of the gravitational energy per unit mass which is not compensated by the thermal driving (see, [4, 8]). The most interesting part is the new extra term L.!o  !? / which corresponds to the energy of the magnetic rotator in the dragging frame. L is still the plasma specific angular momentum but !o is the rotation frequency of the frame at the base of the flow close to the ergosphere. If the black hole entrainment is high we should have !o  ˝. !? corresponds to the frame rotation at the Alfv´en surface and this term is rather small if the magnetic lever arm is large, as we expect. In other words, in order of magnitude close to the black hole we have EPoynt:;o  L˝  L.!o  !? /. Thus the rotation of the Kerr black hole increases greatly the efficiency of the magnetic rotator. As in the classical case, " measures the efficiency of the magnetic rotator to collimate the flow. Thus if " > 0 we have an Efficient Magnetic Rotator (EMR) where magnetic collimation may dominate, while if " < 0 we have an Inefficient Magnetic Rotator (IMR) where collimation cannot be but of thermal origin. We see however that the rotation of the black hole helps the magnetic collimation of the spine jet at large distances. This is a direction which needs to be explored in the future.

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References 1. Ferreira, J., Dougados, C., Cabrit, S., A&A, 453, 785–796 (2006) 2. Matsakos, T., Tsinganos, K., Vlahakis, N. et al., A&A, 477, 521–533 (2008) 3. Matsakos, T., Massaglia, S., Trussoni, E. et al., A&A, submitted (2009) 4. Meliani, Z., Sauty, C., Vlahakis, N., Tsinganos, K., Trussoni, E., A&A, 447, 797 (2006) 5. Meliani, Z., Casse, F., Sauty, C. 2006b, A&A, 460, 1–14 (2006) 6. Meliani, Z., Keppens, R., Giacomazzo, B., A&A, 491, 321–337 (2008) 7. Sauty, C., in: Jets from Young Stars, Lecture Notes in Physics, Springer-Verlag, Volume 723, 209–224 (2007) 8. Sauty, C., Tsinganos, K., Trussoni, E., A&A, 348, 327–349 (1999)

Flaring Activity in Accretion Flows of Young Stellar Objects Fabio Reale

Abstract X-ray observations have shown extensive flaring activity in young stellar associations such as the Orion nebula. Observed flares are often very long and intense, and have been associated to very long magnetic loops, which may connect the stellar surface to the circumstellar disk. As such, these loops are candidate to be also the channel of star accretion from the disk, and one then wonders whether they flare during accretion flows. As a first attack to this question we have modelled in detail flares inside long coronal loops containing plasma at high density, comparable to that presumed for accretion flows. Preliminary results show that such flares would decay on time scales smaller than the observed ones.

F. Reale () Dipartimento di Scienze Fisiche & Astronomiche, Universit`a di Palermo, Sezione di Astronomia, Piazza del Parlamento 1, 90134 Palermo, Italy, INAF - Osservatorio Astronomico di Palermo “Giuseppe S. Vaiana”, Piazza del Parlamento 1, I-90134 Palermo, Italy e-mail: [email protected]

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1 Introduction Soft X-ray observations of the solar corona show that X-ray coronal flares are brightenings occurring in well-localized and limited regions of the corona, i.e. loops [1] or arcades of loops [2]. In the soft X-ray band we typically detect thermal emission of plasma at temperature about ten times that of the quiescent corona (10 MK) and with a density one hundred times higher [3]. Coronal flares are characterized by light curves with a steep rise phase and a more gradual decay, with timescales ranging from a few minutes to several hours. The general evolution of a flare is interpreted as triggered by a strong and fast heat pulse followed by the plasma cooling, which drives the decay phase. The cooling is due to two main mechanisms: the thermal conduction to the cooler chromosphere and the radiation from optically thin plasma, characterized by the respective timescales [4]: c D

3nc kB T0 L2 2=7 T07=2 r D

D

10:5nc kB L2

T05=2

 50

3=2 TM;7 3kB TM D 9  103 nM P .T / nM;10

nc;10 L29 5=2 T0;7

(1)

(2)

where L (L9 ) is the loop half-length (in units of 109 cm), T0 (T0;7 ) is the loop maximum temperature (in units of 107 K), nc (nc;10 ) is the particle density (1010 cm3 ) at the end of the heat pulse, D 9  107 (c.g.s. units) is the thermal conductivity, TM (TM;7 ) is the temperature at the time of the density maximum (107 K), nM (nM;10 ) the maximum density (1010 cm3 ), P .T / the plasma emissivity per unit emission measure. Both timescales have a direct or indirect dependence on the length of the loop where the flare occurs. This dependence has been used to infer the size of the flaring regions involved in stellar flares which cannot be resolved by present-day instruments [5]. The cooling times can be combined into a global thermodynamic cooling time [6], L L9 (3) s D 3:7  104 p D 120 p T0 T0;7 This timescale depends linearly on the loop length. The length of the flaring loop can then in principle be easily derived by inverting (3). However, the length can be largely overestimated if significant heating is released during the decay making the decay slower than expected from a pure cooling [7]. There is a way to solve this ambiguity. The decay in the density-temperature diagram follows a track with positive slope, from high to low values. It has been shown [8] that the slope of the track depends on the presence of significant heating during the decay, steeper for no heating, shallow for dominant heating. One can then constrain the presence of heating during the decay and thereby obtain a unique value of the loop length [7]. The presence of significant heating is typical of long-enduring events, which can involve multiple loop systems or even long arcades. On the other hand, flares

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characterized by steep decay tracks and therefore no heating are driven by a plasma coherent cooling which most probably occurs in single loops.

2 Modeling Flares on YSO Long-lasting and very intense flare events are frequently observed in young stellar objects. A large sample of such flares has been collected by Chandra in the Orion star-formation region [9] during the COUP enterprise. The most intense events have been systematically analyzed and long loop lengths diagnosed, of the order of 1012 cm. So long loops extend over several stellar radii and are candidate to connect the stellar surface to still-existing circumstellar disks. Such long magnetic tubes have been hypothesized to be the channel of accreting flows from the disk to the star. One of such long events on source 1,343 has been modelled in detail through time-dependent hydrodynamic loop modeling. The light curve shows a flare duration of about one day and a decay time of several tens of ks. Time-resolved spectral analysis has allowed very well constrained density and temperature diagnostics in the decay. The decay slope for this event is compatible with no heating in the decay and supports the occurrence in a single loop system. From the relevant formula we obtain a loop half-length very close to 1012 cm, which corresponds to several stellar radii. Hydrodynamic simulations allow to obtain further constraints, i.e. a very long heat pulse duration of 20 ks and a small loop aspect of about 2% (10% is a typical value for solar coronal loops). Simulations are performed by numerically solving the time-dependent hydrodynamic equations for a compressible fully-ionized plasma confined in a closed magnetic flux tube. The plasma is assumed to move and transport energy only along the magnetic field lines and can be described with a one-dimensional model. The flare is triggered with an input heating function composed by a time-dependence, a top-hat function, and a spacedependence, a Gaussian localized at the loop footpoints. The intensity of the heat pulse is set so to have a maximum temperature of the order of the observed one (the energy rate and the temperature are linked by the loop scaling laws, [10]). In order to solve the time-dependent equations, it is necessary to specify a complete loop atmosphere, including a chromosphere, transition region and corona, as initial condition. In that case, it was chosen an initial flux tube with the appropriate length and a relatively tenuous atmosphere, with a density of 109 cm3 at the loop apex. There is no real constraint from observations on the initial atmosphere, but it is wellknown that the flare evolution does not depend on it except for details, provided that the initial pressure is much lower than the pressure in the flare regime. The latter assumption is certainly reasonable for standard flaring regions, i.e. solar-like active region loops. If we further speculate that the long flaring loop diagnosed from the modeling connects the star to the disk, we are then considering a rarefied initial stardisk flux tube. This implies that no accretion flow is present in the tube when the flare occurs.

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3 Modeling Flares in Accreting Flux Tubes Our question is now whether it is possible that such a giant flare occurred in a stardisk flux tube where an accretion flow is present (Fig. 1). Our approach here is to repeat exactly the same hydrodynamic simulation but considering an initial tube with an accretion flow. The flare energy is by far larger than the energy involved in the accretion in the same time range and therefore we can neglect the accretion dynamics in this context. The characteristics of the accretion which instead makes the difference is the plasma density. Therefore, we approximate the accreting flow with an initially denser tube. The density involved in accretion has been estimated to be in the range 1011 –1012 cm3 [11, 12]. To be conservative, for our modeling we assume an initial density of 1011 cm3 at the loop apex, about the low range boundary but still two orders of magnitude larger than in the previous modeling. The energy input rate is here 5 times larger than before, to let the plasma reach the same flare temperature, in spite of for the larger plasma heat capacity. Figure 2 shows the light curve obtained from the new simulation of the flare triggered in the accreting channel (after the appropriate rescaling for the cross-section), as compared to the observed data. The rise, peak and early part of the decay are still are well-described by the simulation, the late decay, after about 30 ks, is not, whereas the previous modeling follows the decay quite accurately until about 60 ks. In the

Fig. 1 Cartoon of two possible sites of long and intense X-ray flares observed in YSO: a lowdensity magnetic flux tube interconnection the star surface to the circumstellar disk (left) and an analogous flux tube filled with a dense accretion flow (right) Fig. 2 X-ray light curve of the flare observed on source 1,343 of the COUP survey (data points) and that obtained from a hydrodynamic simulation starting from a magnetic flux tube containing a dense accretion flow (solid line)

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accreting flaring loop the decay drops too fast. There is a precise physical reason for this: in the standard modeling the density remains steadily below 1011 cm3 even at the peak. The value 1011 cm3 is already the starting value in the accretion case and it grows up to about twice at the flare peak. The implication is that the loop is never in equilibrium conditions and the density is always so high that the radiation cooling invariably dominates over the conduction cooling and is much faster than the standard loop cooling time. One may argue that in a longer loop or with significant residual heating the decay may be made slower again, but we have seen that the fast decay is due only to the high density, indipendently of the loop length, and that the density-temperature diagram appears to exclude the presence of heating in the decay. From this test case we conclude that it seems unlikely that the big and long flare on source 1,343 in the COUP survey is triggered in an accreting flux tube. This conclusion may apply to most of the long and intense X-ray flares observed on YSO, since dense accretion flows tend to shorten the flare duration. On the other hand, it has been recently shown with the aid of 2-D MHD simulation of flares in a magnetized star-disk system (Yelenina et al. 2008, this volume), that the opposite may occur: strong flares may trigger important accretion episodes. This perspective deserves more investigation in the next future. Acknowledgements The author acknowledges support from the Marie Curie Fellowship Contract No. MTKD-CT-2005-029768 of the project “Young stellar objects, their surroundings and jets: Advanced observational and MHD studies”.

References 1. S. Masuda, T. Kosugi, H. Hara, S. Tsuneta, Y. Ogawara, Nature, 371, 495 (1994). DOI 10. 1038/371495a0 2. M.J. Aschwanden, D. Alexander, Sol. Phys. 204, 91 (2001). DOI 10.1023/A:1014257826116 3. G. Peres, G.S. Vaiana, Memorie della Societa Astronomica Italiana, 61, 401 (1990) 4. F. Reale, A&A, 471, 271 (2007). DOI 10.1051/0004-6361:20077223 5. F. Reale, in Stellar Coronae in the Chandra and XMM-NEWTON Era, Astronomical Society of the Pacific Conference Series, vol. 277, ed. by F. Favata, J.J. Drake (2002), pp. 103–+ 6. S. Serio, F. Reale, J. Jakimiec, B. Sylwester, J. Sylwester, A&A, 241, 197 (1991) 7. F. Reale, G. Peres, S. Serio, A&A, 318, 506 (1997) 8. J. Jakimiec, B. Sylwester, J. Sylwester, S. Serio, G. Peres, F. Reale, A&A, 253, 269 (1992) 9. F. Favata, E. Flaccomio, F. Reale, G. Micela, S. Sciortino, H. Shang, K.G. Stassun, E.D. Feigelson, ApJS, 160, 469 (2005). DOI 10.1086/432542 10. R. Rosner, W.H. Tucker, G.S. Vaiana, ApJ, 220, 643 (1978). DOI 10.1086/155949 11. J. Robrade, J.H.M.M. Schmitt, A&A, 473, 229 (2007). DOI 10.1051/0004-6361:20077644 12. M.M. Jardine, S.G. Gregory, J.F. Donati, MNRAS, 386, 688 (2008). DOI 10.1111/j.1365-2966. 2008.13103.x

Similarities of the Launching Mechanism in Protostellar/AGN Jets Ryoji Matsumoto

Abstract We present the results of global 2D and 3D magnetohydrodynamic simulations of jet formation from a gas disk rotating around a central object. In a disk-star system, differential rotation twists magnetic loops connecting a star and its disk. As magnetic twist accumulates, the magnetic loops inflate and form current sheets inside the loops. Magnetic reconnection taking place in the current sheet can be the origin of X-ray flares observed in protostars. Numerical simulations using larger computing area revealed that the expanding magnetic loops form a magnetic tower. Magnetic reconnections taking place near the footpoints of the tower inject hot plasmoids into the tower. Less collimated outflow of cool gas emanates from the disk along the large-scale magnetic fields formed by the magnetic loop expansion. We also show that even when the large-scale poloidal magnetic fields do not exist at the initial state, they are generated by the buoyant rise of magnetic loops from the accretion disk. These magnetic loops are twisted, elongated, and form magnetic towers. Core-jet and outer wind structure is common both in AGNs and in protostars.

R. Matsumoto () Department of Physics, Graduate School of Science, Chiba University, 1-33 Yayoi-Cho, Inage-ku, Chiba 263-8522, Japan e-mail: [email protected]

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Fig. 1 Uchida-Shibata mechanism of jet formation from an accretion disk threaded by large-scale poloidal magnetic fields. Jets are launched and collimated by the Lorentz force created by twisting the magnetic fields anchored to the rotating disk

1 Introduction Jets and outflows are observed in various astrophysical objects such as active galactic nuclei (AGN), protostars, and in galactic black hole candidates. Uchida and Shibata [18] and Shibata and Uchida [17] carried out axisymmetric, twodimensional magnetohydrodynamic (MHD) simulations of jet formation from an accretion disk initially threaded by large-scale poloidal magnetic fields (Fig. 1). They showed that when the disk gas is coupled with the magnetic fields, the magnetic fields twisted by the disk rotation accelerate the disk gas by the Lorentz force. The outflows can be collimated toward the rotation axis by the pinch force. The terminal speed of the jet is close to the Keplerian rotation speed of the disk anchoring the large-scale magnetic fields. Matsumoto et al. [15] applied the Uchida-Shibata model of the jet formation to a geometrically thick disk (torus) around a super massive black hole in AGNs. A geometrically thick disk can be formed when the accretion rate is so small that the radiative cooling is negligible (hot, optically thin torus), or when the accretion rate MP exceeds the Eddington accretion rate MP Edd D LEdd =c2, where LEdd is the Eddington luminosity [16]. The former corresponds to the radio galaxies, and the latter may correspond to the luminous quasars. We would like to point out that the pre-existence of large-scale poloidal magnetic fields threading the disk is not essential for the production of jets and outflows. Large-scale open magnetic fields can be created by twisting the magnetic fields connecting the central object and its accretion disk, or by twisting the magnetic loops buoyantly rising from the disk (Fig. 2). In the following, we present the results of global MHD simulations of jet formation according to these scenarios.

2 Jet Formation from a Disk Initially Threaded by Large-scale Poloidal Magnetic Fields When an accretion disk is threaded by large-scale poloidal magnetic fields, the disk gas can be accelerated in bipolar directions by the magneto-centrifugal force [2] and the magnetic pressure gradient force. Global MHD simulations of this process

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a

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c

star-disk connection

buoyantly rising magnetic loops

Fig. 2 Origin of large-scale poloidal magnetic fields. (a) Poloidal magnetic fields are carried in by accreting gas. (b) Expansion of magnetic fields connecting the central object and the disk. (c) Twisting of the magnetic loops buoyantly rising from the disk

Fig. 3 Magnetic field lines (solid curves) and density distribution (grey scale) of the resistive MHD simulation of jet formation from a torus initially threaded by uniform magnetic fields parallel to the rotation axis. (left) initial state, (right) nonlinear stage

were first carried out by Uchida and Shibata [18] and Shibata and Uchida [17]. Kuwabara et al. [10] found that the jets are launched intermittently because the magneto-rotational instability (MRI; [1]) deforms magnetic field lines and drives magnetic turbulence inside the disk. Quasi-steady jets can be obtained by simulations including finite resistivity [10, 11, 3, 4]. Figure 3 shows a result of the axisymmetric resistive MHD simulations of the jet formation starting from a constant angular momentum torus threaded by uniform vertical magnetic fields. Kuwabara et al. [11] compared the results of the resistive MHD simulations with the theory of steady axisymmetric MHD outflows. According to Kudoh and Shibata [8] and Ustyugova et al. [19], when the resistivity is negligible, the following quantities should be conserved along a magnetic field line anchored to the disk; K D 4 jVp j=jBp j,  D rV  rB =K, ˝ D V =r  KB =.4 r/, S D log.P = /, E D Œ=. 1/P = C r2˝2=2CVp 2=2C.V =r ˝/2 r2=2, where Vp and Bp are poloidal velocity and poloidal magnetic field, respectively, V and B are azimuthal velocity and azimuthal magnetic field, respectively,

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5

0

K

S

E

0

1

2

3

Z

Fig. 4 Distribution of the conserved quantities along a magnetic field line depicted by a white curve in the left panel. Grey scale in the left panel shows temperature. Arrows show velocity vectors TIME = 45.502 10

8

poloidal velocity > poloidal Alfven speed

6 z

poloidal velocity > slow magnetosonic speed

4

2

0 –2 0

0.5

0 1

x

2 1.5

Fig. 5 A result of the 3D MHD simulation of jet formation using Cartesian grids. (left) Gray scale shows the ratio of gas pressure to magnetic pressure (plasma ˇ). Solid curves show magnetic field lines. (right) Arrows show velocity vectors

is the gravitational potential, is the density, and  is the specific heat ratio. Figure 4 shows these quantities along a magnetic field line depicted in the left panel. Although the entropy S increases due to the Joule heating, other quantities are conserved. These results indicate that the disk-jet system is approaching a quasi-steady state. Kudoh and Shibata [8] showed that when the poloidal magnetic field is small, outflows are accelerated between the Alfv´en point and the fast magnetosonic point by the magnetic pressure gradient force created by twisting the magnetic fields. Since the azimuthal magnetic field is accumulated in this region, the outflow can subject to the current driven kink instability. Figure 5 shows a result of the

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three-dimensional global MHD simulation using Cartesian grids (Kuwabara et al., 2009, in preparation). Outflow wiggles in the region where the poloidal velocity exceeds the poloidal Alfv´en speed. The kink instability, however, does not destroy the jet.

3 Jet Formation by Expansion of Magnetic Loops Connecting a Star and its Disk When a star and its disk is connected by magnetic loops, the magnetic fields will be twisted by the difference of the angular speed between the star and the disk and will inflate [12]. Figure 6 shows a result of the 2D axisymmetric resistive MHD simulation of the star-disk interaction [5]. The initial state is a disk threading the dipole magnetic field of the central star. We assumed anomalous resistivity which sets in when J = exceeds the threshold value, where J is the current density. We neglected the rotation of the central star. As the magnetic loops anchored to the star and the disk are twisted, they inflate and form a current sheet inside the loop. Magnetic reconnection taking place in the current sheet produces hot plasmoids ejected from the reconnection region, and heats the plasma to 10 KeV. This mechanism is similar to that of solar flares but takes place in much larger volume, thus can release more magnetic energy. Numerical results can explain the X-ray flares observed in protostars. The right panel in Fig. 6 shows the 3D structure of expanding magnetic loops. Since the magnetic field lines of the expanding magnetic loops have large inclination angle at the surface of the disk with respect to the rotation axis, dense, magneto-centrifugally driven wind emanates from the disk. The high speed hot outflow produced by magnetic reconnection is a low-density flow corresponding to the

Fig. 6 A result of 2D MHD simulation of protostellar flares. Solid curves show magnetic field lines. Gray scale shows the temperature distribution. The right panel depicts the 3D structure of the expanding magnetic fields

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z/rs

60 plasmoids 40 20 0 –20–10 0 10 20 r/rs

Fig. 7 Formation of a magnetic tower jet. Left panel shows the magnetic field lines. Right panel shows the temperature distribution. Solid curves show magnetic field lines projected onto the poloidal plane

optical jets, and the cool, dense outflow emerging from the disk may correspond to the high speed neutral winds observed in protostars. When we carry out the simulation of the star-disk magnetic interaction by using much larger simulation box, the expanding magnetic loops are collimated toward the rotation axis, and form a magnetic tower (e.g., [20, 6]). Figure 7 shows the results of axisymmetric 2D MHD simulation by Kato et al. [7]. The left panel shows the magnetic field lines, and the right panel shows that the magnetic reconnection taking place intermittently around the foot points of the magnetic tower injects hot plasmoids into the tower. When the central star is a neutron star, this mechanism creates a collimated, sub-relativistic (v  0:2c) jet.

4 Formation of Outflows from Turbulent Accretion Disks In this section, we consider an accretion disk initially has no global poloidal magnetic fields. It corresponds to a disk around galactic black hole candidates. The mechanism of the angular momentum transport which enables the accretion of such disks had long been a puzzle until Balbus and Hawley [1] pointed out the importance of MRI. When weak magnetic field exists in the disk, the disk gas can accrete by transporting angular momentum through the Maxwell stress produced by the MRI-driven turbulent magnetic fields. Figure 8 shows a result of the global 3D MHD simulation of the formation of an accretion disk starting from a torus threaded by weak azimuthal magnetic fields [13]. The general relativistic effects are simulated by using the Pseudo-Newtonian potential. Figure 8(a) shows the initial condition. Grey scale shows the density distribution, and the white curves show magnetic field lines. Since MRI grows even

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Fig. 8 A result of the 3D global MHD simulation of the formation of an accretion disk. (a) Density distribution at the initial state. (b) Quasi-steady state. (c) Enlargement of the innermost region of the accretion disk

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when the initial magnetic field is purely azimuthal, the disk becomes turbulent, and the disk gas infalls by transporting the angular momentum. After several rotation period of the initial torus, quasi-steadily accreting disk is formed. Solid curves in Fig. 9a show the magnetic field lines projected onto the poloidal plane. Grey scale shows the azimuthal magnetic fields. We found that large-scale poloidal magnetic fields are automatically produced in the coronal region above the innermost region of the disk [14]. These magnetic fields are formed by the expansion of the magnetic loops emerging from the disk by buoyancy. When the footpoints of the magnetic loops are anchored to the disk at different radius, the magnetic loop will be twisted by the difference of the angular speed at footpoints. When critical twist is accumulated, they inflate, and form large-scale poloidal magnetic fields (see Fig. refmatsufig:2(c)). This mechanism is identical to the mechanism of the formation of a magnetic tower. Figure 9b shows the isosurface where the vertical speed is 0.05 c. Magnetically driven outflows emerge from the disk. Figure 10 shows a result of the global 3D MHD simulation of black hole accretion disks starting from a disk with poloidal magnetic loops imbedded in the

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Fig. 10 A result of the global 3D MHD simulation of a black hole accretion disk starting from a disk with poloidal magnetic loops imbedded in the disk. Solid curves show magnetic field lines. Gray scale shows the Poynting flux

disk [7]. Color shows the Poynting flux. Large scale poloidal magnetic fields and outflows are automatically formed by the emergence of the magnetic loops and subsequent twisting and inflation of the loops.

5 Summary We have shown that even when the disk is not initially threaded by large scale poloidal magnetic fields, they can be formed by (a) twisting the magnetic loops connecting the central object and the disk, or by (b) twisting the magnetic loops buoyantly rising from the disk to the disk corona. Such inflating magnetic loops are collimated toward the rotation axis, and form magnetic towers. Magnetic reconnection taking place in the expanding magnetic loops inject hot plasmoids into the tower. Outside the magnetic towers, magnetically driven outflows emerge from the disk along the large-scale poloidal magnetic fields created by the expansion of the magnetic loops. The core jet and the outer wind structure will be formed in AGNs, galactic black hole candidates, and in protostars. Acknowledgements We thank Mami Machida, Yoshiaki Kato and Takuhito Kuwabara for collaboration and providing the figures. Numerical simulations were carried out on VPP5000 at Center for Computational Astrophysics, CfCA of NAOJ. This work is supported by the Grant-in-Aid for Scientific Research of the Ministry of Education, Culture, Sports, Science, and Technology (RM: 20340040).

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References 1. Balbus, S.A., Hawley, J.F.: A powerful local shear instability in weakly magnetized disks. I - Linear analysis, ApJ, 376, 214–233 (1991) 2. Blandford, R.D., Payne, D.G.: Hydromagnetic flows from accretion discs and the production of radio jets, MNRAS, 199, 883–903 (1982) 3. Casse, F., Keppens, R.: Magnetized Accretion-Ejection Structures: 2.5-dimensional Magnetohydrodynamic Simulations of Continuous Ideal Jet Launching from Resistive Accretion Disks, ApJ, 581, 988–1001 (2002) 4. Casse, F., Keppens, R.: Radiatively Inefficient Magnetohydrodynamic Accretion-Ejection Structures, ApJ, 601, 90–103 (2004) 5. Hayashi, M.R., Shibata, K., Matsumoto, R.: X-Ray Flares and Mass Outflows Driven by Magnetic Interaction between a Protostar and Its Surrounding Disk, ApJ, 468, L37–L40 (1996) 6. Kato, Y., Hayashi, M.R., Matsumoto, R.: Formation of Semirelativistic Jets from Magnetospheres of Accreting Neutron Stars: I njection of Hot Bubbles into a Magnetic Tower, ApJ, 600, 338–342 (2004a) 7. Kato, Y., Mineshige, S., Shibata, K.: Magnetohydrodynamic Accretion Flows: Formation of Magnetic Tower Jet and Subsequent Quasi-Steady State, ApJ, 605, 307–320 (2004b) 8. Kudoh, T., Shibata, K.: Magnetically Driven Jets from Accretion Disks. I. Steady Solutions and Application to Jets/Winds in Young Stellar Objects, ApJ, 474, 362–377 (1997) 9. Kudoh, T., Matsumoto, R., Shibata, K.: Magnetically Driven Jets from Accretion Disks. III. 2.5-dimensional Nonsteady Simulations for Thick Disk Case, ApJ, 508, 186–199 (1998) 10. Kuwabara, T., Shibata, K., Kudoh, T., Matsumoto, R.: Resistive Magnetohydrodynamics of Jet Formation and Magnetically Driven Accretion, PASJ, 52, 1109–1124 (2000) 11. Kuwabara, T., Shibata, K., Kudoh, T., Matsumoto, R.: The Acceleration Mechanism of Resistive Magnetohydrodynamic Jets Launched from Accretion Disks, ApJ, 621, 921–931 (2005) 12. Lovelace, R.V.E., Romanova, M.M., Bisnovatyi-Kogan, G.S.: Spin-up/spin-down of magnetized stars with accretion discs and outflows, MNRAS, 275, 244–254 (1995) 13. Machida, M., Matsumoto, R.: Global Three-dimensional Magnetohydrodynamic Simulations of Black Hole Accretion Disks: X-Ray Flares in the Plunging Region, ApJ, 585, 429–442 (2003) in Black Hole Accretion Flows, 14. Machida, M., Matsumoto, R.: Excitation of Low-Frequency QPOs in Black-Hole Accretion Flows, PASJ, 60, 613–626 (2008) 15. Matsumoto, R., Uchida, Y., Hirose, S., Shibata, K., Hayashi, M.R., Ferrari, A., Bodo, G., Norman, C.: Radio Jets and the Formation of Active Galaxies: Accretion Avalanches on the Torus by the Effect of a Large-Scale Magnetic Field, ApJ, 461, 115–126 (1996) 16. Rees, M.J., Begelman, M.C., Blandford, R.D., Phinney, E.S.: Ion-supported tori and the origin of radio jets, Nature, 295, 17–21 (1982) 17. Shibata, K., Uchida, Y.: A magnetodynamic mechanism for the formation of astrophysical jets. II - Dynamical processes in the accretion of magnetized mass in rotation, PASJ, 38, 631–660 (1986) 18. Uchida, Y., Shibata, K.: Magnetodynamical acceleration of CO and optical bipolar flows from the region of star formation, PASJ, 37, 515–535 (1985) 19. Ustyugova, G.V., Koldoba, A.V., Romanova, M.M., Chechetkin, V.M., Lovelace, R.V.E.: Magnetocentrifugally Driven Winds: Comparison of MHD Simulations with Theory, ApJ, 516, 221–235 (1999) 20. Ustyugova, G.V., Lovelace, R.V.E., Romanova, M.M., Li, H., Colgate, S.A.: Poynting Jets from Accretion Disks: Magnetohydrodynamic Simulations, ApJ, 541, L21–L24 (2000)

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Formation of Episodic Magnetically Driven Radiatively Cooled Plasma Jets in Laboratory Experiments Francisco Suzuki-Vidal, Sergey V. Lebedev, Andrea Ciardi, Simon N. Bland, Jeremy P. Chittenden, Gareth N. Hall, Adam Harvey-Thompson, Alberto Marocchino, Cheng Ning, Chantal Stehle, Adam Frank, Eric G. Blackman, Simon C. Bott, and Tom Ray

Abstract We report on experiments in which magnetically driven radiatively cooled plasma jets were produced by a 1 MA, 250 ns current pulse on the MAGPIE pulsed power facility. The jets were driven by the pressure of a toroidal magnetic field in a “magnetic tower” jet configuration. This scenario is characterized by the formation of a magnetically collimated plasma jet on the axis of a magnetic cavity, confined by the ambient medium. The use of a radial metallic foil instead

F. Suzuki-Vidal (), S.V. Lebedev, S.N. Bland, J.P. Chittenden, G.N. Hall, A. Harvey-Thompson, A. Marocchino, and C. Ning Imperial College London, The Blackett Laboratory, Plasma Physics Group, London SW7 2BW, United Kingdom e-mail: [email protected] A. Ciardi and C. Stehle Observatoire de Paris, LUTH, Meudon, 92195, France A. Frank and E.G. Blackman Department of Physics and Astronomy, University of Rochester, Rochester, NY, USA S.C. Bott Center for Energy Research, University of California, San Diego, 92093-0417, USA T. Ray Dublin Institute for Advanced Studies, Dublin, Ireland

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of the radial wire arrays employed in our previous work allows for the generation of episodic magnetic tower outflows which emerge periodically on timescales of 30 ns. The subsequent magnetic bubbles propagate with velocities reaching 300 km/s and interact with previous eruptions. This setup also allowed for the addition of a neutral gas above the foil in order to study the effect of the ambient density on the dynamics of both the early time hydrodynamic jet formed from plasma ablated from the foil and of the subsequent magnetic tower outflows.

1 Introduction Highly collimated jets and outflows from protostars have been of increasing interest to observational and theoretical astrophysics as observational techniques and computer simulations continue to improve. Recently high-energy density plasma experiments have been able to reproduce the dynamics shown on the stellar scale within the laboratory, maintaining the relevant dimensionless parameters (i.e. Mach number, Reynolds number, cooling parameter, etc.) given by the MHD scaling laws [16]. Experimental facilities have shown to be capable of performing experiments which reproduce particular features of these objects by using high power lasers [6, 7, 3] and high currents in a z-pinch machine [10, 11]. Different models of the formation of protostellar jets have proposed that magnetic fields are responsible for driving and collimating outflows from a system composed by a star with an accretion disk [4]. In particular in the “magnetic tower” model [13], the magnetic field topology evolves, due to differential rotation, into one with a predominantly toroidal magnetic field, which collimates ejected material from the system as a jet on the axis of a cavity confined by the external ambient pressure. This model has been proposed as a mechanism of jet formation for different astrophysical objects ranging from protostars to neutron stars [14, 8, 18, 9]. The experimental approach that reproduces some aspects of the plasma jet dynamics relevant to this model has been the radial wire array z-pinch configuration [12], in which a plasma jet is collimated on the axis of a magnetic “bubble” by rising toroidal magnetic field loops. Magneto-hydrodynamic simulations have shown that dimensionless parameters in these experiments are relevant to jets from young stellar objects [5]. In this paper we present experiments in which episodic formation of magnetic tower jets was observed. The experimental set-up also allows us to vary controllably the density of the ambient medium through which the magnetic tower jets propagate. It is believed that the knots and shocks observed in the protostellar jets could originate from both the variability of the outflow at the jet formation stage, or could arise from the interaction with the ambient medium. The experimental capabilities developed in the present work can contribute to a better understanding of the issues related to variability of astrophysical jets.

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2 Experimental Setup The experimental configuration is similar to the radial wire array z-pinch used in our previous experiments [12]. In the present experiments the current from the MAGPIE generator (peak current of 1 MA in 250 ns) [15] is driven into a 6–6.5 m thick aluminum foil, which is held radially between two concentric electrodes (Fig. 1). The central electrode (cathode) is a hollow cylinder with a diameter of 3.2 mm, with the diameter of the outer electrode being 60 mm. Diagnostics included: laser probing ( D532 nm, t0.4 ns) providing 2-frame interferometry, shadowgraphy and schlieren imaging; time resolved (2 ns exposure) pinhole cameras which recorded emission in the XUV region (>30 eV) providing up to 8 frames per experiment; magnetic “pick-up” probes to measure any trapped magnetic field inside the outflows; an inductive probe connected to the cathode to measure voltage and thus Poynting flux driving the outflow. The imposed current path (Fig. 1a) produces a toroidal magnetic field B below the foil which is directly proportional to the current and decreases with the radial distance from the cathode (B / I.t /=r). For peak current the toroidal magnetic field can reach magnitudes of B 100 T (1 MG) at the cathode radius. As the current increases in time the foil is ohmically heated by the current leading to the formation of plasma on the foil surface. The ablated plasma flows from the foil surface in the axial (JR  B ) direction. Side-on laser probing imaging (Fig. 2), taken at 172 ns from the current start, shows an axial displacement of the foil near the central electrode, which at early time is reasonably well described by 0-D equations of motion, assuming that most of the foil mass is accelerated by the pressure of the toroidal magnetic field. The material ablated from the foil fills the region above it with a low-density background plasma with a typical electron density integrated along the laser line of sight measured by laser interferometry of Ne  1018 cm2 . Axial density profiles reconstructed from interferometry show an exponential decay with the height above the foil, suggesting an isothermal expansion of the plasma with a typical sound speed of CS 9 km/s.

Fig. 1 Schematic of the experimental setup showing: (a) the current path through the foil (JR ), the toroidal magnetic field (B ) and the net J B force acting on the plasma produced from the foil. In (b) it is shown the current path after the formation of the magnetic cavity

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Fig. 2 Side-on laser interferometry image at 172 ns from the current start. The pressure of the toroidal magnetic field below the foil starts displacing plasma axially and a hydrodynamic jet is formed on axis from plasma ablated from the foil

Figure 2 also shows an enhancement of density in the ablated plasma in the region near the axis. This hydrodynamic “precursor” jet is formed from plasma ablated from the surface of the foil, which is redirected towards the axis by radial pressure gradients. The formation of such jets by converging plasma flows has been also reported in previous laboratory astrophysics experiments using conical wire arrays and radial wire arrays [2, 10, 11]. When an ambient gas was injected into the region above the foil in the present experiments, the presence of this hydrodynamic jet affected the early time dynamics of the interaction of the outflow with the ambient, as will be discussed later in this paper.

3 Formation of Episodic Magnetic Tower Jets The formation of magnetically driven jets starts later in time, when the Lorentz JR  B force (which is strongest at the cathode radius) leads to ablation of all of the foil mass near the cathode and to the formation of a small radial gap between the cathode and the remainder of the foil. From this moment the Poynting flux can be injected through this gap into the region above the foil. The toroidal magnetic field pushes the ablated plasma axially and radially outwards and also pinches the plasma on axis, forming a magnetic tower jet configuration. At this stage the current flows along the jet on the axis of the magnetic cavity and along the walls of the cavity, in the same way as in our previous experiments [12, 5]. The magnetic pressure from these rising toroidal loops inside the cavity inflates it both radially and axially, with measured velocities of VR 50–60 km/s and VZ 130–200 km/s respectively. Experimental results showing such dynamics are shown in Fig. 3. It can be seen that the initial diameter of the bubble is given by the diameter of the cathode. The most prominent feature of this new experimental set-up is that we now observe several subsequent outflows formed in the same experiment. It is possible to follow the axial positions of the subsequent episodes of the outflows shown

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Fig. 3 Side-on XUV emission images showing the formation of episodic magnetic towers outflows which emerge with a periodicity of 30 ns. The sequence shows results from a single experiment with the exception of the image at 233 ns

Fig. 4 Measurements of the axial extent of the episodic magnetic bubbles for the sequence shown in Fig. 3. The start of each outflow episode is correlated with soft x-ray emission produced by the pinching of a plasma jet on axis

in Fig. 3, with Fig. 4 presenting the measurements that allowed the determination of their axial velocities. It is seen that the tip of each outflow episode is expanding with approximately constant velocity, and the extrapolation of the trajectories back in time allows determining the starting time for each episode. Each subsequent bubble expands with a faster velocity, reaching VZ D 325 km/s for the third observed magnetic cavity. This increase in velocity is consistent with sweeping of the ambient plasma by the earlier episodes, thus allowing the subsequent magnetic bubbles to propagate through a lower ambient density. Figure 4 also shows that the episodic outflows are accompanied by episodic outbursts of soft x-rays (photon energy between 200–300 eV and above 800 eV), which can be well correlated with the formation of each new magnetic tower jet. This is an indication that each new episode starts from the pinching of plasma on the axis of the magnetic cavity and that pinched plasma is the source of the x-ray emission. Both the axial expansion dynamics and the periodicity of x-ray emission show a timescale of 30 ns for the formation of subsequent magnetic tower outflows.

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The formation of episodic magnetic tower outflows occurs due to reconnection of current at the base of the cavity, as the gap formed between the foil and the central electrode is closed by a plasma. The mass distribution in the radial foil setup is different from that in radial wire arrays used in our previous experiments (dm/dr / r for a foil, whereas constant for radial wires). This could lead to a smaller gap formed in the foil at the start of the first outflow and to a faster closure of the gap by the plasma expanding from the cathode and the remaining foil. The typical width of this gap estimated from the position at the base of the magnetic cavity walls is r  0.3–0.7 mm. The JR  B force acting on the plasma closing the gap will push this plasma upwards and will lead to formation of a new magnetic tower outflow. The process of gap closure and formation of new outflows will continue for the duration of the current pulse from the generator, allowing to obtain 4–5 episodes for each single experiment.

4 Jet Propagation in an Ambient Gas The radial foil setup readily allows for the addition of an ambient neutral gas, which is injected via a supersonic gas nozzle into the space above the foil before the start of the current pulse, in a similar way to previous experiments with conical wire arrays [1]. Argon was used in most of the experiments although other gases (e.g. Xe, He) were also tested. In the experiments reported here an estimated initial number density of the argon gas was N1017 –1018 cm3 . The presence of the ambient gas above the foil led to several new features. At early time the hydrodynamic jet formed by the plasma ablated from the foil interacts with the ambient producing a conical shock and a working surface. These dynamics can be observed in a sequence of XUV emission images shown in Fig. 5, where the axial displacement of plasma from the foil forms a conical shock, with the subsequent formation of the precursor hydrodynamical jet on the axis of the foil. The tip of the jet can be observed as a highly emitting region due to the compression of plasma in this region. It should be pointed out that for the sequence shown in Fig. 5 the diameter of the cathode was increased from 3.1 mm to 4.7 mm in order to decrease the magnitude of the toroidal magnetic field B below the foil, delaying in this way the formation of magnetic-tower outflows that ultimately overtake these

Fig. 5 Sequence of XUV emission images showing the formation of a conical shock that drives a hydrodynamical jet on the axis of the foil when an argon ambient is present (6.5 m thick aluminium foil, 4.7 mm diameter cathode). The cathode and the initial foil position prior to the start of the current are shown schematically

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Fig. 6 Images showing the dynamics of the interaction of jets with an argon ambient gas (6.5 m aluminium foil, 3.1 mm diameter cathode). An initial conical shock is formed by the hydrodynamical jet and the plasma ablated from the foil. A faster propagating bow shock structure develops ahead of the hydrodynamical jet (a-b). Magnetically driven jets emerging later in time overtake these earlier formed structures (c-d). The evolution of a magnetic tower jet embedded in the remnants of earlier episodes is also seen in images c-d

initial features. The formation of jets driven by the hydrodynamical flow of plasma has been previously observed and studied in other types of wire array configurations [10, 2]. The subsequent evolution of the interaction of the precursor jet with the ambient is shown in Fig. 6. The conical shock has an opening angle (measured from the jet axis to the conical shock boundary) of 60ı and is moving with an axial velocity of VZ 55 km/s (Fig. 6a). It is observed that ahead of the tip of the hydrodynamical jet a second shock feature is formed at 230 ns. This bow shock, best seen in XUV images, is a spherical front moving at a faster velocity of VZ 80–110 km/s (Figs. 6a–6b). The interaction of this bow shock with the initial conical shock forms a contact boundary (“Mach stem”) which can be seen in Fig. 6b as a horizontal dark (emitting) line. It is possible that radiation from the working surface at the end of the hydrodynamic precursor jet is playing a role in the formation of this fast bowshock like structure, but a more detailed investigation is needed here. The addition of an ambient gas shows no significant effect on the periodicity of magnetic tower jet formation. Figures 6c–6d show the evolution of an embedded magnetic tower jet formed inside the remnants from the earlier outflow episodes. For this embedded

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Fig. 7 Axial displacements and respective velocities for the different features observed with an argon background (as shown in Fig. 6). The dashed lines represent the new shock features produced by the addition of the ambient gas. The solid lines show the formation of magnetic-tower outflows, which develop in the same way as in the case without an ambient gas (Figs. 3 and 4)

magnetic tower both the envelope and the central jet are clearly seen. As the subsequent magnetic bubbles reach axial expansion velocities of 300 km/s they can catch-up and interact with the initial shock features produced by the earlier episodes. This interaction could be responsible for the complex shock structure seen at the top of the cavity in Fig. 6d. The velocities for the different features observed in this configuration were obtained by measuring their axial displacements in respect to the initial foil position at different times. These measurements are presented in Fig. 7, showing that all axial displacements can be well fitted by a linear approximation and thus are consistent with constant expansion velocities. The fits can be extrapolated to infer the starting time of the features. This results in a starting time of 75 ns for the conical shock, which is in good agreement with the time of melting of the foil by the current at the cathode radius. Figure 7 also shows outbursts of x-ray emission, which are correlated with the formation of magnetic-tower jets as in the case without an ambient gas. A deeper analysis of all these features will be presented in future publications.

5 Summary In these experiments we were able to show for the first time a way of producing episodic magnetically driven jets in the laboratory and to observe how they interact with a modified ambient medium. The similarities of the dynamics of these episodic jets with the single-episode magnetic tower jets studied in our previous experiments with radial wire arrays indicate that, although some of the plasma parameters in the

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radial foil experiments are yet to be accurately measured, the dimensionless numbers (M& 5, Re 1, cooling parameter  . 1) are expected to be similar and therefore these new experiments are relevant to the physics of young stellar objectsjet launching. Preliminary measurements of trapped toroidal magnetic field inside the magnetic towers (B 0.5 T) are consistent with a sufficiently high magnetic Reynolds number (ReM 100), which is in agreement with a temperature of the magnetically driven jet of 200 eV, estimated using x-ray spectroscopy. Voltage across the gap of the magnetic towers measured with an inductive probe has given first estimates of Poynting flux and magnetic energy available to drive these outflows. The process of current reconnection at the base of the cavity is responsible for the formation of the episodic jets in the experiments and this can be compared to the phenomenon of mass accretion and magnetic reconnection that could be responsible for episodic outflows in young stellar objects [8]. The effect of an ambient gas on the dynamics of the outflows produced in the radial foil configuration have also been studied. Although the overall dynamics are not affected by the presence of the ambient, i.e. the formation of a hydrodynamical jet and subsequent magnetic-tower outflows, these preliminary results show the formation of several additional features relevant to the propagation of astrophysical jets and their interaction with the ambient medium. Future publications will provide a further analysis on the formation of these features, together with possible connections to astrophysical scenarios. Acknowledgements This research was supported by the European Community’s Marie Curie Actions & Human Resource and Mobility within the JETSET network under contract MRTN-CT-2004 005592 and the Stewardship Sciences Academic Alliances program of the NNSA under DOE Cooperative Agreement DE-FC03-02NA00057.

References 1. Ampleford, D. J. et al.: Formation of Working Surfaces in Radiatively Cooled Laboratory Jets, Astrophys. Space Sci. 298, 241–246 (2005) 2. Ampleford, D. J. et al.: Supersonic Radiatively Cooled Rotating Flows and Jets in the Laboratory, PRL 100, 035001 (2008) 3. Blue, B. E. et al., Experimental Investigation of High-Mach-Number 3D Hydrodynamic Jets at the National Ignition Facility, PRL 94, 9 (2005) 4. Blandford, R. D. and Payne, D. G.: Hydromagnetic flows from accretion discs and the production of radio jets, MNRAS 199, 883–903 (1982) 5. Ciardi, A. et al.: The evolution of magnetic tower jets in the laboratory, Physics of Plasmas 14, 056501 (2007) 6. Farley, D. R. et al.: Radiative Jet Experiments of Astrophysical Interest Using Intense Lasers, PRL 83, 1982–1985 (1999) 7. Foster, J. M. et al.: Supersonic jet and shock interactions, Physics of Plasmas 9, 2251–2263 (2002) 8. Goodson, A. P. et al.: Jets from Accreting Magnetic Young Stellar Objects. I. Comparison of Observations and High-Resolution Simulation Results, APJ 524, 142–158 (1999) 9. Kato, Y. et al.: Formation of Semirelativistic Jets from Magnetospheres of Accreting Neutron Stars: Injection of Hot Bubbles into a Magnetic Tower, APJ 600, 338–342 (2004) 10. Lebedev, S. V. et al.: Laboratory Astrophysics and Collimated Stellar Outflows: The Production of Radiatively Cooled Hypersonic Plasma Jets, APJ 564, 113–119 (2002)

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11. Lebedev, S. V. et al.: Production of radiatively cooled hypersonic plasma jets and links to astrophysical jets, Plasma Physics and Controlled Fusion 47 (2005) 12. Lebedev, S. V. et al.: Magnetic tower outflows from a radial wire array Z-pinch, MNRAS 361, 97–108 (2005) 13. Lynden-Bell, D.: Magnetic collimation by accretion discs of quasars and stars, MNRAS 279, 389–401 (1996) 14. Lynden-Bell, D.: Magnetic jets from swirling discs, MNRAS 369, 1167–1188 (2006) 15. Mitchell, I. H. et al.: A high impedance mega-ampere generator for fiber z-pinch experiments, Rev. Sci. Instr. 67, 1533–1541 (1996) 16. Ryutov, D. D. et al.: Criteria for Scaled Laboratory Simulations of Astrophysical MHD Phenomena, APJ-S 127, 465–468 (2000) 17. Suzuki-Vidal, F. et al.: Formation of Magnetically Driven Radiatively Cooled Plasma Jets in the Laboratory, HEDLA 08 Proceedings (in press) (2008) 18. Uzdensky, D. A. and MacFadyen, A. I.: Stellar Explosions by Magnetic Towers, APJ 647, 1192–1212 (2006)

Part IV

Observational Constraints on Jet Launching

Jets in the MHD Context Nektarios Vlahakis

Abstract Outflows in the form of jets is a widespread phenomenon in astrophysics. Their main driving mechanism is likely related to magnetic fields. These fields are able to tap the rotational energy of the central object and its surrounding disk, and accelerate and collimate matter ejecta. To zeroth order these outflows can be described within the theory of steady, axisymmetric, ideal magnetohydrodynamics (MHD). The analytical insight into the equations of the theory (mostly on the transfield component of the momentum equation) gives simple analytical scalings for the flow speed, density, and magnetic field. The analysis is focused on nonrelativistic YSO jets; similar works [1, 2] exist for relativistic AGN, and highly relativistic GRB jets.

N. Vlahakis () IASA and Section of Astrophysics, Astronomy and Mechanics, Department of Physics, University of Athens, Panepistimiopolis, 15784 Zografos, Athens, Greece e-mail: [email protected]

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1 The Steady, Axisymmetric, Ideal MHD Description The system of equations of nonrelativistic, steady, ideal MHD, consist of the solenoidal and flux-freezing conditions  r B D 0, r  .V  B/ D 0, the continuity r . V / D 0, energy V r P =

D .  1/q=

, and momentum

.V r/ V D rP  r.GM=r/ C .r  B/  B=.4/ equations. Here V is the velocity of the outflow, B the magnetic field, the mass density, P the pressure, GM=r the gravitational potential of a central star with mass M,  is the specific heat ratio, and q is the volumetric rate of energy input/output. In this study q is assumed zero, since we are mainly interested in flows at relatively large distances from the central source where the energy input/output is not expected to affect the dynamics. Assuming axisymmetry (@=@ D 0, in cylindrical Œz ; R$ ;  coordinates), five conserved quantities along the flow exist. If A D .1=2/ B d S is the poloidal magnetic flux function, these are (e.g., [3]) the field angular velocity ˝.A/, the mass-to-magnetic flux ratio A .A/, the total angular momentum-to-mass flux ratio L.A/, the total energy-to-mass flux ratio E.A/, and the adiabat Q.A/. Using the expressions of these constants of motion we may express all the flow quantities as functions of the poloidal Alfv´en Mach number M and the magnetic flux A:1 $ rA  O $ 2 ˝A G 2  1 L ; B D  A ; G ; $A2  ; 2 $ $ M 1 $A ˝ 2 2 2 2 2 O A M rA   $ ˝ M G ; D Vp D ; V D A ; P D Q

: $A $ M2  1 4M 2 Bp D

(1a) (1b)

The two equations that give the remaining unknowns M and A are the Bernoulli and the transfield component of the momentum equation.

1.1 Bernoulli Equation The Bernoulli equation is the expression of the total energy-to-mass flux ratio E.A/ D

Vp2 2

C

V2 2

C

 P GM $˝.B /  C :  1

r A

(2)

The terms on the right hand side correspond to the poloidal and azimuthal kinetic, enthalpy, gravity, and Poynting. Using (1) it takes the form ED

1

G2  1  P GM M 4 jrAj2 $A2 ˝ 2 .M 2  G 2 /2 C C  C$A2 ˝ 2 2 : (3) 2 2 2 2 2 2G .M  1/  1

r M 1 2A $

The subscripts p/ denote poloidal/azimuthal components.

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In the super-Alfv´enic part of the flow (G  1, M  1), the dominant terms of the right hand side are the poloidal kinetic and the Poynting. As a result the Bernoulli equation (3) can be simplified to E

Vp2 2

C

2 M 4 jrAj2 $˝.B / 2 2 G ; or, E  C $ ˝ : A 2 A M2 2A $ 2

(4)

Eliminating M using the first from (1b) we can write the following equation for the kinetic-to-total energy flux ratio   1=2 .1  /  m

Vp2 =2 $jrAj A˝ 2 : ;  ; m  1=2 3=2 A E 2 E A

(5)

Evidently, the acceleration efficiency depends on the shape of the flow A.$ ; z/ through the function $ jrAj=A .D $ 2 Bp =A/, the solution of the transfield component of the momentum equation. At the fast magnetosonic point (subscript f) the value of f is 1=3, see Fig. 1. Equivalently one third of the total energy flux has been transfered to kinetic up to this point. We also find .$ jrAj=A/f D 2=.33=2 m / (see Fig. 1). The acceleration efficiency 1 (the asymptotic value of ) depends on the decline of the function $jrAj=A D $ 2 Bp =A in the super–fast magnetosonic regime. Assuming r  1=$ we deduce the asymptotic value .$ jrAj=A/1  1, and thus the acceleration 1=2 efficiency depends on the value of the constant of motion m through 1 .11 /  3=2 then the fast magnetosonic point is located at infinity and m . If m D 2=3 1  1=3. If m < 2=33=2 then 1 > 1=3. For example, if m D 0:2, then .$ jrAj=A/f  1:9 and 1  0:77, while if m D 0:1, .$ jrAj=A/f  3:8 and 1  0:89. Very high acceleration efficiencies correspond to tiny values of m and large values of .$ jrAj=A/f , meaning that the poloidal magnetic field/streamlines

Fig. 1 The function  1=2 .1  / (solid line). Its intersection with the line m $jrAj=A (the value of the latter at each point is known from the solution of the transfield) gives the value of . There are two solutions; the larger corresponds to the super–fast magnetosonic regime. At the fast magnetosonic point there is one double solution f D 1=3, m .$ jrAj=A/f D 2=33=2

ζ1/2(1–ζ) 2/33/2

σmϖ |∇ A| /A

0

ζ1

1/3

ζ2

1

ζ

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N. Vlahakis

are bunched at the distance of the fast magnetosonic point and jrAj  A=$. For example, 99% acceleration efficiency (1 D 0:99) requires m  0:01 and .$ jrAj=A/f  38.

1.2 Transfield Equation The transfield component of the momentum equation can be written as Bp2



Bp2 4

where

M V2 Vp2 2

1C

!

4R

.M 2  1/ D

B 2 rA r .$B/

4 jrAj $B

A2 GM rA rr rA r$ rA

C

rP C ; (6) jrAj $ jrAj 4M 2 r 2 jrAj r 2 A  rA r ln j$rAj 1  R jrAj

is the curvature of the poloidal field/streamlines. In the super–fast magnetosonic regime (M  1), by dropping the pressure and gravity terms, and assuming that the centrifugal term (second on the right hand side) is also negligible2 , (6) can be simplified to $ M 2 Bp2 rA $ 2  $r ln .$B/

D ı; , or, Mfm 2 R B jrAj R

(7)

where ı is a slowly varying function of distance.3 Here the relations for the Alfv´en and fast magnetosonic Mach numbers were used, namely M 2 D 4 Vp2 =Bp2 and 2 Mfm  4 Vp2 =B 2 (assuming that the flow is cold at super–fast magnetosonic distances). Equation (7) provides a connection between the geometry of the flow and the functional dependence of the kinetic-to-total energy flux ratio . Indeed, the square of the fast Mach number can be written as

For a power-law shape z / $ b the term $=R / $ 2 =z2 , while M 2 / $ 2 . Thus, the ratio between the poloidal curvature term (left hand side of 6) and the centrifugal term scales as / $ 4 =z2 / $ 2.2b/ , and increases with distance for 1 < b < 2. (The ratio can be estimated  .6ı=5/.$f =$A /2 .$=$f /2.2b/ .) Thus, only if the flow shape is nearly radial (b  1), or, b > 2 the centrifugal term must be kept. 3 It is a constant of order unity if the gradient in the transfield direction (along the vector rA=jrAj) scales as 1=$, i.e., if .rA=jrAj/ r  1=$. 2

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209 2 Mfm 

4 Vp2 B2

D

Bp2 M 2 B2



2 ; 1

since the ratio jB j=Bp  $˝=Vp and $A2 ˝ 2 G 2 =M 2  E.1  / (using (1), (4), and (5)). Thus, (7) gives 

1 : 1 C .2=ı/.$=R/

(8)

At the fast magnetosonic point where f D 1=3 we get .$=R/f  ı. For a power-law field/streamline shape in cylindrical coordinates z D zf .$=$f /b we get $=R D Œ.b  1/=b 2 .$f =zf /2 .z=zf /2.b1/=b . The requirement that .$=R/f  ı fixes the coefficient Œ.b  1/=b 2 .$f =zf /2  ı. (The latter relation connects the shape of the fast magnetosonic surface zf D zf .$f / with the slope b.) Substituting the curvature in (8) we get 

1 : 1 C 2.z=zf /2.b1/=b

(9)

(Note that the previous equation holds as long as the function $jrAj=A declines, see (5). The maximum value 1 is determined from the minimum value of $jrAj=A.) Examples on how the kinetic-to-total energy flux ratio increases with distance are given in Fig. 2. The analytical expression (9) is in very good agreement with the exact r selfsimilar solution of [4], with error less than 10% in the super–fast magnetosonic regime, as seen in the right panel of Fig. 3 (the line shape near the fast magnetosonic point can be approximated as z=zf D .$=$f /b with b D 1:29, see the left panel of Fig. 3). The resulting analytical solution for .z/ (9) is based on the simple relation 2 D ı (see 7), which can be rewritten as tan #=tan #f D sin M , using .$=R/Mfm $=R D .b  1/.d$=d z/2 D .b  1/ tan2 # (valid for a line shape z / $ b with

1 b=2.5 b=2 b=1.3 b=1.1

0.9 0.8 0.7 ζ

Fig. 2 The kinetic-to-total energy flux ratio  as a function of z=zf , for various values of the exponent b that controls the flow shape z D zf .$=$f /b . For larger values of b the collimation is more efficient, resulting in faster acceleration. Note that the value of b is related to the ratio $f =zf D ı 1=2 b= .b  1/1=2 D .3:5ı 1=2 ; 2:4ı 1=2 ; 2ı 1=2 ; 2:0ı 1=2 / for b D .1:1 ; 1:3 ; 2 ; 2:5/, respectively

0.6 0.5 0.4 0.3

1

2

3

4

5

6 z/zf

7

8

9

10

210

N. Vlahakis 1 7 6

4

ζ

ϖ/zf

5 0.1

3 analytical (b=1.29) r self-similar relative error

2 analytical (b=1.29) r self-similar

1 0

0

1

2

3

4

5 z/zf

6

7

8

9

10

0.01 0.01

0.1

1

10

100

z/zf

Fig. 3 Comparison of the analytical results with the exact r self-similar solution of [4]. Poloidal field/streamlines (left panel) and the kinetic-to-total energy flux ratio (right panel). The analytical expression for the line shape is z=zf D .$=$f /b and for  is given by (9)

opening half-angle #), and M  arcsin.1=Mfm / for the opening half-angle of the fast magnetosonic Mach cone (for a cold flow). The lasr equation expresses the fact that the flow should be in causal connection in the transfield direction in order for the collimation to continue to be efficient, resulting in the continuation of the acceleration as well.

2 The Analytical Solution By using (1), (4) and the main analytical results (5) and (9), we can write simple analytical formulas for all physical quantities of the flow in the super–fast magentosonic regime (assuming that the flow is cold): 21=2 E 3=2 A  1=2 .1  / EA .1  / ; ; B   $ 2˝ 2 $˝  2 E.1  / $ 2 ˝ 2  E C E Vp  21=2  1=2 E 1=2 ; V  A ;  A 2 2 ; $˝ 4$ ˝  b 1 $  ; z  z : f 2.b1/ $f 1 C 2.$=$f / Bp 

(10a) (10b) (10c)

The above equations can be further simplified by using E  $A2 ˝ 2 , a relation that holds if the flow near the source is Poynting flux-dominated (in that case the dominant term of the right hand side of 3 is the last one, which can be approximated as $A2 ˝ 2 since G 1, M 1 near the source):

Jets in the MHD Context

21=2 $A3 ˝A  1=2 .1  / $A2 ˝A .1  / ; ; B    $2 $  2 $ 2 .1  / $A2 ˝ Vp  21=2  1=2 $A ˝ ; V  ;  A A 2 ; $ 4$  b 1 $  ; z  zf : 2.b1/ $f 1 C 2.$=$f /

Bp 

211

(11a) (11b) (11c)

Acknowledgements The present work was partially supported by the European Community’s Marie Curie Actions - Human Resource and Mobility within the JETSET (Jet Simulations, Experiments and Theory) network under contract MRTN-CT-2004 005592.

References 1. S. S. Komissarov, M. V. Barkov, N. Vlahakis, A. K¨onigl, MNRAS, 380, 51 (2007) 2. S. S. Komissarov, N. Vlahakis, A. K¨onigl, M. V. Barkov, MNRAS, in press (2009) 3. K. Tsinganos, ApJ, 252, 775 (1982) 4. N. Vlahakis, K. Tsinganos, C. Sauty, E. Trussoni, MNRAS, 318, 417 (2000)

Jets from Embedded Protostars Brunella Nisini

Abstract An enormous observational progress has been made in the last decade on the characterization of collimated jets from young and embedded protostars, the so-called class 0 and class I objects. Here I will review the main results in this field recently obtained thanks to the improved IR and sub-mm performances of the present instrumentation, and to the development of the relevant spectroscopic diagnostic techniques.

1 Introduction According to the models [27, 3], accretion and ejection are intimately related through the presence of a magnetized accretion disk: the jets carry away the excess angular momentum, so that part of the disk material can move towards the star. This paradigm of star formation is now being observationally tested in low mass,

B. Nisini () INAF-Osservatorio Astronomico di Roma e-mail: [email protected]

K. Tsinganos et al. (eds.), Protostellar Jets in Context, Astrophysics and Space Science Proceedings, DOI 10.1007/978-3-642-00576-3 26, c Springer-Verlag Berlin Heidelberg 2009 

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still active, pre-main sequence stars (the Classical T Tauri stars, CTTs), through detailed UV, optical and near-IR observations [40]. However, the characteristics of jets from these evolved YSOs are unlikely to be appropriate for those from protostars in earlier evolutionary phases, which are expected to propagate in a denser medium and be associated with more energetic mass ejection phenomena. In such un-evolved objects, so-called class I and class 0 sources, the initial part of the jet is often obscured to optical wavelengths. The availability of new sensitive instrumentation with high spatial resolution working in the IR and sub-mm domain, is now allowing jets from embedded protostars to be studied in details. In parallel with the technological progress, spectral diagnostic tools employing IR/sub-mm lines have been developed, making it possible to perform a quantitative analysis on the physical properties of the embedded jets. The scope of the present contribution is to describe some of the progresses done in recent years in our understanding of the class 0/I jets physical characteristics and to discuss the main differences/similarities with CTT jets.

1.1 Diagnostic on Physical Properties of Embedded Jets Diagnostics based on IR spectroscopic techniques has been widely used in the last decade to probe the kinematics and excitation of the jet beams in embedded protostars. These studies have made it possible a preliminary comparative analysis on how the jet physical properties of young protostars differ from those of CTT stars to get insight of any evolutionary effect in jets excitation and dynamics. The diagnostic techniques mainly employ the wealth of lines produced by [FeII] and H2 in the 0.8– 2.5 m spectral range, that allow both the atomic and molecular components of the shocked gas to be traced simultaneously (see e.g. [21,31]). [FeII] lines, in particular, trace gas at low ionization with temperatures in the range 8000–20 000 K and densities 103 –105 cm3 , thus similar to the physical conditions probed by optical lines in CTT jets [32]. Ratios of [FeII] lines in the H band, such as 1.533 m/1.644 m, or 1.60 m/1.644 m are used to measure the electron density, while the combination of [FeII] lines in the J and H bands can be used to measure the reddening, an important parameter for any diagnostic use of ratios from lines with very different wavelengths or from absolute line luminosities. In the jet regions where the extinction is less extreme, like when the jet emerges from the dense circumstellar envelope, optical and IR lines can be used together to get additional information on the physical conditions. For example the [FeII]1.64 m/0.862 m ratio is a sensitive temperature indicator that does not rely on assumptions on abundance or ionization fraction like other ratios emplyed in optical studies [33, 37]. Although this kind of diagnostic has been applied so far only to a limited sample of class 0/I jets, some general results can be already drawn. In Table 1a summary is given of the physical parameters derived in a sample of sources that includes both class 0/I and CTT objects. All these sources, including the class0/I external part of the jet beams, have been observed also in the optical, and thus it has been possible to apply the diagnostic based on optical lines developed by [2] (the so-called BE techniques)

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Table 1 Physical properties of some jets from different classes of objects investigated in both the optical and IR range. Parameters averaged over the brightest jet knots BE technique [FeII] lines Jet HH 1 HH 111 HH 34 HH 46 HH 83 HH73 HH 24 C/E

class 0 0 I I II II II

ne 103 cm3 2.0 1.0 0.8 1.2 0.5 0.6 0.4

xe 0.05 0.1 0.04 0.2 0.4 0.3 0.3

Te 103 K 12.0 13.0 13.8 1.6 17.5 17.5 19.3

nH cm3 36.6 11.3 16.2 2.6 0.9 1.7 1.3

Ref [33] [38] [38] [2] [38] [38] [38]

ne cm3 5.7 2.6 1.8 3.6 ... ... ...

Te 103 K 9.8 7.5 5.8 ... ... ... ...

Ref [33] [38] [38, 17] [18] [38] [38] [38]

that allows one to derive not only the electron density and temperature, but also the ionization fraction and thus the total gas density (for details, see [15] and Fig. 1). All the jets have been simultaneously observed also in the IR, but only in the class 0/I objects the [FeII] lines have been detected and a complementary analysis have been performed. Several interesting results can be derived from this table. In class 0/I sources the electron density and temperature derived from the [FeII] lines are always higher and lower, respectively, than the correspondig values derived from the optical analysis. This is the consequence of the different location of the optical and IR line emitting regions in the not resolved cooling layer behind the shock front: as shown in [33], Fe emission traces post-shocked regions located at larger distances from the shock front than optical lines, where the compression is higher and the temperature is declining. In comparison with the sampled CTT jets, class0/I jets are cooler and have higher electron and total densities, but lower ionization fractions. This is what expected from the excitation due to shocks with similar velocities but different pre-shock densities. Finally, Table 1 also shows that the objects having an optical estimated electron density below 103 cm3 have not been detected in [FeII] lines, which indicates that iron is a good diagnostic tracer only of the densest jets. The detection of bright [FeII] lines in dense jets is a clear evidence that most of the dust grains where iron is normally locked in interstellar medium conditions, have been destroyed by passage of the shocks, releasing refractory elements into gas phase. The abundance of refractory species, such as Fe and Ca, has been recently measured in few class I jets by [33, 38] and [17]. The interesting result is that in the inner jet regions, the gas-phase abundance of these species is only 20–30% of the solar value, while it rises to more than 80% in the external regions. This is a strong indication that a significant fraction of dust grains is still present in the initial jet beam, where metals are locked. If confirmed by observations at higher angular resolution, this evidence may suggests that the dust grains are directly lifted from the protostellar disk and accelerated in the jet, a possibility which is compatible only with disk-winds models predicting that regions at large disk radii are involved in the jet acceleration. In few cases, e.g. HH34, HH46 and HH1 jets, where the iron gas phase abundance in the jet has been estimated, it has been possible to measure the mass flux of

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Fig. 1 Left: Combined optical and IR diagnostic analysis applied to the HH34 jet (from [38]). At the top of the figure the intensity profiles of different optical and IR lines are plotted as a function of the distance from the driving source. Open circles are the values of physical parameters derived from [Fe II] line ratios, while filled circles indicate the values determined from optical lines using the BE technique [2]. Right: Velocity resolved [FeII] lines diagnostic applied to the internal A6 knot of the HH34 jet, located within 2 00 from the central source (From [17]). Top panel shows the [Fe II]1.64 m/1.60 m intensity ratio as a function of velocity, while the profiles of the two lines are represented in the bottom panel. The plotted ratio is sensitive to density variations and indicates that the HVC gas (at VLSR  100 km s1 ) has an electron density lower than the gas in the LVC (at VLSR  30 km s1 )

the jets, from the 1.64 m line absolute luminosity [15]. The derived values are in the range 5 108 –5 107 Mˇ yr1 ; such values are similar or higher than the mass fluxes of the most active CTT stars. The mass accretion rates estimated in these stars from the Br luminosity and the source bolometric luminosity [1, 17], indicate that, as in CTT, the mass ejection over mass accretion rates are in the range 0.01–0.1, suggesting that the mechanism for the jet launching should remain the same over mass accretion rates spanning several order of magnitudes. In addition to the atomic component discussed so far, class 0/I jets very often present emission from molecular gas, mainly traced through the IR roto-vibrational H2 lines. The ratio between [FeII]/H2 emission often decreases with distance from the driving source (e.g. [32]) and H2 becomes important only in the bow shock regions of interaction between the jet and the ambient medium. In many cases, however, the initial jet beam can be traced only through H2 lines, due to either a larger

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extinction or very low excitation conditions, as evidenced by, e.g. [9, 19]. In these cases, the H2 IR line spectrum represents the only way to get some information on the jets excitation conditions. This is usually done through a Boltzman diagram where the column density distribution derived from the observed lines is compared with that expected from an LTE distribution of gas at a single temperature. Temperatures derived in this way ranges between 2,000 and 4,000 K, but often different temperature components are found in the same excitation diagram as traced by the different vibrational levels. Several classes of shock models, including steady-state C and J-type shocks or J-shocks with a magnetic precursor, have been employed to reproduce the different excitation diagrams observed. The general result is that very rarely a steady-state shock can reproduce the temperature stratification often observed in the H2 excitation diagrams. The time scales for a shock wave to attain steady-state depends on the density and for n  104 cm3 is of the order of 104 yrs, i.e. becomes comparable with the jet crossing time. It is indeed possible to see time-dependent effects in young flows and measure the age of the flow if the H2 excitation diagram is sufficiently sampled over a large energy range [29, 20]. Finally, mass flux rates can be also derived from H2 line luminosity, if the molecular jet velocity is known [31]. In class I jets presenting bright atomic line emission, the contribution of the molecular component to the total jet mass flux is negligible. However, for the molecular jets, mass fluxes comparable to those derived in atomic class I jets are found, indirectly indicating that the different excitation conditions exhibited by these two classes of jets is not related to a different accretion efficiency.

1.2 Observations of the Jet Base Many studies have been conducted in the last few years to investigate the structure and kinematics at the base of atomic jets in class I sources and compare them with the more evolved CTT sources. This have been usually done through long-slit spectroscopy with the slit positioned both along and across the jet, and constructing the corresponding Position-Velocity (PV) diagrams illustrating radial velocity variations both as a function of the distance from the central source and in the jet transversal direction. A detailed study of this kind have been done on the L1551-IRS5 jet by [39] with the IRCS instrument on Subaru. This is one of the first studies that recognized in a [Fe II]1.64 m PV diagram the presence of two velocity components, a narrow High Velocity Component (HVC) at V 300 km s1 and a Low Velocity Component (LVC) at VD 100 km s1 that is spatially confined within 400 AU from the IRS5 source. This velocity structure is typical of the Forbidden Emission Line (FEL) regions observed in optical studies of CTT stars (e.g. [25]) and indicates that the jets kinematical structure is already defined at early evolutionary stages. [37] constructed synthetic [Fe II] PV diagrams adopting the cold disk-wind model developed by [16] and compared them with the PV diagram observed in L1551.

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They found that the disk-wind model reproduces quite well the observed two velocity components but the LVC becomes weaker than the HVC too rapidly with respect to the observations, where the LVC dominates over the HVC further out, until distances of 0.6 00 from the star. Also the predicted electron density is lower by one/two order of magnitudes with respect to the values measured by [26] from [Fe II] line ratios. The same kind of problems have been found in the comparison of cold disk-wind models with optical observations of CTT stars (e.g. [28]). Additional heating at the jet base could improve the matching between the predicted and observed PV diagrams, enhancing the relative contribution of the LVC with respect to the HVC. When more than one velocity resolved [Fe II] line is observed, one can study variations of the physical properties as a function of the jet velocity, to separately investigate the parameters pertaining to the different jet components. This kind of analysis have been performed by [17, 18] on the HH34 and HH46 jets, using the [Fe II]1.64 m/1.60 m ratio to derive the electron density and mass flux rate in different velocity bins (see Fig. 1). Like for L1551 IRS5, also in these two sources the PV diagram at the jet base identifies a LV (at nearly 0 km s1 VLSR ) and a HV component (at  200 km s1 ) (see the contribution by Garcia Lopez et al. in this conference). In both the considered jets, the LVC observed close to driving source has an electron density which is a factor of two higher than the density in the HVC: most of the mass flux, however, is carried out by the HVC that is the only one surviving at larger distances. This set constraints on the origin of the LVC: an electron density higher than in the HVC seems in fact in contrast with scenarios in which the LVC is due to gas entrained by the HVC or to gas directly ejected in the external layers of a disk-wind (while the HVC should be ejected in the inner jet), since in both the cases one would expect both the ionization and the total density to decrease from the high to the low velocities. Measurements of physical parameters at the base of IR jets in class I sources have been also obtained by [42], through Subaru echelle spectroscopy. Electron densities of 104 cm3 or greater are always found in the inner jet regions, while the H2 line ratios indicate temperatures of 2  3  103 K with no evidence for temperature stratifications. An interesting result found by [8] and [9], is that several class I jets present high-velocity H2 emission within a few hundred of AU in several class I jets. Such small scale H2 emission regions have been called ’molecular hydrogen emissionline’ (MHEL) regions in analogy with the atomic FELs regions observed in TT stars. In fact, like the FELs, both LVC and HVC are observed in these regions, with velocities of the order of 5–20 and 50–150 km s1 respectively. The origin of the MHELs is unclear, but the similarity of their kinematical signature with the ones observed in atomic jets suggests a common origin for the two. Such a tight relationship between MHELs and FELs has been demonstrated by [10], who showed that a FEL region, traced by [Fe II] lines, is often associated with the H2 emission region: [Fe II] emission has usually velocities higher than H2 and a ratio between the HVC/LVC brightness higher than the H2 emission. The most likely interpretation is that the HV H2 emission is excited in a layer between the atomic jet and the

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near-stationary and dense ambient medium. The H2 LVC, however, could be also due to a cool molecular component of a disk-wind. Adaptive optics assisted observations were obtained on one of this MHEL regions, namely the SVS13 jet, using NACO at the VLT [11]. The reached spatial resolution of <0.25 00 , well sampled by a pixel scale of 0.027 00 , was able to resolve the H2 emission into several peaks along the jet, consistent with thermal expansion of packets of gas ejected during periods of increased accretion activity. Proper motion measurements, obtained combining these data with Subaru data taken two years earlier indicate that the H2 peaks possess a tangential velocity of 0.028 00 /year. At variance, [Fe II] presents a barely resolved single peak which seems to be stationary. This evidence suggests that [Fe II] could be associated with a stationary collimation shock at the base of the jet. Interesting results have been also obtained from studies of the kinematics and excitation of IR jets in the transverse direction across the width of the jet (see the contribution by Coffey et al. in this conference). Chrysostomou et al. [5] in particular, obtained high dispersion long-slit observations of the two small scale H2 jets from HH26-IR and HH72-IR with the slit perpendicular to the jet axis, and in both jets detected an asymmetry of the measured radial velocity with respect to the jet axis, that can be modelled by assuming that the jets possess a toroidal velocity due to rotation. These observations follow the detection of rotational signatures in jets obtained through HST observations of T Tauri micro-jets (e.g. [8]), showing that also jets from younger protostars are able to carry the excess of angular momentum away from the central source. Instrumentation performing IFU spectral images at adequate spatial and spectral resolution, such as SINFONI at VLT, are expected to give a more complete picture of both the excitation and kinematics along and across the jet structure [12].

1.3 Class 0 Molecular jets In heavily embedded Class 0 protostars, jets are often observed only at millimeter wavelengths in the form of a collimated, high-velocity (100–300 km s1 ) molecular outflow that that can be traced down to the central source. The best and most studied example is the HH211 jet [22], while others of such cases are L1448 [23], HH212 [6] and IRAS04166+2706 [41]: all of them have very short dynamical times of the order of 103 –104 yrs, indicating that the central source should be indeed very young. These highly collimated flows are usually enveloped by CO molecular gas at low velocity, delineating a wide angle cavity. This morphological evidence suggests that the molecular jet in its propagation entrains the ambient gas producing the slowly moving outflow. Near-IR H2 is not observed near the source and at the jet base, instead tracing the hot gas excited further downwind in bow shocks at or near the jet apex. In addition to CO, these jets are well traced by SiO emission. SiO has critical densities larger than CO, i.e. of the order of 106 cm3 or larger and thus is able to

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trace the densest jet components, more linked to the primary jet directly ejected from the protostars. Indeed, interferometric SiO observations, performed with both the PdB and SMA facilities [24,30,4], show that such molecular jets have characteristics very similar to the hot jets seen in T-Tauri or Class I YSOs, such as a knotty structure indicating different episodes of mass ejection, velocities similar to those of T Tauri micro-jets and a very narrow width. The jet width has been for the first time resolved in HH212 through PdB SiO observations [6, 4] and results of only 0.200 , e.g. 90 AU, at 500 AU from the protostar. Such a small width, very similar to the width of T Tauri atomic jets at the same distance, together with the high velocity close to the escape speed, suggest that the jet acceleration mechanism and launching zone remain the same in all phases, and that the jet collimation is already defined very early in the protostellar evolution. Invaluable information about the excitation conditions of the class 0 SiO jets comes from multi-line analysis at different wavelengths. Sub-mm SiO multitransition studies of HH211 and L1448 have indeed indicated that the gas must be very dense (n(H2) 105 –106 cm3 ) and relatively warm (T100–500 K) [35, 36]. The combination of PdB and SMA SiO observations in HH212 have shown that SiO in the ‘microjet’ is optically thick, indicating abundances larger than previously derived assuming optically thin emission [4]. Low spatial resolution ISO-LWS observations of L1448 suggest temperatures between 300 and 1500 K, inferred from the copious high-J CO and H2 O emission [34]. Recently Spitzer observations of the L1448 and HH211 SiO jets have confirmed that in such jets there is a component of warm gas, copiously emitting in the mid-IR pure rotational lines [13, 14] (see Fig. 2 and the contribution by Dionatos et al. in this conference). In addition, the Spitzer mid-IR observations have for the first time revealed the presence of an embedded atomic jet at low excitation, traced by the fundamental transitions of

O 1.5"

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Fig. 2 Spitzer H2 S(5) map of the HH211 molecular jet superimposed over the high velocity CO JD2–1 map of [22]. From [14]. A color version of this plot can be found in the electronic version of this book and in Appendix A (Fig. 2)

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[FeII], [SiII] and [SI]. Such an atomic gas is characterized by an electron density ne  300–500 cm3 , a temperature Te < 2,500 K and a ionization fraction <5 103 ; its physical conditions are thus very different from those found in more evolved class I and class II jets. Nevertheless, the momentum flux estimated in the atomic component amount to 106 , i.e. high enough to substain the molecular outflow in L1448.

2 Conclusions The progresses done in our understanding of young class 0/class I jets have been enormous, thanks both to the access at sensitive instrumentation mounted on large telescopes, and to the development of sophisticated models and tools for the data analysis and interpretation. In the very next future new IR instrumentation specifically designed to improve performances at high-angular resolution will be operative. AO assisted IFU spectrometers using laser guide stars are starting now to be operational and will be able to provide sub-arcsec resolution 2D maps of excitation and kinematics at the jet base, allowing a better comparison with model for the jet acceleration and heating to be performed. With a similar resolution, it will be possible in the next decade to obtain reconstructed interferometric images from planned facilities, like the VLTI/VSI instrument and Linc-Nirvana on LBT. Finally, the MIRI instrument on JWST will make it possible to observe in the 20 m window with  1 00 of angular resolution, while PACS on the Herschel satellite will have a spatial resolution of 9 00 at 60 m. These instruments will therefore open the mid- and far-IR spectral range to observations at the relevant spatial scales for the study of young stars flow.

References 1. Antoniucci, S., Nisini, B., Giannini, T., & Lorenzetti, D. 2008, A&A, 479, 503 2. Bacciotti, F., & Eisl¨offel, J. 1999, A&A, 342, 717 3. Casse, F., Ferreira, J., 2000, A&A, 353, 1115 4. Cabrit, S., Codella, C., Gueth, F., Nisini, B., Gusdorf, A., Dougados, C., & Bacciotti, F. 2007, A&A, 468, L29 5. Chrysostomou, A., Bacciotti, F., Nisini, B., Ray, T. P., Eisl¨offel, J., Davis, C. J., & Takami, M. 2005, Protostars and Planets V, 8156 6. Codella, C., Cabrit, S., Gueth, F., Cesaroni, R., Bacciotti, F., Lefloch, B., & McCaughrean, M. J. 2007, A&A, 462, L53 7. Coffey, D., Bacciotti, F., Ray, T. P., Eisl¨offel, J., & Woitas, J. 2007, ApJ, 663, 350 8. Davis, C. J., Ray, T. P., Desroches, L., & Aspin, C. 2001, MNRAS, 326, 524 9. Davis, C. J., Stern, L., Ray, T. P., & Chrysostomou, A. 2002, A&A, 382, 1021 10. Davis, C. J., Whelan, E., Ray, T. P., & Chrysostomou, A. 2003, A&A, 397, 693 11. Davis, C. J., Nisini, B., Takami, M., Pyo, T.-S., Smith, M. D., Whelan, E., Ray, T. P., & Chrysostomou, A. 2006, ApJ, 639, 969 12. Davis, C. et al., in preparation

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13. Dionatos, O., Nisini, B., Giannini, T., Garcia Lopez, R., Davis, C., J., Smith, M. D. 2008, ApJ, in press 14. Dionatos, O., Nisini, B., Cabrit S., Kristensen L. E., A&A, in preparation 15. Dougados, C., Bacciotti, F., Cabrit, S., Nisini, B., Jets from Young Stars IV, Lecture Notes in Physics, Berlin Springer Verlag, in press 16. Garcia, P. J. V., Ferreira, J., Cabrit, S., & Binette, L. 2001, A&A, 377, 589 17. Garcia Lopez, R., Nisini, B., Giannini, T., Eisl¨offel, J., Bacciotti, F., & Podio, L. 2008, A&A, 487, 1019 18. Garcia Lopez, R., Nisini, Eisl¨offel, J.,B., Giannini, T., Bacciotti, F. 2008, A&A, submitted 19. Giannini, T., McCoey, C., Caratti o Garatti, A., Nisini, B., Lorenzetti, D., & Flower, D. R. 2004, A&A, 419, 999 20. Giannini, T., McCoey, C., Nisini, B., Cabrit, S., Caratti o Garatti, A., Calzoletti, L., & Flower, D. R. 2006, A&A, 459, 821 21. Giannini, T., Calzoletti, L., Nisini, B., Davis, C. J., Eisl¨offel, J., & Smith, M. D. 2008, A&A, 481, 123 22. Gueth, F., Guilloteau, S., 1999, A&A, 343, 571 23. Guilloteau, S., Bachiller, R., Fuente, A., Lucas, R., 1992, A&A, 266, 520 24. Hirano, N., Liu, S.-Y., Shang, H., Ho, P. T. P., Huang, H.-C., Kuan, Y.-J., McCaughrean, M. J., & Zhang, Q. 2006, ApJ, 636, L141 25. Hirth, G. A., Mundt, R., & Solf, J. 1997, A&AS, 126, 437 26. Itoh, Y., et al. 2000, PASJ, 52, 81 27. K¨onigl, A., Pudritz, R., 2000, Protostars and Planets IV (PPIV), p.7591 28. Lavalley-Fouquet, C., Cabrit, S., & Dougados, C. 2000, A&A, 356, L41 29. Le Bourlot, J., Pineau des Forˆets, G., Flower, D. R., & Cabrit, S. 2002, MNRAS, 332, 985 30. Lee, C.-F., Ho, P. T. P., Hirano, N., Beuther, H., Bourke, T. L., Shang, H., & Zhang, Q. 2007, ApJ, 659, 499 31. Nisini, B. 2008, Lecture Notes in Physics, Berlin Springer Verlag, 742, 79 32. Nisini, B., Caratti o Garatti, A., Giannini, T., & Lorenzetti, D. 2002, A&A, 393, 1035 33. Nisini, B., Bacciotti, F., Giannini, T., Massi, F., Eisl¨offel, J., Podio, L., & Ray, T. P. 2005, A&A, 441, 159 34. Nisini, B., Benedettini, M., Giannini, T., Codella, C., Lorenzetti, D., Di Giorgio, A. M., Richer, J. S., A&A, 360, 297 35. Nisini, B., Codella, C., Giannini, T., Richer, J. S., 2002, A&A, 395, L25 36. Nisini, B., Codella, C., Giannini, T., Santiago Garcia, J., Richer, J. S., Bachiller, R., Tafalla, M., 2007, A&A, 462, 163 37. Pesenti, N., Dougados, C., Cabrit, S., O’Brien, D., Garcia, P., & Ferreira, J. 2003, A&A, 410, 155 38. Podio, L., Bacciotti, F., Nisini, B., Eisl¨offel, J., Massi, F., Giannini, T., & Ray, T. P. 2006, A&A, 456, 189 39. Pyo, T.-S., Hayashi, M., Kobayashi, N., Terada, H., Goto, M., Yamashita, T., Tokunaga, A. T., & Itoh, Y. 2002, ApJ, 570, 724 40. Ray, T., Dougados, C., Bacciotti, F., Eisl¨offel, J., & Chrysostomou, A. 2007, Protostars and Planets V, 231 41. Tafalla, M., Santiago, J., Johnstone, D., & Bachiller, R. 2004, A&A, 423, L21 42. Takami, M., et al. 2006, ApJ, 641, 357

Accretion Luminosity of Embedded Protostars Simone Antoniucci

Abstract We present the main results of a series of near infrared spectroscopic observations of a sample of young stellar sources classified as Class Is. Our aim is to derive the accretion luminosity of the objects and investigate their actual accretion activity. For the three sources of our sample that drive a massive jet, we discuss the connection between the accretion and ejection processes, by deriving the mass ejection rate to mass accretion rate ratio.

1 Accretion in Class I objects In the classical Young Stellar Objects (YSOs) classification scheme based on the shape of the Spectral Energy Distribution (SED) [4], Class I are more embedded and less evolved objects with respect to the optically visible Class II objects (T Tauri stars), and are expected to be characterised by higher mass accretion rates. The large extinction coupled with the strong emission excess due to the circumstellar activity, make it often difficult, for these embedded sources, to disentangle the properties of the central star from those of the various active circumstellar regions. This is the main reason why the mass accretion rate of this Class of YSOs is still poorly constrained from an observational point of view, being unknown in particular the fraction of the source luminosity due to accretion only. Efficient diagnostic tools are therefore needed to probe the actual accretion activity of embedded young sources. This is necessary also to clarify the real evolutionary status of the objects, since the SED classification is prone to biases from orientation effects [9]. The work we present aims at determining the accretion luminosity Lacc for a sample of objects classified as Class Is from the analysis of the features detected in their near IR spectra. The accretion activity in some of the investigated objects is indirectly testified also by the presence of jet-like structures extending down to the source itself. For these jet-driving sources, we study the connection between the

S. Antoniucci () INAF - Osservatorio Astronomico di Roma, Via di Frascati, 33 - I-00040 - Monte Porzio Catone (RM) e-mail: [email protected] K. Tsinganos et al. (eds.), Protostellar Jets in Context, Astrophysics and Space Science Proceedings, DOI 10.1007/978-3-642-00576-3 27, c Springer-Verlag Berlin Heidelberg 2009 

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accretion and ejection process by determining the ratio between the mass accretion and the mass ejection rates and by analysing the accretion and ejection signatures visible in the spectra.

2 Sample and Observations The sample we have analysed is composed of 16 Class I sources belonging to nearby star-forming regions [6, 1, 2]. The measurements were carried out using the VLT-ISAAC spectrometer and consist, for each object, of three medium resolution (R  9000) spectra of the sources in the H and K bands (Fig. 1). In the observed

Fig. 1 Examples of H and K band spectra of two sources of the sample: IRS 2 in CrA and the jet-driving source HH34 IRS in Ori. Several emission features are detected, in particular HI recombination lines of the Brackett series, the NaI doublet around 2.21 m, the CO bandhead at 2.3 m, and, in the case of the HH34 IRS, the classical H2 and [FeII] jet lines. A series of photospheric absorption features are observed in the IRS 2 spectrum, which can be used to spectrally classify the source and derive its stellar parameters (from [6] and [1])

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Fig. 1 (continued)

wavelength ranges lie important emission features tracing the accretion region, the accretion disc and jets, as well as several absorption lines that can be used as a diagnostic tool for spectral classification [3, 6]. For the observations of the 3 jet-driving sources of the sample (namely HH34 IRS, HH26 IRS and HH46 IRS), the slit has been aligned along the jet axis.

3 Spectra Several emission features commonly observed in young sources are detected in the spectra: HI recombination emission, that is believed to trace the accretion region, NaI and CO features tracing the gas in the accretion disc, and H2 , and [Fe II] lines from the jet shock regions. However, the presence and the intensity of these emission features are varying among the objects; as a result, the spectra display quite different

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characteristics at medium resolution, as can be easily seen from the spectra shown in Fig. 1. The different line flux ratios observed in the objects (e.g. Br / NaI) suggest that the ionisation structure of the inner regions the lines are tracing is not the same in all the sources. Absorption lines from the protostellar photosphere are detected in a few objects, depending on the S/N ratio of the observations and on the amount of continuum excess emission (veiling), which is expected to be connected to the accretion activity.

4 Accretion Diagnostics In the analysis of our near infrared spectra we have used two different methods to derive the object accretion luminosity Lacc . The first method is based on the detection of spectral absorption features from the stellar photosphere. Whenever these absorption lines are observed, the spectral type of the protostar and the excess emission veiling can both be derived. Using the intrinsic stellar parameters of the inferred spectral type (such as colours and bolometric correction, given by some evolutionary model, e.g. [8]) and the H and K band photometry, it is then possible to derive all the source stellar parameters, and in particular the luminosity L (see [6] for details). The accretion luminosity is then obtained as Lacc D Lbol  L , where the bolometric luminosity of the object Lbol is usually estimated by integrating the SED from the near IR to the mm region. Alternatively, for sources where absorption features cannot be detected, an estimate of Lacc can be derived from the flux of the HI recombination emission lines in the spectra, using the relation found by Muzerolle et al. [5] on the basis of observations of T Tauri stars, that directly connects the accretion luminosity to the HI Br emission line luminosity. Although this method provides in principle an easy way to derive the accretion luminosity, it relies on the assumption that the Muzerolle relation is still valid for the more embedded sources of our sample and requires also a solid determination of the extinction towards the sources, needed for de-reddening the measured Br fluxes. Standard methods for deriving the extinction (based for instance on the observed colours or on the depth of the 9.7 m silicate feature) are often not reliable in the case of embedded young sources, due to the presence of scattered light and to the different physical properties (with respect to those of the ISM) of the dust grains in the circumstellar regions of YSOs. Hence, we have inferred the extinction using a “self-consistency” method based on the fact that the value of L implied by the flux of the de-reddened Br (L D Lbol  Lacc ) must be consistent with the L computed from the observed magnitude of the object once reasonable assumptions on the stellar properties of the source have been made (see [1] for details).

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5 Results 5.1 ADYOs The Class Is of the sample include sources with a large variety of accretion properties. We find that only 8 objects out of 16 actually present a luminosity dominated by accretion (Lacc=Lbol > 0:5), as displayed in Fig. 2, that shows the ratio Lacc=Lbol plotted versus Lbol . We note also that there is no clear correlation between Lacc and Lbol . Our analysis therefore suggests that only a limited number of sources normally classified as Class I actually have accretion-dominated luminosities, pointing to the presence of a bias in the classical YSO classification scheme. On the basis of the small sample we have gathered, we may try to tentatively define some properties characterising these Accretion-Dominated Young Objects (ADYOs), i.e. the sources with Lacc=Lbol > 0.5. These objects have large veilings (rK > 2), as expected if the excess emission is mainly due to the accretion, and most of them (70%) present NaI and CO features in emission, indicating the presence of hot gas in an active accretion disc, although there is no direct correlation between the flux of these features and the accretion luminosity. Sources with massive and dense jets detected very close to the exciting object have high Lacc=Lbol ratios, but the presence of a jet seems not to be a necessary characteristic of an ADYO. Fig. 2 The fraction of the total luminosity due to accretion (Lacc=Lbol ) is plotted as a function of the bolometric luminosity Lbol for the 16 Class Is of the sample. We note that only 8 of the objects appear to have accretion-dominated luminosities (Lacc=Lbol > 0:5). Furthermore, there seems to be no correlation between Lacc and Lbol . The 3 jet-driving sources are indicated by triangles, while the squares refer to sources with detected photospheric absorption features Table 1 Derived parameters for the observed jet driving sources of the sample (from [1]) Source Lacc MP acc MP loss (Lˇ / Lacc=Lbol (107 Mˇ yr1 ) (107 Mˇ yr1 ) MP loss=MP acc HH26 IRS 3.2 0.5 8.5 0.2–0.5 0.02–0.06 HH34 IRS 13.3 0.8 41.1 0.4–1.2 0.01–0.03 HH46 IRS 1.5 0.2 2.2 0.3–2.0 0.14–0.90

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5.2 Accretion and Ejection The mass ejection rates (MP ej ) for jet sources can be estimated from measurements of emission lines flux (e.g. H2 , [FeII]) and jet kinematics (e.g. [7]). The mass accretion rate (MP acc ) is derived from the accretion luminosity using the inferred estimates for the mass and radius of the source (see for example [1]). The MP ej=MP acc ratio is an important parameter indicating the efficiency of the accretion process. Table 1 displays the MP ej=MP acc found in the jet-driving sources of the sample, that are consistent with the values predicted by established MHD models (typically in the range 0.01 – 0.5). In general, we note that there seems to be no precise connection between the presence and flux of accretion and ejection signatures in the spectra of jet-driving sources. Finally, the observed low Br /NaI and high CO/NaI ratios in these objects (with respect to sources without jets) suggest the presence of massive accretion discs characterised by large amounts of warm molecular gas.

References 1. Antoniucci, S., Nisini, B., Giannini, T., Lorenzetti, D.: Accretion and ejection properties of embedded protostars: the case of HH26, HH34, and HH46 IRS. Astron. Astrophys. 479, 503– 514 (2008). DOI 10.1051/0004-6361:20077468 2. Antoniucci, S., et al.: in preparation (2009) 3. Greene, T.P., Lada, C.J.: Spectroscopic Detection of a Stellar-like Photosphere in an Accreting Protostar. Astron. J 124, 2185–2193 (2002). DOI 10.1086/342861 4. Lada, C.J., Wilking, B.A.: The nature of the embedded population in the Rho Ophiuchi dark cloud - Mid-infrared observations. Astrophys. J. 287, 610–621 (1984). DOI 10.1086/162719 5. Muzerolle, J., Hartmann, L., Calvet, N.: A Brgamma Probe of Disk Accretion in T Tauri Stars and Embedded Young Stellar Objects. Astron. J. 116, 2965–2974 (1998). DOI 10.1086/300636 6. Nisini, B., Antoniucci, S., Giannini, T., Lorenzetti, D.: Probing the embedded YSOs of the R CrA region through VLT-ISAAC spectroscopy. Astron. Astrophys. 429, 543–557 (2005). DOI 10.1051/0004-6361:20041409 7. Podio, L., Bacciotti, F., Nisini, B., Eisl¨offel, J., Massi, F., Giannini, T., Ray, T.P.: Recipes for stellar jets: results of combined optical/infrared diagnostics. Astron. Astrophys. 456, 189–204 (2006). DOI 10.1051/0004-6361:20054156 8. Siess, L., Dufour, E., Forestini, M.: An internet server for pre-main sequence tracks of low- and intermediate-mass stars. Astron. Astrophys. 358, 593–599 (2000) 9. Whitney, B.A., Wood, K., Bjorkman, J.E., Wolff, M.J.: Two-dimensional Radiative Transfer in Protostellar Envelopes. I. Effects of Geometry on Class I Sources. Astrophys. J. 591, 1049–1063 (2003). DOI 10.1086/375415

Resolved Inner Jets from T Tauri Stars Francesca Bacciotti

Abstract Understanding how jets are generated around an accreting object is of wide interest to the astrophysical community. A viable way to test observationally the proposed theories is to investigate the inner regions of jets emanated by nearby young stars in the T Tauri phase, that are not deeply embedded anymore in the formation material. Instruments working at sub-arcsecond angular resolution, like those on-board the Hubble Space Telescope, or mounted at ground-based telescopes equipped with Adaptive Optics, allow us to probe the flows in their acceleration and collimation region. The application to these data of new methodologies of analysis and spectral diagnostics have brought an impressive wealth of morphological, kinematical and physical information, providing the most stringent constraints to date for the various models. Here the recent determinations of jet widths, velocity fields, gas excitation parameters and mass outflow rates are reviewed. In addition, a critical discussion is offered on the available indications for jet rotation, i.e. that jets may indeed transport the excess angular momentum away from the system.

F. Bacciotti () Istituto Nazionale di Astrofisica (INAF) - Osservatorio Astrofisico di Arcetri, Largo E. Fermi 5, I-50125, Firenze, Italy e-mail: [email protected]

K. Tsinganos et al. (eds.), Protostellar Jets in Context, Astrophysics and Space Science Proceedings, DOI 10.1007/978-3-642-00576-3 28, c Springer-Verlag Berlin Heidelberg 2009 

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1 Introduction Collimated jets of material moving at extremely high velocities, associated to accretion disks orbiting around the jet source, are observed in a large variety of astronomical objects, over many different scales in size and mass. Despite the large number of studies devoted to this subject, however, open problems still require further investigation. For example, what is the agent that extracts the excess angular momentum from the system? What is the nature of the relationship between accretion and ejection? Which is the most likely jet generation mechanism? Is the process similar at all masses? The best way to attack observationally these issues is to observe jets emanated by young stellar objects. Such jets are close to Earth and easily observable (see, e.g., the review by [10]), and the properties of their sources are well known. The main advantage, however, is that contrary to their extragalactic counterparts, these nebulae have an emission line spectrum, that provides, through appropriate diagnostics, important information about the jet physical properties (see, for example, [4, 33, 41, 21]). To answer the questions related to the accretion/ejection engine, however, one has to investigate the conditions of the flow in the immediate environment of the young star. This suggests to explore the case of the jets from T Tauri stars (TTSs), that are at a relatively evolved stage and around which the environment has already been cleared. According to the models most of the dynamically interesting phenomena are expected to occur within 100–200 AU from the source [48, 24, 42, 47]. To explore this region, however, even for the nearest star formation sites (at 120–140 pc) one needs sub-arcsecond angular resolution. Moreover, the investigations have to be conducted in tracers belonging to different wavelength bands. Diffraction-limited resolution and full wavelength coverage is primarily achieved with observations from outside the Earth’s atmosphere, with telescopes mounted on orbiting spcecrafts, like the Hubble Space Telescope. Many progresses, however, have been made recently also in ground-based astronomy with the development of Adaptive Optics associated to 8m-class telescopes. This review is a description of the state-of-the-art sub-arcsecond resolution observations of the initial section of stellar jets as well as of the most interesting derived results.

2 Observations from the Ground with Adaptive Optics Angular resolution beyond the limit imposed by atmospheric seeing is achievable from the ground using Adaptive Optics (AO) in cascade to large telescopes, as is the case of the ESO Very Large Telescope in Paranal or of the Japanese SUBARU (for a description of AO systems see [23]). In the Near-Infared (NIR) range, results are competitive with those obtainable with HST, reaching 0.00 1 – 0.00 2 resolution. Starting from the pioneering AO studies of T Tauri jets presented in [20, 31], big progresses have been made, and we are now provided, for example, with detailed Position-Velocity diagrams of the initial jet region to be compared with synthetic

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maps developed by theorists (e.g. [2, 44], and see the contribution by S. Cabrit in this volume). More recently, velocity channel maps of the jet base have been constructed for the jets from two classical TTSs: the RY Tau jet with the Integral Field Spectrograph OASIS mounted at the Canada-France-Hawaii-Telescope [1], and the DG Tau jet, using SINFONI at the VLT [2]. From such maps, a number of interesting properties can be deduced, as the jet diameter in the collimation zone, the variation of morphology with different velocities, and kinematical gradients across the jet width. Interestingly, AO studies are now conducted also for the Herbig AeBe stars, the immediately higher mass counterparts of TTSs. A good example is the Adaptive-Optics assisted observations of the jet form LkH˛ 233 [39], that appear to indicate that the properties at the base of jets from higher mass sources are completely analogous to the one observed in the T Tauri case. This, together with the recent detections of jets in Brown Dwarfs [51], appears to point toward an universality at all masses of the mechanism for star formation, that has been originally developed for isolated T Tauri stars. The obtained results demonstrate the validity of the use of AO to observe stellar jets. In particular, the higher sensitivity and smaller diffraction limit of large telescopes is clearly a plus with respect to using relatively small mirrors in space devices. AO observations, however, also present disadvantages: one needs a very bright star in the field (possibly the jet source itself) or a laser guide star to correct the wavefront; the target is better seen close to the zenith, which can pose problems observing from the southern hemisphere; AO systems have only been developed for a limited spectral range (usually in the NIR); until recently, only low spectral resolution instruments were associated to AO systems, although second generation instruments (e.g. CRIRES at the VLT) are improving upon this point; finally, the original signal is substantially reduced by the AO correction, that, also, depends on the actual seeing after all. Therefore, the utility and importance of space observations is far from being ruled out.

3 Observations from Space with the HST The contribution given by the Hubble Space Telescope (HST) to the study of stellar jets is of paramount importance, as described in, e.g. [10, 9, 9]. The Wide Field Planetary Camera2 (WFPC2) has provided the most detailed images to date of the beam and bow shocks of the jets, resolved at 0.00 1 in the various emission lines (e.g. [29, 26]). On the other hand, spectroscopic studies conducted with the Space Telescope Imaging Spectrograph (STIS) have allowed us to get information on the jet properties in the first arcseconds from the source, taking advantage of the fact that the star continuum can easily be subtracted from a spectrum. Stepping the slit every 0.00 07 across the flow, for example, high angular resolution 3-D datacubes (2-D spatial, 1-D in velocity) have been built, form which images and ‘velocity channel maps’ of the jet base were obtained for the DG Tau, RW Aur and LkH˛

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Fig. 1 Left: reconstrucetd images of the jet from the Herbig AeBe star LkH˛ 233 in various emission lines detected in multiple HST/STIS spectra (adapted from [36]). Right): transverse position-velocity (PV) diagrams of the physical quantities in the red lobe of the jet from the TTS Th28. The PV maps are obtained applying the so-called BE technique to optical HST/STIS spectra taken with the slit placed at 0.00 3 from the star in the ‘perpendicular’ configuration (adapted from [17])

jets [6, 53, 36]. The DG Tau flow in particular shows an onion-like kinematic structure, with faster flow concentrated towards the axis, as predicted by magnetohydrodynamic models. Moreover, the emission in the various lines can be combined to diagnose the excitation conditions of the plasma. [5] applied for the first time to HST/WFPC2 filtered images of the HH30 jet in Taurus a diagnostic technique developed in [4] to find the physical conditions of the emitting gas (electron density ne and temperature Te , hydrogen ionisation fraction xe , total density nH and mass outflow rate MP j ). The method, referred to as the ‘BE technique’, from the name of the authors, was built upon the recognition of the fact that charge exchange between O and N with H is the dominant mechanism in the ionisation balance of these species (see also [21]). Being this procedure fast and easy to apply (with respect, for example, to a grid of shock models), other similar analyses of large STIS datasets followed, like the one relative to the DG Tau jet datacube in [7], and, more recently, the Lk H˛ 233 jet (see [36] and Fig. 1, left panel), and the RW Aur jet ([37] and contribution by S. Melnikov in this volume). Tipically, ne is high close to the star (104  106 cm3 , xe varies between 0.01 and 0.4, total densities are up to 106 cm3 , Te varies between 8 103 and 2 104 K, and the derived mass outflow rates turn out to be between 0.05 and 0.1 the corresponding mass accretion rates onto the central object. The above mentioned study on the RW Aur jet indicates that the well-known, but yet unexplained, asymmetry between the velocities in the two lobes also reflects asymmetries in the density and in the mass-loss rates. The analysis of the Lk H˛ 233 jet, in addition, confirms the results obtained from the NIR diagnostics of data taken

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with AO-assisted observations (see above), that is, jets from Herbig AeBe stars appear to be analogous to the T Tauri jets but scaled toward higher excitation and higher mass outflow rates. Recently, excellent studies have been conducted on several T Tauri jets by [3,28] with STIS slitless spectroscopy. In particular, the physical properties of the HH 30 jet have been studied at 0.00 1 resolution at two different epochs separated by about 1.5 years. This allowed the authors to determine the morphology, the excitation and their temporal evolution in the jet’s collimation and acceleration region, using an improved version of the BE technique that combines all the lines available in the considered STIS optical range. In such a way the range of validity of the technique has been extended to regions of very high excitation and/or high density, like the immediate vicinity of the star, thus providing a mean to directly test the validity of the models proposed for the jet acceleration. Finally, a new approach to exploit the full information provided by STIS spectra for the jet base has been attempted in [17]. Here Position-Velocity diagrams of the physical conditions of the gas in various TTS jets have been constructed applying the BE technique to spectra taken with the slit placed at about 0.00 3 from the star perpendicular to the flow axis (see Fig. 1, right panel). This way of analysing the gas physics retains the full information of the original spectra, illustrating how the excitation parameters of the emitting parcel can vary not only with distance from the star or the jet axis, but also with the velocity of the gas. This approach appears to be very promising, and a conceptually similar but more exhaustive study of existing STIS spectra of TTS jets is in progress [35].

4 Jet Rotation: A Discussion One of the most exciting findings made possible thanks to high spatial resolution was the detection of asymmetries in the radial velocity across the jet section, that have been interpreted as rotation of the jets around their symmetry axes. In many theoretical models the jets are believed to carry away the excess angular momentum from the system, and, therefore, a trace of rotation should be observable in the outflow immediately above the acceleration region. The first indications for T Tauri jet rotation have been presented in [8] for the flow from DG Tau. For this jet HST/STIS spectra showed at optical wavelengths systematic shifts in radial velocity, typically 5 to 25 ˙5 km/s, at jet positions displaced symmetrically at 20 – 30 AU from the axis and at 50 – 60 AU from the source. These results were followed closely by a similar finding, obtained again with HST/STIS, in the bipolar flow from RW Aur [54]. The detection of rotation is interesting per se, as it supports the most popular models for jet acceleration, based on the idea that jets are launched through the combined action of centrifugal and magnetic forces ([46, 48, 24, 42], and see the reviews by A. Konigl, M. Cai, R. Lovelace, J. Ferreira, S. Cabrit in this book). The possible physical implications of the rotation results, therefore, were immediately

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investigated. Indeed, the application of a few basic equations valid for a generalised stationary magneto-centrifugal wind allowed to confirm that the measured velocity shifts were compatible with the rotation of the flow imposed by the launching mechanism. The derived toroidal velocities were in turn used to estimate other quantities: the ‘footpoint radius’, i.e. the location in the accretion disk from where the observed portion of the jet is launched (e.g. [3]), that turned out to be between 0.5 and 2 AU from the star; the ratio of the the azimuthal to the poloidal component of the magnetic field, that turned out to be consistent with magnetic collimation due to ‘hoop stresses’; finally, the angular momentum carried by the jet: interestingly, for both the DG Tau and RW Aur outflows, this was estimated to be between 60% and 100% of the angular momentum that the inner disk has to loose to accrete at the observed rate. Therefore, the observed velocity shifts appeared to be good candidates to constitute the first observational validation of the theories of jet launching, and the long-awaited answer to the removal of excess angular momentum from accreting systems. But, is the rotation interpretation correct? One reasonable objection is that once the region beyond the Alfv´en surface is reached (typically a few AU above the disk), the wind very likely undergoes to various kinds of instabilities. These typically produce differential velocities across the beam of the order of the sound speed, i.e. of about 10 km s1 , that is similar to the typical velocity shift observed. Other alternative explanations for the observed velocity shifts have been proposed that invoke a launching mechanism subject to pulsed ejection and precession of the ejection direction around the symmetry axis [12], or interaction of the ejected material with a warped disk [49]. Also, one cannot exclude the possibility that asymmetric bow-shocks form in the flow as a consequence of interaction with an ambient medium where the matter is unevenly distributed. These objections, however, are hardly reconciled with the fact that in both the DG Tau and RW Aur jets, where the rotation signatures have been examined at various distances from the star, the asymmetries appear to have a persistent sign all along the flow. In order to elucidate the situation and check if velocity asymmetries attributable to rotation are a general property of stellar outflows, a dedicated survey was needed. Signatures of coherent velocity asymmetries have been sought with different instruments/telescopes/slit configurations in different targets of different ages, in various line tracers, and in the various parts of the same system (bipolar lobes, disk, coaxial molecular outflows, etc.). The number of observative studies conducted recently on this topic led to interesting results, although the sample of targets observed is limited by the fact that both sub-arcsecond angular resolution and a good spectral resolution are needed simultaneously, to resolve at the same time the jet borders and the expected small velocity differences. Studies have been conducted with HST/STIS in both the optical and NUV ranges ([15, 16], see Fig. 2, left panel), with the spectrograph ISAAC at the VLT ([13], Fig. 2, right panel), GNIRS at the GEMINI telescope [18]. In the mm range the rotation of disks and molecular outflows has been investigated with the SMA and the Plateau de Bure Interferometer (PdBI) [11, 40, 14, 32, 33, 34].

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Fig. 2 Left: transverse radial velocity profiles suggesting jet rotation determined from HST/STIS spectra across the jets from the Class II sources DG Tau and Th 28 (adapted from [16]). Right: similar measurements for the HH 26 and HH 72 jets from Class I sources obtained in H2 lines with VLT/ISAAC (adapted from [13]) Table 1 Synopsis of available jet rotation studies Jet  range Obs. mode Vshift DG Tau VIS HST/STIS, para Y 00 , perpb Y 00 VIS 00 00 NUV , perp Y RW Aur VIS HST/STIS, par Y 00 , perp Y 00 VIS TH 28 VIS HST/STIS, perp Y 00 00 NUV , perp Y CW TAU VIS HST/STIS , perp Y HH 30 VIS HST/STIS, perp ? 00 mm PdBI, RVmapc Y/?c HH 26 NIR VLT/ISAAC, perp Y HH 72 NIR VLT/ISAAC, perp Y HH111 NIR GEMINI/GNIRS N HH34 NIR GEMINI/GNIRS Y HH 212 NIR UKIRT/CGS4, par Y 00 NIR GEMINI/GNIRS Y 00 mm (SiO) PdBI, RV map LR Y 00 mm (SiO) SMA, RV map HR Y HH 211 mm (SiO) SMA, RV map Y CB26 mm (12 CO) PdBI, RV map Y a

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Ref. [8, 50] [16] [16] [54, 11] [15] [15] [16] [16, 22] [15, 40] [40] [13] [13] [18] [18] [19, 52] [18] [14] [33] [32] [34]

Set of slits along to the flow; b Slit placed across the flow; c rad. vel. map: rotation found in the disk but not in the molecular outflow; d knots NK1-SK1, see Coffey et al, this volume.

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The results obtained at the time of writing are summarised in Table 1. Transverse velocity gradients have been detected in almost all the targets observed, with characteristics compatible with the rotation expected on theoretical grounds (see [25] and contribution by S. Cabrit in this volume). Positive detections include also jets from Class I sources, meaning that rotation may be present, as one would expect, from the earliest epochs (see contribution by D. Coffey, this volume). Although the improved statistics appears to confirm the detection of jet rotation and the validity of the magneto-centrifugal scenario, some of the results remain of difficult interpretation. For example, very poor velocity differences are found in jets lying on the plane of the sky (like the HH 30 jet), contrary to what one would expect. Moreover, in the HH 212 molecular flow, one of the knots shows opposite sense of rotation with respect to the others and to the flattened disk-like envelope around the source. The most alarming result, however, is the disagreement between the senses of rotation in the RW Aur disk and in the associated bipolar jet. Further investigations appear to be needed to clarify these discrepancies. It should also be considered that present-day facilities are pushed to their limits in rotation studies, while in principle higher angular and spectral resolutions should be used than those available today. Interferometry at various wavelengths and the development of new telescopes for both ground-based and space observations are expected to achieve the performances necessary to secure the finding that jets indeed rotate. Acknowledgements The author thanks the organisers of the conference for the invitation to present this review. These studies were supported in part by the European Community’s Marie Curie Actions - Human Resource and Mobility within the JETSET (Jet Simulations, Experiments and Theory) network under contract MRTN-CT-2004 005592.

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14. Codella, C.; Cabrit, S.; Gueth, F.; Cesaroni, R.; Bacciotti, F.; Lefloch, B.; McCaughrean, M.J.: Astron. Astroph., 462, L53 (2007) 15. Coffey, D., Bacciotti, F., Ray, T.P., Woitas, J., & Eisl¨offel, J.: Astrophys. J., 604, 758–765 (2004) 16. Coffey, D., Bacciotti, F., Ray, T.P., Eisl¨offel, J., & Woitas, J.: Astrophys. J., 663, 350 (2007) 17. Coffey, D., Bacciotti, F. & Podio, L.: Astrophys. J., 689, 1112 (2008) 18. Coffey, D., Chrysostomou, A., Bacciotti, F., Nisini, B. & Davis, C.J.: Lect. Not. in Phys., Springer, this volume (2009) 19. Davis, C.J., Berndsen A., Smith, M.D., Chrysostomou, A. & Hobson, J.: Mon. Not. Astr. Astroph., 314, 241 (2000) 20. Dougados, C., Cabrit, S., Lavalley-Fouquet, C. & M´enard, F.: Astron. Astroph., 357, 61 (2000) 21. Dougados, C., Bacciotti, F., Nisini, B. & Cabrit, S.: Lect. Not. in Phys., Springer, in press (2009) 22. Dougados, C.: priv. comm., (2009) 23. Esposito, S., Pinna, F.: Lect. Not. in Phys., Springer, 742, 45 (2008) 24. Ferreira, J.: Astron. Astroph., 319 340 (1997) 25. Ferreira, J., Dougados, C., Cabrit, S.: Astron. Astroph., 453 785 (2005) 26. Hartigan, P., Morse, J., Reipurth, B., Heathcote, S., & Bally, J.: Astroph. J., 559, L157 (2001) 27. Hartigan, P., Edwards, S., and Pierson, R.: Astrophys. J., 609, 261–276 (2004) 28. Hartigan, P., & Morse, J.: Astroph. J., 660, 426 (2007) 29. Heathcote, S., Morse, J., Hartigan, P., Reipurth, B., Schwartz, R.D., Bally, J., & Stone, J.: Astron. J., 112, 1141 (1996) 30. K¨onigl, A. and Pudritz, R.: In Protostars and Planets IV (V. Mannings, A.P. Boss, and S.S. Russell, eds.), Univ. of Arizona Press, Tucson, 759–XXX (2000) 31. Lavalley-Fouquet, C., Cabrit, S. & Dougados, C. Astron. Astroph., 356 41L (2000) 32. Lee, C-F, Ho, P.T.P., Palau, A., Hirano, N., Bourke, T.L., Shang, H., & Zhang, Q.: Astrophys. J., 670 1188L (2007) 33. Lee, C-F, Ho, P.T.P., Bourke, T.L., Hirano, N., Shang, H., & Zhang, Q.: Astrophys. J., 685 1026L (2008) 34. Launhardt, R., Pavlyuchenkov, Ya., Gueth, F., et al.: Astron. Astroph., in press (2009) 35. Maurri, L., Bacciotti, F., Coffey, D., Podio, L., Eisl¨offel, J., & Ray, T.P.: Astron. Astroph., in prep. (2009) 36. Melnikov, S., Woitas, J., Eisl¨offel, J., Bacciotti, Locatelli, U. & Ray., T.P.: Astron. Astroph., 483 199 (2008) 37. Melnikov, S., Eisl¨offel, J., Bacciotti, Woitas, J. & Ray., T.P.: Astron. Astroph., to be submitted (2009) 38. Nisini, B., Bacciotti, F., Giannini, T., Massi, F., Eisl¨offel, J., Podio, L. & Ray, T.P.: Astron. Astroph., 441, 159 (2005) 39. Perrin, M.D. & Graham, J.R.: Astroph. J., 670, 499 (2007) 40. Pety, J., Gueth, F., Guilloteau, S., & Dutrey, A.: Astron. Astroph., 458, 841 (2006) 41. Podio, L., Bacciotti, F., Nisini, B., Eisl¨offel, J., Massi, F., Giannini, T., Ray, T.P.: Astron. Astroph., 456, 189 (2006) 42. Pudritz, R.E., Ouyed, R., Fendt, Ch. & Brandenburg, A., in Protostars and Planets V, B. Reipurth, D. Jewitt, and K. Keil (eds), University of Arizona Press, Tucson, p. 277–294 (2007) 43. Pyo, T-S., Kobayashi, N., Hayashi, M. et al.: Astroph. J., 590 340 (2003) 44. Pyo, T-S., Hayashi, M., Kobayashi, N., et al.: Astroph. J., 649 836 (2006) 45. Ray, T., Dougados, C., Bacciotti, F., Eislø¨ effel, J., Chrysostomou, A.: in Protostars and Planets V, B. Reipurth, D. Jewitt, and K. Keil (eds), University of Arizona Press, Tucson, p. 231–244 (2007) 46. Sauty, C., & Tsinganos, K.: Astron. Astroph., 287, 893 (1994) 47. Shang, H., Li, Z.-Y., Hirano, N.: in Protostars and Planets V, B. Reipurth, D. Jewitt, and K. Keil (eds), University of Arizona Press, Tucson, p. 261–276 (2007) 48. Shu, F.H., Najita, J., Ostriker, E.C., Shang, H.: Astrophys. J., 455, L155 (1995) 49. Soker, N., Astron. Astroph., 435, 125 (2005)

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Searching for Jet Rotation Signatures in Class 0 and I Jets Deirdre Coffey, Francesca Bacciotti, Antonio Chrysostomou, Brunetta Nisini, and Chris Davis

Abstract In recent years, there has been a number of detections of asymmetries in the radial velocity profile across jets from young stars. The significance of these results is considerable. They may be interpreted as a signature of jet rotation about its symmetry axis, thereby representing the only existing observational indications

D. Coffey () The Dublin Institute for Advanced Studies, Ireland e-mail: [email protected] F. Bacciotti INAF, Osservatorio di Arcetri, Florence, Italy e-mail: [email protected] A. Chrysostomou University of Hertfordshire, Hatfield, U.K. e-mail: [email protected] B. Nisini INAF, Monte Porzio Observatory, Rome, Italy e-mail: [email protected] C. Davis Joint Astronomy Centre, Hilo, Hawaii, U.S.A. e-mail: [email protected]

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supporting the theory that jets extract angular momentum from star-disk systems. However, the possibility that we are indeed observing jet rotation in pre-main sequence systems is undergoing active debate. To test the validity of a rotation argument, we must extend the survey to a larger sample, including younger sources. We present the latest results of a radial velocity analysis on jets from Class 0 and I sources, using high resolution data from the infrared spectrograph GNIRS on GEMINI South. These observations demonstrate the difficulty of conducting this study from the ground, and highlight the necessity for high angular resolution via adaptive optics or space-based facilities.

1 Introduction Star formation theory demands an explanation for how excess angular momentum is removed from accreting material so that it can collapse onto the newly forming star. Protostellar jets [1, 8] can theoretically be ejected via magneto-centrifugal forces [7, 9], but verification is difficult due to instrumental constraints. A significant observational breakthrough has been made in recent years by our team. We have conducted a series of studies revealing indications of jet rotation. The first study constituted observations of the Class 0 HH 212 outflow. They revealed a difference in radial velocity of 2 km s1 across the flow in the H2 2.12 micron near-infrared (NIR) line at 5 arcsec (2 300 AU) from the star [6]. This became the first observational hint of jet rotation, given that the HH 212 disk has the same direction of radial velocity gradient [10]. Optical and near-ultraviolet (NUV) wavelength observations of less embedded Class II jets, harnessing the high resolution of the Hubble Space Telescope (HST), examined jets closer to their ejection point. The DG Tau and RW Aur jets revealed radial velocity asymmetries of 5–20 km s1 within 100 AU of the star [2, 11]. Furthermore, these asymmetries were sustained for 90 AU along the jet in both cases. A survey was then undertaken of eight jets from six T Tauri systems, i.e. including two bipolar jets. Analysis in the optical and NUV consistently showed asymmetries of typically 15–25 km s1 close to the ejection point [4,5]. Asymmetries seemed common amoung Class II jets, but what about earlier evolutionary stages? A pilot study was carried out to check for rotation signatures in Class I flows. Yet again, transverse Doppler gradients were revealed in NIR spectra of HH 26 and HH 2 a at 1 000 AU from the star [3]. As a natural follow up to this pilot study, we now conduct a wider survey of younger sources. We examine three Class 0/I systems in NIR lines. The [Fe II] 1.64 micron emission traces the hot inner parts of the jet while the H2 emission at 2.12 micron traces the warmer outer regions. Given the high spectral resolution (R  17 800), we chose to use GNIRS on GEMINI South.

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2 Observations and Data Reduction Observations were conducted, in queue mode, with the GEMINI Near InfraRed Spectrograph (GNIRS) on GEMINI South. Spectra were obtained of four jets from Class 0 and I sources, at a distance along the outflow of a few arcsec. Observations were made of the initial channel of the HH 34 jet within 1 arcsec of the source, the H knot in the HH 111 jet, and the first knot in the approaching and receding jet from HH 212, namely NK1 and SK1 respectively. All jet targets were observed in two bands to examine [Fe II] 1.64 micron and H2 2.12 micron emission. The resulting long slit spectra had a nominal velocity resolution of 17 km s1 . Typical seeing varied between 0.5 and 0.8 arcsec. The data were calibrated and reduced according to standard procedure, using dedicated GEMINI IRAF tools.

3 Analysis of Spectra This study is very demanding in terms of both spatial and spectral resolution, and so obtaining useful observations from the ground is difficult. Fortunately, we can reach beyond the nominal resolution both spatially and spectrally via profile fitting. As expected, the H2 emission was found to be spatially broad and so is almost always spatially resolved. We can determine the magnitude and direction of the implied jet toroidal velocity by profile fitting in the dispersion direction to obtain a jet radial velocity profile transverse to the flow propagation. The [Fe II] emission, however, is more confined to the jet axis, and so in our data it is never spatially extended enough to be resolved, even under excellent seeing conditions.

4 Results The HH 34 jet was clearly detected in both [Fe II] and H2 emission. The data contained reflected continuum emission from which we measured the seeing. Unfortunately, the jet width is unresolved in both lines. The HH 111-H approaching jet knot is clearly detected in both [Fe II] and H2 emission, although the jet width is unresolved in [Fe II]. H2 emission is extremely faint in spite of the 2 hour integration, and no clear systematic gradient is apparent. The HH 212 NK1 approaching jet knot reveals a clear [Fe II] detection. The intensity is about one third of the signal-to-noise of H2 . The jet width is unresolved in [Fe II], but resolved in H2 emission. Our H2 data, shows only a mere hint of a gradient of 1-2 km s1 . The direction of such a gradient would imply that the northwest side of the flow is more blue-shifted than the south-east side. This is opposite in direction to the asymmetry detected by [6] for NK1. The HH 212 SK1 receding jet knot was unresolved in [Fe II], but resolved in H2 emission. An asymmetry of magnitude 3–5 km s1 was measured in H2 , which

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matches the gradient direction of our NK1 dataset, implying that the sense of derived toroidal velocities in both lobes of the bipolar jet would be in agreement. This sense also matches the direction of the lower resolution H2 measurements of [6] for SK1, and the direction of the Doppler gradient reported across the associated disk [10]. However, in both NK1 and SK1, emission is spatially asymmetric on the opposite sides of the flow axis. This causes inaccuracies in spatial Gaussian fits, giving rise to difficulties in identifying the jet axis, and hence the magnitude of the radial velocity differences at equal distance either side of the axis. Furthermore, as in the case of the HH 111 analysis, we are observing far from the source and so outside factors may influence the jet dynamics which complicate our interpretation. We therefore cannot safely claim in these cases that we are likely to be observing a rotation signature.

5 Conclusions We have conducted a ground-based survey in NIR lines of four jets from Class 0/I sources to search for signatures of jet rotation. In all cases, the [Fe II] emission, which traces the highly collimated atomic jet component, is unresolved under our

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seeing conditions of 0.5 to 0.8 arcsec. We find that the jet from HH 34, examined within 1 arcsec of the driving source, is spatially unresolved in H2 also. We find no gradient in HH 111-H for the spatially resolved H2 emission, for velocity resolution of gradients down to 1-2 km s1 (depending on signal-to-noise). We find only a mere hint of a gradient in the H2 NK1 knot of the HH 212 bipolar flow. Finally, we detect a gradient of  5 km s1 in the SK1 knot, which matches the disk rotation sense. However, both SK1 and NK1 are far from the source and so external factors such as asymmetric shocking become more significant. Also, they present asymmetric spatial profiles, which confuses determinations of the jet axis location. Our analysis illustrates the difficulty in conducting this study from the ground, given the combination of high spectral and spatial resolution demands. Furthermore, it is clear that in order to safely interpret Doppler asymmetries as signatures of jet rotation, there is a need for improved statistics via high resolution multi-wavelength jet observation close to the source, and comparison with the associated disk rotation sense. Acknowledgements Based on observations at the Gemini Observatory, under Program ID GS2006B-Q-46, which is operated by the Association of Universities for Research in Astronomy, Inc., under a cooperative agreement with the NSF on behalf of the Gemini partnership. The present work was supported in part by the European Community’s Marie Curie Actions - Human Resource and Mobility within the JETSET network, under contract MRTN-CT-2004-005592.

References 1. Bally, J., Reipurth, B., & Davis, C.J., 2007, in Protostars & Planets V, B. Reipurth, D. Jewitt & K. Keil (Tuscon: Univ. Arizona Press), 215 2. Bacciotti, B., Ray, T.P., Mundt, R., Eisl¨offel, J., Solf, J., 2002, ApJ, 576, 222 3. Chrysostomou, A., Bacciotti, F., 2008, A&A, 482, 575 4. Coffey, D., Bacciotti, F., Woitas, J., Ray, T.P., & Eisl¨offel, J., 2004, ApJ, 604, 758 5. Coffey, D., Bacciotti, F., Ray, T.P., Eisl¨offel, J., & Woitas, J., 2007, ApJ, 663, 350 6. Davis, C., Berndsen, A., Smith, M.D., Chrysostomou, A., Hobson, J., 2000, MNRAS, 314, 241 7. Pudritz, R.E., Ouyed, R., Fendt, C., & Brandenburg, A., 2007, in Protostars & Planets V, B. Reipurth, D. Jewitt & K. Keil (Tuscon: Univ. Arizona Press), 277 8. Ray, T.P., Dougados, C., Bacciotti, F., Eisl¨offel, J., & Chrysostomou, A., 2007, in Protostars & Planets V, B. Reipurth, D. Jewitt & K. Keil (Tuscon: Univ. Arizona Press), 231 9. Shang, H., Li, Z.-Y., & Hirano, N., 2007, in Protostars & Planets V, B. Reipurth, D. Jewitt & K. Keil (Tuscon: Univ. Arizona Press), 261 10. Wiseman, J., Wootten, A., Zinnecker, H., McCaughrean, M., 2001, ApJ, 550L, 87 11. Woitas, J., Bacciotti, F., Ray, T.P., Marconi, A., Coffey, D., Eisl¨offel, J., 2005, A&A, 432, 149 12. Yang, J., Ohashi, N., Yan, J., Liu, C., Kaifu, N., Kimura, H., 1997, ApJ, 475, 683

Observational Constraints to Steady Jet Models in Young Stars Sylvie Cabrit

Abstract Steady MHD models for protostellar jets (stellar winds, X-wind from the corotation radius, disk winds from an extended region) are critically examined against updated observational constraints. Synthetic model predictions are used whenever available. The origin of H2 cavities/jets in T Tauri stars is addressed and possible rotation signatures in sub-mm SiO jets are carefully re-examined.

1 Introduction Based on theoretical considerations and observational evidence, Ferreira et al., [22] argue that outflows from T Tauri stars probably contain 3 components, illustrated in Fig. 1: (1) an accretion-enhanced steady stellar wind (Panel c) [20, 30], (2) an unsteady magnetospheric wind fed by reconnexions of either X-type (Panel d) [21] or Y-type (Panel f) (see Zanni, this volume); (3) a steady magneto-centrifugal disk

S. Cabrit () LERMA, Observatoire de Paris, 61 Av. de l’Observatoire, 75014 Paris e-mail: [email protected]

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Fig. 1 Possible launch configurations for MHD winds in young accreting stars according to [22]. Only those in Panels a–c are steady and self-collimated. Panel d corresponds to the episodic magnetospheric reconnexion wind discussed by Ferreira (2000 and this volume) and Panel f to that discussed by Zanni (this volume). They only differ by the orientation of the stellar magnetic moment vs. disk field

wind launched either from an extended range of radii (Panel a) [36] or from a narrow annulus near corotation (“X-wind”, Panel b) [46] (see contributions by Sauty, Ferreira, Romanova, Zanni, Koenigl, and Cai for more theoretical details). Which of these components dominates the observed jet mass-flux is still unclear. Since few quantitative predictions are available for unsteady magnetospheric ejections, the present review concentrates on predictions for steady stellar and disk winds and compares them with observational constraints on jet collimation, kinematics, and ejection/accretion ratio [41, 8]. H2 counterparts in T Tauri stars (TTS) are also addressed, and rotation searches in submm jets are carefully re-examined.

2 Jet Collimation Mechanism Considerable progress on the collimation mechanism of protostellar jets has been achieved over the last few years, thanks to sub-arcsecond studies in their inner 200 AU (see Bacciotti, this volume). The current body of data, summarized in Fig. 2, leads to several important conclusions:  The (apparent) opening angle of TTS jets drops to a few degrees at ' 50 AU from

the source, with a typical deconvolved jet radius of 10 AU. The young Class 0 source HH212 shows the same behavior, despite the presence of a much denser infalling envelope [11]. Therefore, jet collimation occurs mainly in the accretion disk atmosphere, not in the envelope.  The large ram pressure of a typical jet at 50 AU, combined with the constraint of low AV  3 mag in TTS, implies that hydrodynamical disk pressure would

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bowshock DG Tau HN Tau CW Tau HH 30 UZ TauE HL Tau RW Aur

Fig. 2 Left: Deconvolved width of the molecular microjet in HH212 (in red) compared with the full range encountered in atomic T Tauri jets. From [11]. Right: Synthetic beam-convolved jet widths for various self-similar MHD disk winds launched from radii  0.07 AU (solid curves), compared with TTS microjets (symbols). From [41]

collimate the jet only if the disk atmosphere is very hot (6000 K) and 1000 times denser than expected from photoionised disk evaporation [8]. Since this is highly unlikely, magnetic jet collimation is required.  Magnetic self-collimation predicts axial density enhancements with apparent widths similar to observed jets, for both X-winds [44] and self-similar disk winds [9, 23]; the latter results are shown in Fig. 2 after adequate convolution by the PSF, assuming an unperturbed innermost streamline at 0.07 AU. More realistic predictions based on MHD simulations of truncated disk winds are under way [50, 51].  External magnetic confinement is necessary for both stellar and magnetospheric winds [7, 22]. For a typical jet mass-flux it would require a poloidal field of a few tens of mG at z D50 AU and 20–100 AU from the axis [8]. Alternatively, the toroidal field in an outer MHD disk-wind is strong enough for confinement (see Fig. 11 in [23]), as confirmed by recent MHD simulations of stellar+disk-winds [33] (see also Matsakos, this volume).

3 Jet Acceleration, Wide-Angle Structure, and Molecular Counterparts 3.1 Kinematics Along the Jet Figure 3, adapted from [38] compares PV diagrams along the jet axis for the DG Tau and RW Aur microjets with synthetic predictions for extended disk winds and the X-wind (both with M? D 0:5Mˇ ). The (deprojected) speed of 300 km s1 con2 strains the magnetic p lever arm parameter of disk winds   .rA =ro / (through 1 Vp ' Vkep .ro / 2  3, cf. [6]):

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Fig. 3 Observed and synthetic PV diagrams of T Tauri microjets along the jet axis (adapted from [38]): (a) cold disk wind model with  D 50 and ro D0.07–1 AU, from [23]; (b) warm disk wind model with  D 13 and ro D0.07–3 AU [8], (c) X-wind from [43], (d) DG Tau PV diagram from [37], (e) RW Aur PV diagram from [38]  self-similar disk winds: A solution with  D 50 produces excessive jet speeds

[23, 38]. A smaller  D 13, obtained when the disk surface is slightly “warm” [14], yields good agreement with observations (Panel b). Slower outer streamlines launched from 1 to 3 AU may explain the accelerating intermediate velocity component observed in DG Tau.  The X-wind model in Panel c is launched entirely from a single radius at 0.06 AU and has a smaller mean N ' 3:5. As a result, the flow is a factor 2 slower and emits over a narrower velocity range than the disk wind. The RW Aur jet is well reproduced. The DG Tau jet would require a larger N ' 6, i.e. a larger magnetic flux [13].

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 Stellar winds: Inner winds traced by Helium P Cygni absorption reach a similar

speed to the resolved jet, suggesting that the two might be closely related [20]. Since TTS rotate at only 10% of break-up, centrifugal acceleration from the star surface is ineffective [22]. Thermal pressure gradients are also ruled out by C IV line profiles indicating T  20,000 K [26], hence turbulent/magnetic pressure gradients would be needed. For example, the Alfv´en-wave driven model of [19] reaches a terminal speed of 300 km s1 . The real challenge here is the total energy needed to power the jet mass-flux (see Sect. 5).

3.2 Constraints from the Wide-Angle Flow Structure HST observations of several jets reveal a transverse velocity decrease of a factor 2 or more from axis to edge (r > 15–30 AU,   20ı ), as illustrated in the left two panels of Fig. 4 (see also [17]). This slow “sheath” closely surrounding the jet core suggests some modification to the current X-wind model.  In the current X-wind model, the wide-angle flow is quasi-radial and has the same

speed as the dense axial flow, being ejected from the same point with a similar magnetic lever arm [13] (see also Cai, this volume). As a result, the swept-up cavity resulting from interaction with the ambient medium quickly expands out to  20 AU from the axis in only 30 years [45]. Subsequent outbursts would propagate into previous wind material and expand even faster. A slow sheath around the jet axis would thus require tight reconfinement of the X-wind by a strong disk magnetic field, as depicted in [48]; the resulting modified X-wind dynamics remain to be computed.  In a self-similar disk-wind, outer streamlines are ejected from increasing disk radii, causing the poloidal speed to drop sharply away from the axis (roughly in proportion to the keplerian law in the disk). As shown in Fig. 4, a range of launch radii of 0.07–3 AU reproduces well the transverse velocity decrease observed in the DG Tau jet [35].

Fig. 4 Left and Middle: Transverse PV diagrams at z D30–45 AU from HST/STIS illustrating the slower sheath at RD15–30 AU around the fast jet core in CW Tau and DG Tau (adapted from [18]). Right: Beam-convolved synthetic transverse PV diagram for the warm disk wind model of Fig. 3b, which also fits the rotation signatures in DG Tau. From [35]

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 In stellar winds, a decrease in flow speed with latitude may be obtained with a

typical pressure distribution and outer confining disk streamlines, as is the case e.g. in the models of [42, 31].

3.3 Molecular H2 Counterparts in TTS Wide, slow H2 cavities encompassing the fast collimated atomic jet have been recently imaged in several TTS [5] (see also Agra-Amboage et al., this volume). As UV and X-ray heating appear insufficient to heat the H2 , two alternative explanations have been proposed: (1) A wide, slow molecular MHD disk wind heated by ambipolar diffusion [52]; (2) A shocked cavity driven at 10–30 km s1 into ambient gas by an unseen wideangle wind, with outbursts every ' 100 yrs to explain the small cavity sizes [53]. The lack of fast shocked atomic gas at wide angles, and the need to quickly refill the cavity with molecular gas between outbursts, may require a slow molecular wideangle wind in this case as well. A fast inner H2 jet is also detected at 100 km s1 from RW Aur A [5]. Entrainment of ambient gas seems unlikely since RW Aur is located in a “hole” of the molecular cloud. Either H2 forms behind dense low-velocity shocks in a dust-free atomic jet [39], or the H2 jet traces a (fast) molecular disk wind. Recent calculations show that molecules may indeed survive in MHD disk winds launched beyond 0.5 AU:

3.4 Molecule Survival in MHD Disk Winds The non-equilibrium chemistry, ionisation, and temperature structure in MHD disk winds has been recently calculated for launch radii of 0.5–9 AU, in a self-similar solution compatible with atomic jet kinematics (see [34] and this volume, for more details). Ionisation and dissociation by stellar FUV and X-rays were included. Ambipolar diffusion heating yields temperatures of 1000–3000 K at 50 AU for Class I-II sources, with 30–100% of the gas remaining molecular at speeds of to 100–20 km s1 . MHD disk winds launched beyond the dust sublimation radius thus appear promising to explain both the fast H2 jets and slow cavities in TTS. Synthetic predictions are in progress to further test this hypothesis.

4 Updated Constraints from Jet Rotation Measurements of transverse velocity shifts across jets suggestive of rotation are reviewed in Bacciotti (this volume). Comparison with theoretical expectations was summarized graphically by [22] using Fig. 5, which plots the specific angular

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Fig. 5 Predicted position in the Vp  rV plane of steady extended MHD disk winds for various launch point ro (solid red curves) and effective magnetic lever arm  (dashed red curves). The thick black curve shows a cut at z D 50 AU across a warm disk wind solution with  D 13 and a range of ro D 0.07–9 AU. The loci of the X-wind (thick red segment), and of stellar winds with various pressure parameters ˇ (in blue), are also indicated. Updated from [22]; colour symbols show new measurements from the compilation of Bacciotti et al. (this volume)

p momentum RV vs. poloidal speed Vp , both scaled by Mˇ =M? . This combination allows to constrain the launch radius ro of a magnetocentrifugal disk wind, as first shown by [4, 3], but also the “effective” magnetic lever arm parameter  , within a factor 2 of the true  (see (10) of [22]). Figure 5 is updated to include datapoints for all TTS and Class 0 sources from Table 1 of Bacciotti (this volume) with M? estimates. Several conclusions may be drawn concerning TTS jets, which are the best resolved (0.1”D15 AU at the Taurus distance):  The modest velocity shifts definitely rule out “cold” MHD disk winds with large

  50 and would require “warm” disk winds with small   13 [22], as concluded also from Fig. 3.  If the rotation interpretation is correct, streamlines dominating optical/UV measurements would be launched from 0.1 to 3 AU (see [41] and refs. therein); the intermediate and high-velocity components of the DG Tau jet may trace a range of ro within the same self-similar disk wind.  The CO conical outflow from the quasi edge-on TTS CB 26 [27] would be launched from a larger radius ro D 5 ˙ 4 AU (depending on the exact inclination i D 85ı 4ı ).

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4.1 Cautionary Notes: The Examples of HH 211 and HH 212 Possible artefacts due e.g. to jet instabilities, precession, or asymmetric shocks are discussed in Bacciotti (this volume). They are definitely present in a few cases [10, 16] (see also Poster by Correia et al., this volume), hence V and the inferred ro and  may be viewed conservatively as upper limits in well-resolved Taurus jets. Uncertainties from low signal-to-noise are discussed in [1]. Here I will insist on 3 complementary cautionary notes, particularly relevant to recent sub-mm studies of more distant Class 0 jets: 1. Biases due to lack of resolution: Synthetic PV diagrams show that rotation signatures in disk winds tend to cancel out at distances to the jet axis less than the beam size [35]. This is of particular concern for submm studies of Class 0 jets in Perseus/Orion, where the 0.3”–1” SMA/PdBI beams lead to a linear resolution 6–30 times worse than in Taurus with HST. The high-velocity jet core being  0.2” wide, its rotation signatures will be strongly suppressed and reliable upper limits on jet rotation require careful modelling of this effect. 2. The inferred ro is extremely sensitive to the adopted Vp , roughly as ro / Vp4=3 [3]. The use of an inadequate Vp value will then lead to erroneous conclusions (see below the case of HH 212) 3. The value of the magnetic lever arm   .rA2 =ro2 / and rotation period 2=˝o implied by rotation measurements is often under-exploited, although it further constrains viable models (see below the case of HH 211). The case of HH 212 (M? D0.15Mˇ ): [29] report transverse shifts at the tips of the innermost SiO bowshocks and infer a launch radius of 0.05–0.3 AU assuming Vp D100–200 km s1 from H2 knot proper motions. However, the material showing transverse shifts is at much lower radial velocities of 1.5–4 km s1 than the fast jet core (projected at 10 km s1 ). Given the jet inclination of i D 86ı [15], the flow speeds are only 20–60 km s1 . The corresponding datapoints in Fig. 5 yield a launch radius ' 0.6 AU with  ' 5:5 ˙ 3 and a rotation period of 1.2 year, favoring an extended disk wind rather than an X-wind, if the data do probe steady jet rotation. The non-detection of velocity shifts across the faster, narrower jet core [16] might be due to lack of resolution (see point (1) above). The case of HH 211: With M? D0.06Mˇ and assuming a poloidal speed of Vp D 150˙50 km s1 (i.e. a jet inclination of 80ı –85ı ) the observed velocity shifts yield a launch radius of 0.06–0.15 AU suggestive of an X-wind [28]. However, (10) of [22] and Fig. 5 yields  D 20 ˙ 5, much larger than in X-wind models with reasonable magnetic fluxes (N D3.5–6, [13]). The rotation period at ro of 22–88 days is also longer than expected for a very young protostar, believed to rotate faster than a TTS (5–10 days). The data would again seem to favor an extended disk wind with  ' 20 launched beyond corotation rather than an X-wind, if the rotation interpretation is correct.

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5 Jet Mass-Flux and Jet Power Since outflow signatures in TTS appear powered by accretion onto the star, the ejection/accretion ratio is a key constraint for steady jet mechanisms. Its average (two-sided) value in TTS jets has been revised upward from 0.01 [25] to f  .2MP j /=MP acc ' 0:2 [8,12], similar to the value in Class I jets [24,2]. A ratio f ' 0:2 appears energetically challenging to both steady stellar winds and the X-wind model as currently modelled (see [22] and Ferreira, this volume):  In steady stellar winds, lifting this mass-flux from the deep potential well requires

a net energy deposition rate in the wind of 30% of Lacc [30], but an even larger amount needs to be injected, since losses due to dissipation or wave divergence are unavoidable [22]. A highly efficient, steady energy injection mechanism near the star surface would be needed.  X-wind model: In steady-state, a magneto-centrifugal wind cannot extract more than the mechanical energy liberated by the accretion flow across the launch region. Assuming keplerian rotation across the X-wind region, this imposes f D .rX =rX / ŒVkep .rX /=Vj 2 (see Ferreira, this volume). With rX ' .h=r/rX , f ' .1=4  1=10/ .h=r/ 0:2. The large jet mass loading would thus require a strongly non-keplerian profile and/or large pressure gradient, not yet included in current models.  In self-similar disk winds, mechanical flux is extracted across a much wider disk region. The ejection to accretion ratio f D ln.rout =ri n /=.2  2/ can reproduce the observed value with  ' 6  13 and rout =ri n  10 [22].

6 Conclusions A detailed comparison of theoretical model predictions with spatially resolved jet properties reveals open issues with all steady models for the jet origin: – Steady MHD stellar winds require an extremely efficient energy injection process at the stellar surface, capable of depositing a large fraction ( 30%) of the accretion energy. – The X-wind model would need external confinement by a disk field [48] to suppress the fast wide-angle flow and explain the low-velocity sheath seen around some jets; The ejection/accretion ratio appears to require non-keplerian rotation or large pressure gradients across the X-wind region (cf. Ferreira, this volume). These aspects remain to be modelled, as well as their impact on X-wind dynamics. – Self-similar MHD disk winds appear compatible with most observational constraints, including recent H2 counterparts in TTS and claimed rotation signatures in atomic/SiO/CO jets, but the latter impose a small magnetic lever arm parameter   5–13, requiring an unidentified heating process at the disk surface [14]. In this scenario, stellar braking would have to be provided by an inner (unsteady)

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reconnexion wind externally confined by the disk wind ([22], Zanni, this volume) and/or by “magnetic towers” (see contributions by Lebedev, Machida, and Romanova in this volume). Acknowledgements This work was supported through the Marie Curie Research Training Network JETSET (Jet Simulations, Experiments and Theory) under contract MRTN-CT-2004-005592. Stimulating discussions with the meeting participants and the members of the JETSET network are gratefully acknowledged.

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Searching for Brown Dwarf Outflows Emma M. Whelan, Tom Ray, Francesca Bacciotti, Sofia Randich, and Antonella Natta

Abstract As outflow activity in low mass protostars is strongly connected to accretion it is reasonable to expect accreting brown dwarfs to also be driving outflows. In the last three years we have searched for brown dwarf outflows using high quality optical spectra obtained with UVES on the VLT and the technique of spectroastrometry. To date five brown dwarf outflows have been discovered. Here the method is discussed and the results to date outlined.

1 Introduction It is now apparent that protostellar-like outflows commonly accompany the formation and evolution of brown dwarfs (BDs; [12, 16, 18]). The overall motivation of this project is to investigate the validity of the accretion/ejection models for the

E.M. Whelan () and T. Ray Dublin Institute for Advanced Studies, e-mail: [email protected] F. Bacciotti, S. Randich, and A. Natta Osservatorio Astrofisico di Arcetri

K. Tsinganos et al. (eds.), Protostellar Jets in Context, Astrophysics and Space Science Proceedings, DOI 10.1007/978-3-642-00576-3 31, c Springer-Verlag Berlin Heidelberg 2009 

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Table 1 Brown Dwarfs and a Very Low Mass Star Discovered to have Optical Outflows. The papers referencing the mass of these objects are given in square brackets in column 4 Source RA (J2000) Dec (J2000) Spectral Type Mass (Mˇ ) ISO-ChaI 217 11 09 52.0 76 39 12 M6.2 0.08 [8] 2MASS1207-3932 12 07 33.4 39 32 54.0 M8.5 0.024 [10] ISO-Oph 32 16 26 22.05 24 44 37.5 M8 0.04 [11] ISO-Oph 102 16 27 06.5 24 41 47.1 M6.5 0.06 [11] LS-RCrA 1 19 01 33.7 37 00 30 M6.5 0.04 [1]

formation of solar-mass stars at the lowest masses, through the study of BD outflows. Ultimately, results will be used to form a better understanding of the outflow mechanism in general. Much of what is known about low mass star formation and in particular the connection between magntospheric accretion and outflow activity comes from studying the classical T Tauri stars (CTTS). Forbidden emission lines (FELs) have proved to be effective tracers of outflow activity in CTTSs and to date they have been used to explore the kinematics, morphology and physical conditions of jets, at high angular resolution [13]. The so-called traditional tracers of CTT jets, i.e. [OI]6300,6363, [SII]6716,6731, [NII]6583 lines, are found in the spectra of BDs. This finding was the first indication that BDs launch outflows. BDs outflows are difficult to detect as they are faint and FEL tracers are only extended on milli-arcsecond scales. Hence long exposure times (1.5 hours with the VLT) and high angular resolution techniques are needed. Our approach is to obtain high spatial resolution spectra with the UV-visual echelle spectrometer (UVES) on the VLT and to analyse the origin of key lines with spectro-astrometry. Table 1.1 lists the objects found to date to be driving optical outflows. Four of these objects have derived masses placing them well within the BD mass range. The fifth, ISO-ChaI 217 has an estimated mass of 80 MJUP placing it just above the hydrogen burning mass limit (HBML).

2 Targets, Observations and Analysis The BDs targeted for this study had all shown evidence of strong T Tauri-like accretion and several were already known to exhibit forbidden emission. High resolution UVES spectra were obtained, reduced using standard IRAF routines and, as stated above, analysed with spectro-astrometry. For ISO-Cha I 217 and ISO-Oph 32, spectra were obtained at orthogonal slit position angles (PA), allowing the PAs of the outflows to be estimated. For the majority of protostars with jets the FEL regions of their spectra are very obviously spatially extended, particularly at high velocities. For example the FELs of the CTTS DG Tau trace a  12” knotty jet [15]. However, when the emission is very compact and close to the source (as in the case of BDs) specialised methods such as spectro-astrometry have to be used to probe its nature. Spectro-astrometry utilises simple Gaussian Fitting to investigate the offset,

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with respect to the star/BD, of emission line features smaller than the seeing disc of the observation. The result of this technique is an offset-velocity diagram, with the displacement of a spectral feature (e.g. the [OI]6300) line shown as a function of velocity and relative to the continuum centroid of the object [19]. The beauty of this technique lies in the fact that its accuracy depends primarily on the signal to noise (S/N) of the observation. Formally, the spectro-astrometric accuracy p is given as follows,  D .Seeing/=.2:3548 Np /, where Np is the number of detected photons. Hence, even under conditions of poor seeing, a high accuracy can be achieved. Refer to Whelan et al. [18] and Whelan & Garcia [19], for further details on the spectro-astrometric technique, including information on spectro-astrometric artifacts and on how the S/N can be increased in BD spectra.

3 Results In this section we will briefly discuss the objects analysed to date (in order of their investigation) and the properties of the outflows uncovered.  ISO-Oph 102 and ISO-Oph 32: ISO-Oph 102 was the first BD investigated

using spectro-astrometry and the results of this investigation were published in Whelan et al 2005. Hence it is the first BD discovered to have an outflow. The [OI]6300, 6363, [NII]6583 and [SII]6731 lines were all detected along with the H˛ line which exhibited a P-Cygni like profile. Our investigation revealed the [OI] and [SII] lines to be offset to a distance of  100 mas at a velocity of  40 kms1 . The [NII]6583 was too faint for spectro-astrometric analysis. Only blue-shifted emission was detected, as is commonly seen in CTTSs. The circumstellar disk is assumed to be obscuring the red-shifted part of the flow. The scale of the blueshifted offset would suggest a minimum (projected) disk radius of 100 mas ( 15 AU at the distance of the -Ophiuchi cloud) in order to hide the redshifted component. ISO-Oph 32 is a 40 MJUP BD which again was known to be a strong accretor [13]. Our recent analysis revealed weak [OI]6300, 6363 emission, similar in strength to what we detected for 2MASS1207-3932. The [OI]6300 line is blue-shifted to a velocity of  30 kms1 and offset to  60 mas at 0ı and 110 mas at 90ı . This constrains the PA of this outflow at  60ı (E of N).  2MASS1207-3932: UVES spectra of the 24 MJUP BD 2MASS1207-3932 obtained in May 2006 revealed the presence of strong H˛ emission and the [OI]6300, 6363 emission lines. Only the [OI]6300 was strong enough for spectro-astrometric analysis. The line profile was double peaked with blue and redshifted emission at  8 kms1 and 4 kms1 , respectively. The blue and red-shifted emission was found to be offset in opposite directions to  85 mas, revealing the presence of a bipolar optical outflow. 2MASS1207-3932 is now the lowest mass galactic object known to drive an outflow. The low radial velocities measured in the [OI]6300 line agree with the hypothesis that 2MASS1207-3932 has a near edge-on accretion disk [19].

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 ISO-Cha I 217: Spectra of ISO-Cha 1 217 at orthogonal positon angles (0ı ,

90ı ) were obtained in September 2007. The FELs of [OI]6300, 6363, [SII]6716, 6731 and [NII]6583 were found to be bright, and spectroastrometric analysis revealed the presence of a bipolar outflow. Systemic velocities of  20; C30 kms1 were measured at both 0ı and 90ı and, interestingly, the red-shifted emission was found to be twice as bright as the blue-shifted. This type of asymmetry has been seen before in CTTSs. Spectroastrometry revealed the red and blue components to be offset to  C=200 mas at 0ı and C=50 mas at 90ı . This suggests a P.A. for the outflow from ISO-Cha I 217 of  15ı (E of N).  LS-RCrA 1: LS-RCrA 1 was the first BD shown to have forbidden emission [2]. Early attempts to spatially map the FELs, using spectra taken with Magellan Inamori Kyocera Echelle ( MIKE ) on the Magellan II telescope failed due to the poor quality of the spectra and hence the faintness of the continuum emission [17]. For this study UVES spectra, taken in June 2003, were obtained from the ESO data archive. Again bright [OI]6300, 6363, [SII]6716, 6731 and [NII]6583 are present in the spectrum and offsets of up to 150 mas were recovered. The H˛ line profile has blue and red-shifted “humps” which our spectro-astrometric analysis has revealed to be offset in opposite directions. Hence, the H˛ line has a component originating in the outflow from LS-RCrA 1. Note that this is the first time that H˛ has been found to be tracing an outflow from a BD. Also note that both lobes of the flow are seen in H˛ while in the FELs only blue-shifted emission is detected. This has been observed in CTTSs and may be explained by the presence of a “dust hole” in the disk, close to the central star [14, 15]. As the H˛ is tracing the outflow much closer to the BD than the FELs, its red-shifted component can be seen through the dust hole in the disk. The FELs originate much further from the star where the red-shifted emission remains obscured by the disk. This result can be taken as further evidence that BD disks evolve in a similar fashion to T Tauri disks and may at some stage harbour planets. Figure 1 shows the results of the spectro-astrometric analysis of the [SII]6731 and H˛ lines. As stated, the long-term aim of this project is to make a comprehensive comparison between protostellar outflows/jets and those driven by BDs. So far results have revealed that optical outflows driven by accreting BDs are common. Notable T Tauri-like properties of the BD outflows include, the predominance of blue-shifted emission (indicating the presence of an accretion disk) and the fact that the measured offsets in the BD FELs lie within the range estimated from BD if BD outflows scale down from T Tauri-jets [16]. In order to truly test whether the same the T Tauri outflow mechanism is operating in the sub-stellar regime, images of BD outflows are needed to investigate collimation. In addition, comparisons between mass outflow and infall rates must be made.

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Fig. 1 The spectro-astrometric analysis of the H˛ and [SII]6731 lines in the spectrum of LS-RCrA 1. The black stars represent the offsets in the line, in the continuum subtracted spectrum, and the open diamonds the position of the continuum. The line and continuum were binned so that the S/N and thus the spectro-astrometric accuracy in both was comparable. The dashed line represents the 1- error in the spectro-astrometric analysis

References 1. Barrado y Navascu´es, D., Mohanty, S., & Jayawardhana, R: Accretion and Outflow in the Substellar Domain: Magellan Spectroscopy of LS-RCrA 1, ApJ, 604, 284, (2004) 2. Fern´andez, M., & Comer´on, F: Intense accretion and mass loss of a very low mass young stellar object, A&A, 380, 264, (2001) 3. Jayawardhana, R., Mohanty, S., & Basri, G: Evidence for a T Tauri Phase in Young Brown Dwarfs, ApJ, 592, 282, (2003) 4. Jayawardhana, R., Ardila, D. R., Stelzer, B., & Haisch, K. E., Jr: A Disk Census for Young Brown Dwarfs, AJ, 126, 1515, (2003) 5. Hartigan, P., Morse, J. A., & Raymond, J: Mass-loss rates, ionization fractions, shock velocities, and magnetic fields of stellar jets, ApJ, 436, 125, (1994) 6. Luhman, K. L., Rieke, G. H., Young, E. T., Cotera, A. S., Chen, H., Rieke, M. J., Schneider, G., & Thompson, R. I. 2000, ApJ, 540, 1016 7. Luhman, K. L., Fazio, G., Megeath, T., Hartmann, L., & Calvet, N: Young brown dwarfs: IMF, disks, spatial distribution, and binarity, Memorie della Societa Astronomica Italiana, 76, 285, (2005) 8. Mohanty, S., Jayawardhana, R., & Basri, G: The T Tauri Phase Down to Nearly Planetary Masses: Echelle Spectra of 82 Very Low Mass Stars and Brown Dwarfs, ApJ, 626, 498, (2005) 9. Mohanty, S., Jayawardhana, R., Hu´elamo, N., & Mamajek, E: The Planetary Mass Companion 2MASS 1207-3932B: Temperature, Mass, and Evidence for an Edge-on Disk, ApJ, 657, 1064, (2007) 10. Muzerolle, J., Luhman, K. L., Brice˜no, C., Hartmann, L., & Calvet, N: Measuring Accretion in Young Substellar Objects: Approaching the Planetary Mass Regime, ApJ, 625, 906, (2005) 11. Natta, A., Testi, L., Muzerolle, J., Randich, S., Comer´on, F., & Persi, P: Accretion in brown dwarfs: An infrared view, A&A, 424, 603, (2004)

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12. Phan-Bao, N., et al: First Confirmed Detection of a Bipolar Molecular Outflow from a Young Brown Dwarf, arXiv:0810.2588 (2008) 13. Ray, T., Dougados, C., Bacciotti, F., Eisl¨offel, J., & Chrysostomou, A: Toward Resolving the Outflow Engine: An Observational Perspective, Protostars and Planets V, 231, (2007) 14. Takami, M., Bailey, J., Gledhill, T. M., Chrysostomou, A., & Hough, J. H: Circumstellar structure of RU Lupi down to au scales, MNRAS, 323, 177, (2001) 15. Whelan, E. T., Ray, T. P., & Davis, C. J: Paschen beta emission as a tracer of outflow activity from T-Tauri stars, as compared to optical forbidden emission 2004, A&A, 417, 247, 2004 16. Whelan, E. T., Ray, T. P., Bacciotti, F., Natta, A., Testi, L., & Randich, S: A resolved outflow of matter from a brown dwarf, Nature, 435, 652, (2005) 17. Whelan, E. T., Ray, T. P., Bacciotti, F., & Jayawardhana, R: Probing outflow activity in very low mass stars and brown dwarfs, New Astronomy Review 49, 582, (2006) 18. Whelan, E. T., Ray, T. P., Randich, S., Bacciotti, F., Jayawardhana, R., Testi, L., Natta, A., & Mohanty, S: Discovery of a Bipolar Outflow from 2MASSW J1207334-393254, a 24 MJ up Brown Dwarf, ApJ, 659, L45, (2007) 19. Whelan, E., & Garcia, P: Spectro-astrometry: The Method, its Limitations, and Applications, Lecture Notes in Physics, Berlin Springer Verlag, 742, 123, (2008)

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Protostellar Jets Driven by Intermediateand High-Mass Protostars: An Evolutionary Scenario? Alessio Caratti o Garatti, Jochen Eisl¨offel, Dirk Froebrich, Brunella Nisini, and Teresa Giannini

Abstract Protostellar jets from intermediate- and high-mass protostars provide an excellent opportunity to understand the mechanisms responsible for intermediateand high-mass star formation. A crucial question is if they are scaled-up versions of their low-mass counterparts. We present a detailed study of the IRAS 20126 + 4104 molecular jet, driven by a 104 Lˇ protostar. The kinematical and physical properties of the jet have been obtained by means of NIR narrow-band imaging, high resolution and low resolution IR spectroscopy. We then compare our findings with those of other high- and low-mass protostellar jets available in the literature, proposing an evolutionary scenario from low- to high-mass jets.

A. Caratti o Garatti () and J. Eisl¨offel Th¨uringer Landessternwarte Tautenburg, Sternwarte 5, D-07778 Tautenburg, Germany e-mail: [email protected]; [email protected] D. Froebrich Centre for Astrophysics and Planetary Science, University of Kent, Canterbury, CT2 7NH, United Kingdom e-mail: [email protected] B. Nisini and T. Giannini INAF - Osservatorio Astronomico di Roma, via Frascati 33, I-00040 Monte Porzio, Italy e-mail: [email protected]; [email protected]

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1 Intermediate- & High-Mass Jets in Context The study of massive outflows and jets is still in its beginnings. Indeed there are several limitations to the observations as the large distance of the massive star-forming regions, the considerable extinction, and the short lifetime of massive young stellar objects (MYSOs). Moreover MYSOs are often grouped in small clusters, which confuses the morphology of massive star-forming regions even more. As a consequence not many massive jets are known up to now, and deep surveys are needed to increase their number (see e g. Stecklum et al. in this volume). Optical and IR studies of intermediate- and high-mass protostellar jets are very rare, and so far only three high-mass jets have been spectroscopically investigated in the NIR [6, 7, 3]. Nevertheless the investigation of these jets is crucial, since they provide insights on high-mass star formation. In particular their detection indicates the presence of an accretion disc, supplying robust evidence of an accreting source with higher accretion/ejection rates [1]. Expanding the number of observations of intermediateand high-mass jets and comparing their general properties with those of low-mass protostellar jets is also crucial for understanding whether differences exist, or if high-mass protostellar jets are just scaled-up versions of their low-mass counterparts (e g. [6, 7]). We investigated the kinematical and physical properties of the IRAS 20126 + 4104 (Lbol 104 Lˇ , MP acc 2  103 Mˇ yr1 [5]) jet by means of NIR narrow-band imaging (H2 and [FeII]), high-resolution (18 500 at 2.12 m) and low-resolution (200–1000 between 0.9–200 m) IR spectroscopy. Our data were collected at the UK Infrared Telescope (with CGS4 and UIST), and at the 3.5-m Italian Telescopio Nazionale Galileo (with NICS). More data were retrieved from the Subaru (CIAO NIR imager) and ISO (SWS and LWS instruments) archives. Figure 1 shows the H2 jet, the positions of the slits and the ISO-SWS FoV.

2 IR Spectra: The Importance of the Cold H2 Component Both imaging and spectroscopy reveal that the jet is mostly molecular. Lowresolution spectra are rich in H2 emission (Fig. 2, left panel), and several pure rotational lines have been detected in the MIR. No ionic emission is detected along the flow, with the exception of faint emission of [FeII] close to the source position. Surprisingly these spectra closely resemble those of low mass jets (Fig. 2, right panel). The ro-vibrational diagrams indicate H2 excitation temperatures of 2000–2500 K, and no signature of fluorescence has been detected. Stratification of temperature is detected only in knot C (Fig. 3, left panel), which can be modelled combining a warm (Tex = 2050 K) and a hot (5200 K) H2 component. Additionally, the ISO-SWS spectrum reveals the presence of a cold component (520 K) with a column density higher than expected. For comparison in Fig. 3 (right panel), the ro-vibrational diagram of a typical low mass jet is presented. Here the gas has similar temperatures, but lower column densities as expected. However, N.H2 / of the

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Fig. 1 NICS H2 mosaic (left panel), showing the entire flow with superimposed positions of the slits and the ISO-SWS FoV. The positions of the H2 knots and the IRAS source are also indicated. The H2 jet has a wide precessing angle of 37ı . A second precession mode with an angle of 8ı has been detected at smaller scale using the high-resolution H2 (continuum subtracted) Subaru image (upper panel)

warm and hot components are one order of magnitude lower, while it is about two in the cold gas. This would imply that a large part of the mass is carried by the cold component. On the other hand, the spatial velocity of the knots (inferred from the radial velocity) is between 50 and 80 km s1 , quite similar to those observed in low mass jets. The estimated LH2 is 8.1 ˙ 0.7 Lˇ , where the cold component contributes about 50% to the whole radiative cooling. The high H2 luminosity suggests the driving source has a significantly increased accretion rate compared to the low-mass YSOs. This is also supported by the measured mass flux rates from H2 lines (MP out (H2 )7.5  104 Mˇ yr1 ), that fairly matches previous CO estimates [10]. Our analysis thus shows that the cold H2 component plays a major role in the kinematics and dynamics of this massive flow.

3 An Evolutionary Scenario? By comparing the measured luminosity of the H2 jet with the source bolometric luminosity (assumed representative of the accretion luminosity), in Fig. 4 we show that IRAS 20126 + 4104 fits the correlation well between these two quantities, as already found for low-mass protostellar jets [4]. The IRAS 20126 + 4104 jet is positioned in the upper-right corner of the diagram. The fit suggests that the relationship also applies to more massive jets in their earliest stage of formation. According to the dynamical timescales of the outflow and the jet (a few times 104 years),

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Fig. 2 Left Panel: 0.9–2.5 m low-resolution spectrum of knot C in the IRAS 20126 + 4104 jet. Right Panel: NIR low-resolution spectrum of HH 320 A, driven by a low mass YSOs

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Fig. 3 For comparison ro-vibrational diagrams of knots C of the IRAS 20126 + 4104 jet (left panel) and of the L1157 (right panel) flow (driven by a low mass YSO) are shown. Different symbols indicate lines coming from different vibrational levels, as coded in the upper right corner of the box (v = 0 lines are located in the upper left part of the diagram)

I-20126 I-16547 I-11101

I-05358

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Fig. 4 LH2 vs Lbol including the IRAS 20126 C 4104 jet (also considering the luminosity of the cold component, the new datapoint is located in the upper-right corner of the diagram). Lower limits values for IRAS 18151-1208, IRAS 11101-5829, IRAS 05358+3543 [9] and IRAS 165474247 jets have been included, as well (see text). For these jets the LH2 was only inferred for the warm component. A dashed line indicates the previous fit from [4]

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IRAS 20126 + 4104 has not yet reached the main sequence (MS) and has not yet developed any hypercompact HII (HCHII) region, which may affect the collimation of the jet/outflow system [1]. Most importantly, the bolometric luminosity of the source mainly comes from accretion [5]. In addition, we compared our results with the LH2 of four high-mass jets previously investigated by means of NIR spectroscopy (IRAS 18151-1208, LH2 = 0.7 Lˇ [6]; IRAS 11101-5829, LH2 2 Lˇ [7]) or H2 imaging (IRAS 05358+3543 [9]; IRAS 16547-4247 [2]), assuming for these last an average temperature of 2500 K. All the sources have almost the same Lbol as IRAS 20126 + 4104. The LH2 estimates are quite close to the H2 luminosity of IRAS 20126 + 4104 obtained from the warm H2 component (i. e. from our NIR analysis, LH2 = 4.6 ˙ 0.3 Lˇ ). Indeed an MIR investigation could reveal whether the cold H2 component plays a major role in the cooling of those massive jets, as well. Considering our results and the literature data of a few intermediate- and highmass protostellar jets, we conclude that these few jets appear to be scaled-up versions of their low-mass protostellar counterparts. Indeed the evolutionary outflow scenario proposed in [1] fits our findings well. In this context, those sources which have jets detected toward them are very young (well before the MS turn-on), while those without detectable jets near the protostar have ultracompact HII (UCHII) regions. If the disappearance of a collimated jet in early B protostars stems from the presence of enhanced ionising radiation from an accreting early main sequence star, then all early B stars may be formed via accretion. In this sense, they are scaledup versions of low-mass protostars at early phases (such as for IRAS 20126 + 4104, IRAS 18151-1208, IRAS 11101-5829, etc.). As YSOs evolve, developing UCHII regions which destroy the disc, the jets finally disappear, and the outflows look more like poorly collimated wind-blow bubbles as observations seem to indicate [8]. Acknowledgements The present work was supported in part by the European Community’s Marie Curie Actions-Human Resource and Mobility within the JETSET (Jet Simulations, Experiments and Theory) network under contract MRTN-CT-2004 005592. Our observations have been funded by the Optical Infrared Coordination network (OPTICON), a major international collaboration supported by the Research Infrastructures Programme of the European Commission’s Sixth Framework Programme.

References 1. Beuther, H., Shepherd, D. In: M.S.N. Kumar, M. Tafalla, P. Caselli (eds.) Cores to Clusters: Star Formation with Next Generation Telescopes, pp. 105–119 (2005) 2. Brooks, K.J., Garay, G., Mardones, D., Bronfman, L. ApJ 594, L131–L134 (2003) 3. Caratti o Garatti, A., Froebrich, D., Eisl¨offel, J. et al. A&A 485, 137–152 (2008) 4. Caratti o Garatti, A., Giannini, T., Nisini, B., Lorenzetti, D. A&A 449, 1077–1088 (2006) 5. Cesaroni, R., Neri, R., Olmi, L. et al. A&A 434, 1039–1054 (2005) 6. Davis, C.J., Varricatt, W.P., Todd, S.P., Ramsay Howat, S.K. A&A 425, 981–995 (2004) 7. Gredel, R. A&A 457, 157–166 (2006) 8. Li, J.Z., Smith, M.D., Gredel, R. et al. ApJ 679, L101–L104 (2008) 9. Porras, A., Cruz-Gonz´alez, I., Salas, L. A&A 361, 660–670 (2000) 10. Shepherd, D.S., Yu, K.C., Bally, J., Testi, L. ApJ 535, 833–846 (2000)

Part V

Jet Propagation, Stability, Interaction with the Environment, X-ray Emission

General Properties of Jets from Active Galactic Nuclei and Comparison with Protostellar Jets Silvano Massaglia

Abstract We present the general properties of the Active Galactic Nuclei (AGNs) and discuss the origin and structure of jets that are associated to a fraction of these objects. In particular, we address the problems of jet acceleration, its possible connection with accretion processes and jet composition. Moreover, we consider the long standing question about the origin of the dichotomy, both in morphology and power, in Fanaroff-Riley I and II radio galaxies suggesting that FR I jets are lighter than FR II ones and therefore more sensitive to the effects of shear-layer instabilities due to the interaction of the jet with the ambient medium. This hypothesis is supported by hydrodynamical numerical simulations in the relativistic regime and in three spatial dimensions. Finally, we make comparisons with jets from Young Stellar Objects (YSOs) and argue about the non dichotomic appearance of YSO jets in contrast with AGN jets.

S. Massaglia () Dipartimento di Fisica Generale dell’Universit`a, Via Pietro Giuria 1, 10125 Torino (Italy) e-mail: [email protected]

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1 Introduction Most of the galaxies of the Local Universe shine by effect of stellar and interstellar gas emissions, predominantly in the optical band. Typically, the emitted spectra are characterized by stellar absorption lines and emission lines by HII regions. These sources are called Normal Galaxies. A small fraction of galaxies, about 1%, do not follow this behavior. In fact, they show strong and broad emission lines, consistent with velocity dispersion of the emitting gas attaining several thousand of kilometers per second. Most remarkable, the non-thermal emission coming from a central nucleus, with size  102 pc, dominates over the thermal one coming from stars and interstellar gas and extends well beyond the optical band, from radio to gamma rays (Fig. 1). These sources are called Active Galaxies and they show Active Galactic Nuclei (AGNs) at their centers. The AGNs are not all the same, on the contrary they can be extremely different in their properties. It resulted convenient to classify AGNs according to their radio power, in fact they can be separated into two distinct classes: Radio Quiet and Radio Loud AGNs. Typically, the luminosity in the GHz band of radio loud AGNs exceeds the radio quiet ones by about three orders of magnitudes. Table 1 summarizes the objects belonging to the two classes. Observationally, Seyfert 1 show broad emission lines in the spectrum (Broad Line Region, BLR, with velocity dispersion   104 km s1 ), while for Seyfert 2 the spectral lines are narrower (Narrow Line Region, NLR,   103 km s1 ).

Fig. 1 Spectral Energy Distribution of the Centaurus A core. The continuous line is a synchrotron plus Synchrotron Self-Compton (SSC) model (Prieto et al. 2007)

Properties of Jets from AGNs Table 1 Radio quiet AGNs Seyfert I galaxies (Sey 1) (BLR,   104 km s1 ) Seyfert II galaxies (Sey 2) (NLR,  103 km s1 ) Radio quiet quasars (QSOs)

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Radio loud AGNs Radio galaxies Radio quasars BL Lac Objects Optically Violent Variables (OVVs)

This complex “zoology” has been interpreted by Urry & Padovani [23] within the so called Unified Model for AGNs. According to this picture, radio quiet AGNs have no jets (or very weak ones) while the radio loud ones have jets; the angle between the line of sight and the plane of the obscuring torus determines the observed properties of the AGNs. We will discuss the main properties of radio loud AGNs, in particular of the radio galaxies, since the presence of a jet allows comparisons with the protostellar case. In Section 2 we review the main properties of radio galaxies and in Section 3 we address the question of the origin of the dichotomy of radio source morphology and make comparisons with the case of protostellar jets.

2 The Radio Galaxies Radio galaxies are seen in the radio band emitting power-law spectra, indicative of synchrotron emission with typical spectral index ˛  0:5. Moreover, jet and hot-spots of some bright sources are seen in the optical and X-ray bands as well. The emission mechanism at these higher frequencies is again continuum and it may be synchrotron or Synchrotron-Self-Compton. Direct observations of radio galaxies give us: i) the radio luminosity that is 1039  1044 ergs s1 ; ii) the size, from a few kiloparsec to some megaparsec; iii) the morphological brightness distribution, and iv) the polarization degree of the radio emission. We can then derive, by indirect means, the main physical parameters such as the life timescale, 107  108 ys, the mean magnetic field, 10  103 G, and the kinetic power, 1042  1047 ergs s1 . The values of the jet main physical parameters, such as jet velocity, density and composition, are still under debate after many decades of investigations. The reason for these uncertainties in constraining the basic physical parameters is due to the very nature of the radiation emission which is typically non-thermal continuum, i.e. the absence of any lines in the radiation spectrum (see [15]). Therefore, the value of the magnetic field is derived by the minimum energy assumption, the jet kinetic power by the work done against the ambient to evacuate cavities for accommodating the radio lobes, the jet velocity by the radiative flux contrast of the approaching to receding jets and by proper motion observations, and the jet density by comparison between observed morphologies and the outcome of numerical simulations.

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2.1 The Fanaroff-Riley Classification Historically, the extragalactic radio sources have been classified into two categories [5] based upon their radio morphology: a first class of objects, preferentially found in rich clusters and hosted by weak-lined galaxies, shows jet-dominated emission and two-sided jets and was named FR I (Fig. 2); a second one, found isolated or in poor groups and hosted by strong emission-line galaxies, presents lobe-dominated emission and one-sided jets and was called FR II (or “classical doubles”) (Fig. 2). Besides morphology, FR I and FR II radio sources were discriminated in power as well: objects below  2  1025 h2100 W Hz1 str1 were typically referred as FR I sources. A perhaps more illuminating criterion has been found by Ledlow & Owen [13] who plotted the radio luminosity at 1:4 GHz against the optical absolute magnitude of the host galaxy: they found the bordering line of FR I to FR II regions correlating as LR / L1:7 opt (Fig. 3), i.e. in a luminous galaxy more radio power is required to form a FR II radio sources. This correlation is important since it can be interpreted as an indication that the environment may play a crucial role in determining the source structure. The above argument yields the basic question of the origin of FR I/FR II dichotomy, whether intrinsic or ambient driven.

Fig. 2 VLA radio image at 20 cm of the FR I source 3C31 (left panel); VLA radio image at 6 cm of the FR II source 3C98 (right panel)

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Fig. 3 Diagram of a sample of FRI/II objects (from [13])

2.2 The Jet Composition We know that protostellar (or Young Stellar Object, YSO) jets are made by ordinary matter, since emission lines in the observed spectra reveal the presence of Hydrogen, Helium and metals. Conversely, we do not know for certain what jets in radio galaxies are made of. They can be dominated either by electron-proton plasma, or by electron-positron pairs, or they can be Poynting-flux beams: the absence of emission lines does not allow a definite and clear-cut answer. However, if one calculates the work done by jets to inflate lobes and cocoons realizes that the electron-proton jet composition hypothesis is favored [21]. Moreover, electron-positron jets suffer of strong inverse Compton losses of the CMB photons [10]. On the other hand, since magnetic field are likely to play a fundamental role in the acceleration process, one may think that the main constituent of jets is not mass but fields in form of Poynting-dominated beams. However, Sikora et al. [22] argued that, even though jets could be Poynting-dominated at the origin, observational data imply that they become kinetically dominated beyond about 1,000 gravitational radii. Furthermore, Giannios & Spruit [7] discussed the role of kink instability in Poynting-flux dominated jets and find that the Poynting flux dissipates and the jet becomes kinetically-dominated again at 1,000 gravitational radii.

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2.3 The Jet Acceleration Analytical and numerical acceleration models for protostellar jets typically rely upon MHD processes that convert a fraction of the matter accreting through an accretion disk onto the central young star into outflowing material, traveling at superalfvenic velocities (e.g. [4, 6, 24, 20, 25, 16]) In case of AGN jets, beside the above mechanism, it was suggested [3] that energy can be extracted from the central rotating Supermassive Black Hole (SMBH) as well; but, as pointed out by Livio et al. [14], this mechanism, in order to bring a sensible contribution to the jet energy, requires unreasonably high magnetic fields threading the SMBH. Moreover, recent investigations have revealed a correlation between jet and accretion power that persists over about three decades [1, 11]. This correlation is shown in Fig. 4 [2], where appears that the conversion efficiency of accreting into jet power is   102 . One can then safely assume that the jet acceleration is governed by the accretion rate through an accretion disk in relativistic regime and that AGN jets can be Poyntingflux dominated in the sub-parsec regions, but are matter- dominated beyond.

3 AGN vs Protostellar Jets We have briefly reviewed the main properties of AGN jets and some ideas about how they work. It may be interesting to compare these properties and ideas with the case of protostellar jets (Table 2). For a review on jets from YSOs one can

Properties of Jets from AGNs Table 2 Radiation Line emission Shocks Composition FRI/FRII dichotomy

279 AGN Jets non-thermal no likely electron-proton (?) yes

YSO Jets thermal yes yes electron-proton no

refer to Reipurth & Bally [18]. As discussed before, the presence of emission lines allows a more precise, affordable and direct diagnosis of the mechanisms at work and of the values of the physical parameters. However, there is a parameter that is hardly constrained in YSO jets, that is the medium of the molecular nebula that immediately surrounds the jet. In fact, this medium is too tenuous and cool to be observed. A comparison with the phenomenology of jets in radio galaxies may help to put limits on the external medium. Looking at Table 2 we see that AGN jets are dichotomic while protostellar jets are not: YSO jets do not look like FR Is. FR I jets are non-relativistic at kiloparsec scales. But, VLBI observations of radio galaxies [8] have shown that, while FR Is and FR IIs have different radio power and kiloparsec scale morphology, they appear similar at parsec scales where the jet bulk Lorentz factor is typically  D 3  10 for both classes. Moreover, some FR I sources, e.g. M87 [12] and B2 1144 + 35 [9], observed at VLBI show limbbrightened emission in the radio band. This limb brightening can be interpreted assuming the presence in the jet of a central spine with high Lorentz factor,  D 5  10, and an outer layer with  ' 2. The synchrotron emission from the central spine would be deboosted when viewed at angles larger than about 30ı . When examining the origin of the FR I/FR II dichotomy in the extrinsic scheme, i.e. in terms of the interaction of the jet with the ambient medium, one can suppose that this interaction is mediated by shear-layer instabilities that produce momentum transfer to and entrainment of matter external to the jet. Rossi et al. [19] have studied numerically, in three spatial dimensions, the onset and nonlinear growth of unstable modes in relativistic, fluid jets while propagating into a uniform medium. The unstable perturbations caused mixing and deceleration of the jet, processes controlled by the initial jet Mach number M and ambient-to-jet density ratio , with the jet Lorentz factor parameter kept fixed to 10. They found that very light jets only ( D 104 ) were substantially decelerated, while different Mach numbers did not seem to play a significant role. Moreover, the jet deceleration was not uniform, but more effective at the outer layers, leaving a central “spine” traveling at higher Lorentz factor. In Fig. 5 one plots, for the case M D 3 and  D 104 , the amount of jet mass against the factor ˇ, with ˇ the jet speed in units of the speed of light, at two different times and at a certain longitudinal coordinate of the jet (i.e. 37.5 jet radii from the inlet). Considering the later time (solid line) one realizes that the mass is not uniformly distributed in velocity, but is segregated in a fast central spine with ˇ  5 and an outer layer with ˇ  0:2, in agreement with the observational requirements for the limb brightening.

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Fig. 5 Jet mass distribution vs ˇ for two different evolutionary times of the simulation taken at a fixed axial coordinate in the jet

But how about YSO jets? A class of objects corresponding to FRI sources does not exists: YSO jets look like FRII objects, with terminal bow-shocks and working surfaces and cocoons seen at infrared wavelengths (Bally, this issue). Protostellar jets can be as well subject to shear-layer instabilities [17] and, albeit the physical conditions are different, one may still argue that highly underdense jets would produce FRI-like morphologies, thus YSO jets may have moderate density contrasts with the ambient. Thus, ideally, one could perform a similar 3D numerical study of YSO jets considering the appropriate physical ingredients, such as molecular and atomic radiative losses, a given Mach number and varying density contrast. A limit on the density contrast could be set when a FR I-like morphology arises.

4 Summary We have reviewed the main properties of AGNs, their classification and the interpretation problems connected with the FR I/FR II dichotomy, we have then presented a possible explanation of this dichotomic behavior in terms of shear-layer instabilities and consequent mixing and entrainment of external material. We have finally suggested some possible way to constrain the ambient-to-jet density contrast for protostellar jets. Acknowledgements This contribution has been supported by EU contract MRTN-CT-2004005592 within the Marie Curie RTN JETSET project.

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References 1. S.W. Allen, R.J.H. Dunn, A.C. Fabian, G.B. Taylor, C.S. Reynolds, MNRAS 372, 21 (2006) 2. B. Balmaverde, R.D. Baldi, A. Capetti, A&A 486, 119 (2008) 3. R.D. Blandford, R.L. Znajek, MNRAS 179, 433 (1977) 4. R.D. Blandford, D.G. Payne, MNRAS 199, 883 (1982) 5. B.L. Fanaroff, J.M. Riley, MNRAS 167, 31 (1974) 6. J. Ferreira, A&A 319, 340 (1997) 7. D. Giannios, H.C. Spruit, A&A 450, 887 (2006) 8. G. Giovannini, W.D. Cotton, L. Feretti, L. Lara, T. Venturi, ApJ 552, 508 (2001) 9. G. Giovannini, M. Giroletti, G.B. Taylor, A&A 474, 409 (2007) 10. D.E. Harris, H. Krawczynski, RMxAC 27, 188 (2006) 11. S. Heinz, A. Merloni, J. Schwab, ApJ 658, L9 (2007) 12. Y.Y. Kovalev, M.L. Lister, D.C. Homan, K.I. Kellermann, ApJ 668, 27 (2007) 13. M.J. Ledlow, F.N. Owen, AJ 112, 9 (1996) 14. M. Livio, G.I. Ogilvie, J.E. Pringle, ApJ 512, 100 (1999) 15. S. Massaglia, Ap&SS 287, 223 (2003) 16. T. Matsakos, K. Tsinganos, N. Vlahakis, S. Massaglia, A. Mignone, E. Trussoni, A&A 477, 521 (2008) 17. M. Micono, G. Bodo, S. Massaglia, A. Ferrari, R. Rosner, A&A 360, 795 (2000) 18. B. Reipurth, J. Bally, ARA&A 39, 40 (2001) 19. P. Rossi, A. Mignone, G. Bodo, S. Massaglia, A. Ferrari, A&A 488, 795 (2008) 20. C. Sauty, E. Trussoni, K. Tsinganos, A&A 389, 1068 (2002) 21. F. Shankar, A. Cavaliere, M. Cirasuolo, L. Maraschi, ApJ 676, 131 (2008) 22. M. Sikora, M.C. Begelman, G.M. Madejski, J.-P. Lasota, ApJ 625, 72 (2005) 23. C.M. Urry, P. Padovani, PASP 107, 803 (1995) 24. N. Vlahakis, K. Tsinganos, C. Sauty, E. Trussoni, MNRAS 318, 417 (2000) 25. C. Zanni, A.Ferrari, R. Rosner, G. Bodo, S. Massaglia, A&A 469, 811 (2007)

The Kelvin-Helmholtz Instability in Stellar Jets Edo Trussoni

Abstract The Kelvin-Helmholtz Instability (KHI) is amongst the most relevant processes occurring in outflows propagating through an external medium. The KHI can critically affect the dynamics of jets leading to their disruption, but at same time it drives the formation of the peculiar morphologies observed in stellar collimated outflows. Moreover, the KHI can be an efficient mechanism for momentum deposition into the environment and mixing with the external medium. Here the main properties of the KHI are discussed considering in particular the effects of radiative losses and of magnetic fields, the most important ‘ingredients’ in the framework of the phenomenology of stellar jets.

1 Introduction Since the discovery of jets in radio galaxies few decades ago, the analysis of the various properties of the Kelvin-Helmholtz Instability (KHI), occurring at the interface between two fluids in relative motion, has been one of relevant fields of

E. Trussoni () INAF - Osservatorio Astronomico di Torino, Italy e-mail: [email protected] K. Tsinganos et al. (eds.), Protostellar Jets in Context, Astrophysics and Space Science Proceedings, DOI 10.1007/978-3-642-00576-3 34, c Springer-Verlag Berlin Heidelberg 2009 

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the astronomical research. In fact it is recognized that the KHI is one of the most important physical processes occurring in several contexts, whenever are present collimated flows propagating through an external environment. It is this mechanism that governs the dynamics, evolution, momentum and energy deposition and the mixing with the ambient medium of winds and jets. The theoretical study of the KHI starts from the basic fluidodynamic description of the plasma and must include some physical ‘ingredients’ for a consistent analysis of the various astrophysical scenarios: geometry, magnetic fields, radiative losses, relativistic effects, etc. Two main approaches can be followed for the analysis of the equations governing the KHI: linear and non linear. In the former one, that allows in general an analytical treatment of the problem, we can get information about the kind of perturbation that affects the equilibrium configuration and on which time scales or scale lengths they develop. However a useful comparison of the KHI features with the observed phenomena requires a non linear analysis, that can be performed through advanced numerical simulations. In several previous reviews are outlined the fundamental properties of the KHI in various astrophysical contexts (see e.g. [1, 2, 3, 4]), here we summarize how this instability may be relevant for the dynamics and evolution of stellar jets.

2 KHI: General Properties The basic equations describing the evolution of the KHI are the conservation laws of mass, momentum and energy. For magnetized flows, assuming an ideal MHD description, we must include the conservation of the magnetic flux and the induction equation. In the following of this section the main properties of the KHI are summarized in the simplest case of an adiabatic and unmagnetized jet.

2.1 Linear Analysis Perturbing a cylindrical jet, confined by an external plasma, with sinusoidal normal modes, we get the dispersion relation that provides the complex frequency ! as function of the wavenumber k and of the equilibrium parameters: the time scale of growth of unstable perturbations is ti  1=Im !. This is the so called ‘temporal’ analysis. Alternatively, we can follow a ‘spatial’ analysis, with ! real and k complex, and from the dispersion relation we get li  1=Im k, the scale length of the growth of the instability. In the linear regime the two approaches are equivalent, and can be related through the group velocity of the perturbations. In the simplest case of a jet with no density contrast across the transition layer the development of unstable modes depends on three parameters: the Mach number M D vo =vs (vo and vs are the jet and sound velocities, respectively), the longitudinal wavenumber k (for the temporal analysis) or the frequency ! (for the spatial

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analysis), and the azimuthal wavenumber n (n D 0; 1; > 1 correspond to pinching, helical and flute modes, respectively). Two classes of perturbations are unstable. Ordinary/surface modes. The interface oscillates with monotonically increasing amplitude and its development is concentrated around the surface layer. These modes are stabilized for supersonic speeds and vanishing wavelength with respect to the jet radius, and in non linear regime they evolve into steady vortical structures confined nearby the transition boundary. Reflected/body modes. These are acoustic perturbations propagating inside and outside the jet and are found only for supersonic velocities. The growth of the amplitude of these waves leads to the formation of outgoing shocks carrying away energy from the jet. These properties qualitatively hold almost independently on the azimuthal number n and whether or not a density contrast is assumed between the inner and outer medium (for details see [4]). The main implication from the astrophysical point of view is then that, differently from the classical planar ‘vortex sheet’ case, a supersonic jet always undergoes the effect of the KHI.

2.2 Non Linear Analysis The study of the non linear evolution of the KHI can be performed through numerical simulations, assuming as initial configuration a cylindrical beam in pressure equilibrium with an external medium: this structure is perturbed through a spectrum of linearly unstable modes. It is important in this context the choice of a temporal or spatial analysis. In the former case the jet is perturbed all long its length, and periodic boundary conditions are assumed on both sides of the beam. For the spatial analysis only one boundary is perturbed, where periodic conditions are assumed, while on the opposite boundary free conditions hold. The two approaches are complementary: the temporal analysis allows to follow the global jet evolution for long times. Conversely, through the spatial analysis we can study the formation of structures advected along the flow but only as long as they cross all the domain. In the framework of a temporal analysis of the KHI in 2-D supersonic beams, axisymmetric perturbations (n D 0) grow initially according to the linear theory. Thence the fastest growing modes prevail transforming into internal shocks that heat the intrashock gas and inflate the beam. These shocks evolve leading to mixing between the jet and environment, with consequent widening of the velocity layer. This final configuration is attained after  20–30 dynamical times td (defined as a=vs ) [5]. The effect the KHI on stellar jets is thence twofold: during the intermediate phases regular structures can form, that could be associated with the periodic knots observed in most of these objects. However the instability leads inevitably to the final disruption of the collimated outflow with consequent deposition of energy and momentum into the external medium. The time life of this disruption appears consistent with the observed ages of the stellar outflows.

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Fig. 1 Longitudinal maps of the density distribution at a time D 6:5td for a supersonic jet (M D 10), ten times denser than the environment for the 2-D slab (left panel) and the 3-D jet (right panel; from [7])

The above findings are however restricted to the evolution of the axisymmetric perturbation only (n D 0). As intermediate step to investigate the evolution of helical modes (n D 1) a planar 2-D slab can be considered [6]: the behaviour is similar to the axisymmetric case, but the large scale structures are less regular. In fully 3-D geometry the evolution is drastically different: namely the evolutive phases are similar, but on much faster time scales than in the 2-D cases, leading to a rapid disruption of the jets. As can be seen in Fig. 1, this disruption is enhanced in particular by the onset of small scale structures which rapidly evolve due to the developing of fluting modes (n > 1) as well as of turbulent cascades [7]. These results seem to complicate the encouraging view obtained from the 2-D simulations for modeling the stellar jets (in particular periodic knots), however a more strict comparison of the results of the instability with observations requires more realistic physical assumptions. In the following two sections we outline the properties of the KHI in radiating and magnetized jets.

3 KHI: Radiative Losses The effects of radiation losses on a system become relevant when the cooling time scale tc is comparable with or shorter than td . The value of tc is strictly related to the cooling function .T /, where T is the temperature of the gas. We discuss the main linear and non linear properties of the KHI in radiative unmagnetized jets. Linear regime. The cooling function can be approximated with  / T ˛ ; for val< 15; 000 K and ˛  1:5 for ues of T typical of stellar jets is ˛  6:5 for T 
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Fig. 2 Longitudinal (left panels) and transversal maps (right panels) of the density distribution at a time D 12td for a supersonic jet (M D 10), ten times denser than the environment for the adiabatic (up) and radiative case (down; from [11])

Non linear regime. The crucial role of the cooling function on the evolution of the KHI has been confirmed also in non linear studies. It turns out from 2-D simulations [9] that the instability can lead to the final disruption of the beam earlier or later, with respect to the adiabatic case, depending on the assumed functional form of .T /. A more realistic cooling function, including line emission and non equilibrium conditions, was adopted in [10] to analyze the stability of a supersonic axisymmetric jet. The long term evolution predicts, similarly to the adiabatic case, the formation of quite regularly spaced knots, with velocities and spectroscopical properties consistent with the observational data. More complex is however the behaviour found in 3-D beams [11], in which case there is no evidence of growing of regular structures along the flow. The main difference with the adiabatic case is that radiative emission has always a damping effect on the KHI, in particular in jets denser than the environment (see Fig. 2). Accordingly, the entrainment process and the disruption of the jet are delayed. Moreover, the classical thermal instability does not develop for values of physical quantities typical for stellar jets. A major question is whether molecular outflows can be driven by jets through the KHI [12]. The typical size of these poorly collimated outflows are  0.05–5 pc and their age is  1,000–100,000 years. It turns out from the simulations that almost all the momentum of the outflow is transferred to the external medium in  5,000– 10,000 years, with the mixing layer expanding almost linearly at a speed  10–20 % of the sound speed. On time scales of  105 years the transversal size of the outflow is  0.1–0.2 pc, i.e. of the order of the smallest observed molecular outflows. A correlation between mass and velocity is generally observed in these structures, that can be approximated with a power law, M.v/ / v , with larger  (steeper spectrum) for larger velocities. These properties can be reproduced in the simulations

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for jets with density lower than or comparable with the external medium. Simulations also show that the momentum transfer from the jet to the environment leads to the formation and acceleration of small portions of external material to high speeds > 40 km s1 ) as observed in several objects. (

4 KHI: Magnetized jets Magnetocentrifugal processes play a relevant role in launching and collimating outflows, accordingly magnetic fields must be included in the study of the KHI. The stability of magnetized beams depends critically on the assumed geometry of the field, namely from the predominance of the longitudinal or of the toroidal component. We discuss separately these two configurations. Longitudinal magnetic field. It is well known the stabilizing effect on the KHI of a longitudinal magnetic field: a planar, strongly magnetized vortex sheet across isodensity fluids is always stable. In particular in cylindrical, highly supersonic and super-alfv´enic magnetized jets the growth of the ordinary mode is damped while the onset of the reflected modes is shifted to small wavelengths [4]. This trend has been basically confirmed in the non linear regime, as found in 3-D simulation for transonic and trans-alfv´enic beams [3, 13]: in trans-alfv´enic jets the stability develops only in the super-alfv´enic region. Conversely a super-magnetosonic beam (Mf  3) shows a fast turbulent cascade towards small scale structures, as observed in the unmagnetized case. In other words, longitudinal magnetic fields do not modify the main properties of the hydrodynamic KHI for highly super-magnetosonic jets. Toroidal magnetic field. Much more complicated is the stability analysis when in the equilibrium configuration is B ¤ 0. Critical is the choice of the profiles for the magnetic components and of the magnetic pitch (Pi D rBz =B ), a measure of the prevailing of the toroidal or of the longitudinal field. Moreover, the presence of the axial electric current can drive the well known current driven instability (CDI). Several studies have been performed with various equilibria for the magnetic field (purely toroidal, force free, etc...). It was shown in [14] that in super-magnetosonic force free, cold jets with Bessel-like profile for the field components, perturbations are strongly damped with respect to jets with vanishing B . An additional effect is the splitting of the helical modes into two different instability branches for n D 1 or n D 1. The CDI also originates but, for the assumed structure, the growth rate of unstable modes is smaller than the KHI unless the flow velocity is sub-magnetosonic. Other force free configurations have been considered in [15] where it has been shown the important role of the magnetic pitch for the CDI properties. It turns out that the growth rate of the helical modes become relevant when B prevails, but the axis displacement remains concentrated around the axis. More general geometries were studied in [16] for trans- and super-magnetosonic jets perturbed by helical modes. It has been confirmed that in such conditions the toroidal field tends to decrease the growth of the ordinary and reflected

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Fig. 3 Spatial growth rate vs the wavenumber of the unstable helical modes in a supermagnetosonic jet (Mf D 3) for a force free, high pitch (left panel) and a not force free, low pitch equilibrium (right panel). The solid curve ‘CD’ represents the current unstable n D 1 helical mode; the other curves ‘KH(SM)’ refer to the unstable surface helical perturbation with n D 1 (dashed) and n D C1 (dash-dotted). The unstable body modes are not plotted (from [16])

perturbations, even though at a lower rate than for cold, force free jets. In some cases however the CDI prevails over the KHI: this occurs for not force free equilibria and high enough values of the pitch. This is shown in Fig. 3 where are plotted the spatial growth rates of the two helical modes for a super-magnetosonic jet. This trend basically holds for mildly and highly super-magnetosonic outflows, while it critically depends on the value of the relative pitch (see Fig. 6 in [16]). Three dimensional MHD simulations require advanced and sophisticated numerical codes, large computing facilities and very long runs, thence the non linear evolution of jets with helicoidal magnetic fields could be analyzed for few specific configurations. It has been confirmed that, as effect of the hoop stress, a trans-magnetosonic jet with helicoidal field close to equipartition is more stable than with a simple axial magnetic field [18]. The non linear interplay between the KHI and the CDI has been analyzed in [17] imposing helical perturbations on a trans-magnetosonic jets. In Fig. 4 are shown the results of the evolution of two configurations with a longitudinal magnetic field only and including a toroidal component, assuming values of the parameters that, in the linear regime, predict the prevailing of the CDI over the KHI (right panel in Fig. 3). In the former case we see the expected formation of vortices that, after  10–15 td , decay into smaller structures in agreement with the general KHI properties in 3-D beams. A very different evolution undergoes a jet with a toroidal field. At the beginning the two instabilities evolve independently: the KHI at the beam surface (vortical structures) and the CDI around the axis (enhancement of the toroidal field). When the two unstable regions interact the increase of the azimuthal component strengthens the hoop stress effect on the beam surface, counteracting the evolution of the KHI towards small eddies. The quite surprising conclusion is that the CDI has a stabilizing effect on the KHI, accordingly jets with toroidal fields evolve to more ordered configurations, as evident in Fig. 4.

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Fig. 4 Maps of the density distribution for a slightly super-magnetosonic jet (Mf D 1:24) with an axial (upper panels) and helicoidal magnetic field (lower panels). The left and central maps show longitudinal sections at times D 8:5td and 10.5 td (up), and D 10:5td and 12.5 td (down); the right maps refer to transversal cuts at timesD 14td for both cases (from [17])

5 Conclusions From the results discussed in previous sections it turns out that the typical times of evolution of the KHI, when radiative losses and helicoidal magnetic fields are considered, may last a few tens of dynamical times before complete disruption. Then, expressing td  700 a16 =vs;6 years (a16 and vs;6 are the jet radius and the sound speed in units of 1016 cm and 106 cm s1 , respectively), jets can survive as basically coherent collimated outflows for a non negligible fraction of their time life. During this phase, through the KHI several dynamical effects occur: formation of shocks with plasma heating, formation of complex structures, mixing and entrainment of external matter, momentum deposition and driving of molecular outflows. These promising results encourage further effort to extend the analysis of KHI, in particular a deeper parametric study, or the evolution of a highly supermagnetosonic, radiating jet confined by a toroidal magnetic field. Moreover, there are more physical aspects of the instability process that deserve further investigation. It has been pointed out very recently the relevant role, in 2-D planar trans-magnetosonic slabs, of the structure of the transition layers between the inner medium and the environment. In particular, different thickness of the layers, when compared with the slab transverse size, can heavily affect the evolution of the outflow [19]. Even more critical effects may occur on the non linear properties of the KHI when the physical resistivity is included. In has been shown that in a transition layer between two planar fluids reconnection and resistive dissipation strongly rule the disruption of vortices and the turbulent cascade towards small scale eddies [20].

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We conclude pointing out a couple of warnings concerning the comparison of the theoretical results with the data. The 2-D and 3-D results are very different concerning the evolutive time scales, morphological structures and final configurations. Accordingly, even though 2-D simulations are unavoidable and very useful intermediate steps to study the various properties of the KHI, the observations should be interpreted mainly through 3-D calculations. Finally, the choice of spatial or temporal analysis of the KHI is not obvious. It has been shown that following a spatial analysis the long term evolution of a 2-D trans-magnetosonic slab may result much longer than predicted by the temporal study [21]. This is the result of the enhancement of the magnetic field and the consequent recollimation at longer times, that only through spatial analysis can emerge.

References 1. Birkinshaw, M.: The stability of jets. In Beams and Jets in Astrophysics, P.A. Hughes edt. (Cambridge Univ. Press, Cambridge 1991), pp. 278–341 2. Ferrari, A.: Modeling Extragalactic Jets. ARA&A, 36, 539–598 (1998) 3. Hardee, P.E.: The stability properties of astrophysical jets. ApSS, 293, 117–129, (2004) 4. Trussoni, E.: The Kelvin-Helmholtz Instability. In Jets From Young Stars III - Numerical MHD and Instabilities, S. Massaglia et al. eds., LNP 754 (Springer, Berlin Heidelberg 2008), pp. 105–130 5. Bodo, G., Massaglia, S., Ferrari, A., Trussoni, E.: Kelvin-Helmholtz instability of hydrodynamic supersonic jets. A&A, 283, 655–676, (1994) 6. Bodo, G., Massaglia, S., Rossi, P., Rosner, R., Malagoli, A., Ferrari, A.: The long-term evolution and mixing properties of high Mach number hydrodynamic jets. A&A, 303, 281–298 (1995) 7. Bodo, G., Rossi, P., Massaglia, S., Ferrari, A., Malagoli, A., Rosner, R.: Three-dimensional simulations of jets. A&A, 333, 1117–1129 (1998) 8. Massaglia, S., Trussoni, E., Bodo, G., Rossi, P., Ferrari, A.: Radiative unstable modes in the jets of young stellar objects. A&A, 260, 243–249 (1992) 9. Hardee, P.E., Stone, J.M.: The stability of radiatively cooling jets. II. Non linear evolution. ApJ, 483, 136–147 (1997) 10. Micono, M., Massaglia, S., Bodo, G., Rossi, P., Ferrari, A.: Kelvin-Helmholtz instabilities in stellar jets. IV. On the origin of the emission knots. A&A, 333, 1001–1006 (1998) 11. Micono, M., Bodo, G., Massaglia, S., Rossi, P., Ferrari, A., Rosner, R.: Kelvin-Helmholtz instability in three dimensional radiative jets. A&A, 360, 795–808 (2000) 12. Micono, M., Bodo, G., Massaglia, S., Rossi, P., Ferrari, A.: On the matter entrainment by stellar jets and the acceleration of molecular outflows. A&A, 364, 318–326 (2000) 13. Hardee, P.E., Rosen, A.: On the dynamics and structure of three-dimensional trans-alfv´enic jets. ApJ, 524, 650–666 (1999) 14. Appl, S.: Instabilities in trans-magnetosonic jets. A&A, 314, 995–1002 (1996) 15. Appl, S., Lery, T., Baty, H.: Current driven instabilities in astrophysical jets. A&A, 355, 818– 828 (2000) 16. Baty, H.: On the magnetohydrodynamic stability of current-carrying jets. A&A, 430, 9–17 (2005) 17. Baty, H., Keppens, R.: Interplay between Kelvin-Helmholtz and current-driven instabilities in jets. ApJ, 580, 800–814 (2002) 18. Rosen, A., Hardee, P.E., Clarke, D.A., Johnson, A.: Effects of Magnetic Fields on Mass Entrainment of Super-magnetosonic Jets. ApJ, 510, 136–154 (1999)

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19. Baty, H., Keppens, R.: Kelvin-Helmholtz disruptions in extended magnetized jet flows. A&A, 447, 9–22 (2006) 20. Palotti, M.L., Heitsch, F., Zweibel, E.G., Huang, Y.-M.: Evolution of unmagnetized and magnetized shear flows. ApJ, 678, 234–244 (2008) 21. Viallet, M., Baty, H.: Spatial simulations of the Kelvin-Helmholtz instability in astrophysical jets. A partial stabilization mechanism for weakly magnetized transonic flows. A&A, 473, 1–9 (2007)

Radiative Jets from Variable Sources Alejandro C. Raga, Jorge Cant´o, Fabio De Colle, Alejandro Esquivel, Primoz Kajdic, Ary Rodr´ıguez- Gonz´alez, and Pablo F. Vel´azquez

Abstract This paper gives a short review of the literature on radiative jets from variable sources, giving a more or less complete and up to date bibliography. We then present a number of numerical simulations exploring the effect of different forms of the ejection velocity variability and the effect of having a non-top hat ejection velocity cross section. Finally, we explore the effect of a precession of the ejection axis.

A.C. Raga (), A. Esquivel, A. Rodr´ıguez-Gonz´alez, and P.F. Vel´azquez ICN, Universidad Nacional Aut´onoma de M´exico, Ap. 75-543, 04510 D. F., M´exico e-mail: [email protected]; [email protected]; [email protected]; [email protected] J. Cant´o IA, Universidad Nacional Aut´onoma de M´exico, Ap. 75-264, 04510 D. F., M´exico e-mail: [email protected] F. De Colle DIAS, 31 Fitzwilliam Place, Dublin 2, Ireland e-mail: [email protected] P. Kajdic IG, Universidad Nacional Aut´onoma de M´exico, 04510 D. F., M´exico e-mail: [email protected]

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1 Introduction The general properties of variable jets have been described using analytic techniques based on a “free streaming” approximation for the continuous segments of the flow. This solution is then coupled with either a “ram pressure balance” or a “momentum conservation” equation for the internal working surfaces. A number of papers describe a variety of analytic solutions [44,52,27,57,41,68,15,56,46,5,37,6,47,36]. The problem of a variable, radiative jet has also been studied with 1D, timedependent numerical simulations [44, 23, 68, 73, 46, 38], 2D slab-symmetric simulations [71, 42, 2, 45, 48], and axisymmetric simulatons [13, 14, 1, 3, 40, 73, 56, 70, 16, 4, 28, 37, 60, 75, 55, 25, 26, 61, 10, 72, 50, 11]. Finally, radiative, variable jet models have also been computed with 3D simulations [71,18,19,20,69,73,53,49,29,30,17,7,8,43,33,34,31,32,37,49,39,62,63,64, 22,58,74,55,59,65,66,9,67,24,21,35]. These papers have explored (in a somewhat chaotic way) the effects of multi-period ejection variabilities, with strong emphasis on variabilities of the ejection velocity. Ejections that also have density variabilities have been studied (see, e.g., [23, 46, 38]), but have received relatively little attention in the literature. The effect of introducing a precession of the outflow axis has been well studied (see, e.g., [63]). Also, the effect on variable jets of the presence of a magnetic field has also been studied [16, 7, 8, 10, 11]. In this paper, we present several simulations of variable jets, which are meant as an illustration of the different flow configurations that can be obtained. These simulations are similar to many results present in the literature, having the only advantage of sharing a somewhat more homogeneous parameter set. They explore jets with single- and double-mode sinusoidal ejection velocity variabilities, with non-sinusoidal variabilities, with a precessing ejection direction, and with top-hat and axially peaked ejection velocity cross sections.

2 The Numerical Simulations 2.1 The Basic Setup The simulations were computed with 2D (axisymmetric) and 3D versions of the yguaz´u-a code [51]. The version of the code which was used integrates the gasdynamic equations together with a single rate equation for neutral H, and uses the temperature, density and H ionisation fraction to compute a parametrised cooling function [58]. A seed electron density (assumed to arise from singly ionised C) is assumed, and the initial configuration of the flow (both for the jet and the environment) is otherwise neutral. In all cases, a 6-level, binary adaptive grid was used.

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2.2 Single- and two-mode Sinusoidal Ejection Velocity Variability We present a series of 2D (axisymmetric) simulations of radiative jets. The jets move into a homogeneous environment (with nenv D 100 cm3 , Tenv D 100 K), and have initially top-hat and time-independent density and temperature cross sections (with nj D 1000 cm3 and Tj D 1000 K, respectively). The jet velocity, however, is allowed to have a sinusoidal time variability (with an on-axis mean velocity v0 , half-amplitude v and period ), and a quadratic, centreto-edge dependence with a .vc =ve / velocity contrast. The velocity with which the jet is ejected is therefore given by :    2 t  vj .r; t / D v0 C v sin "

   2 # ve r 1 1 : vc rj

(1)

Though in some of the models we include two modes with different amplitudes (v) and periods ( ), in (1) we show only one such mode. The jet has an initial radius rj D 7  1015 cm, which is resolved with 18 grid points (of the highest resolution grid). All of the models have v0 D 200 km s1 , and the other parameters of the models are given in Table 1. We should note that model M2 has a ˛ D 5ı initial opening half-angle, while the other models have an initially perfectly collimated jet. Figure 1 shows density stratifications obtained from these models. The basic model is M1, with a perfectly collimated, top-hat, single-mode ejection. The other models show that :  having a non-zero initial opening angle (model M2, see Table 1 and Fig. 1) pro-

duces internal working surfaces which broaden and become less dense as they travel away from the source,  introducing a vc =ve D 2 centre-to-limb velocity ratio in the ejection cross section has a strong effect, leading to the formation of narrow, conically shaped working surfaces,

Table 1 Axisymmetric models of jets with sinusoidal variabilitiesa Model v vc =ve [yr] [km s1 ] M1 30 30 1.0 M2 30 30 1.0 M3 30 30 2.0 M4 30/500 30/100 1.0 M5 30/500 30/100 1.11 a

˛ [ı ] 0 5 0 0 0

models shown in Fig. 1 (the remaining parameters are described in Sect. 2.2)

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Fig. 1 Models with single or double-mode sinusoidal ejection variabilities. The density stratifications obtained for a t D 400 (left) and a t D 800 yr integration time (right) are shown. The frames have a horizontal extent of 4  1017 cm. The density stratifications are depitced (in g cm3 ) with the logarithmic colour scale given by the bar on the bottom right. The parameters of the model M1 (top) through M5 (bottom) are described in Sect. 2.2 (also see Table 1). A color version of this plot can be found in the electronic version of this book and in Appendix A (Fig. A.19)

 the two-mode models (M4 and M5) show the characteristic piling up of short-

period mode working surfaces into larger working surfaces (due to the second, longer-period mode),  in model M5, we see that introducing a shallow, vc =ve D 1:11 centre-to-limb velocity ratio, has an important effect on the working surfaces resulting from the two ejection velocity variability modes.

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2.3 Non-sinusoidal Ejection Velocity Variability In this section, we present time-dependent ejection velocity models with different functional forms for the variabilities. We first define the functions :   2 t 0 ; (2) fc .t / D cos fl .t / D 0:142 Œ1 C fc .t / e 10t fr .t / D 0:142 Œ1 C fc .t / e

0 =

10t 0 =

 1;

(3)

 1;

(4)

0

where t D t  .int Œt =  C 1=2/ . We then compute the ejection velocity as vj .t / D v0 C v f .t / ;

(5)

setting f .t / D fc .t /, fl .t / and fr .t /. The “sine wave”, “skewed left” and “skewed right” variabilities (fc , fl and fr , respectively) are shown in Fig. 2. With these three forms for the ejection velocity variability, we run axisymmetric simulations of initially top hat jets with v0 D 200 km s1 , D 100 yr, v D 50 km s1 , nj D 1000 cm3 , Tj D 1000 K and rj D 7  1015 cm (the jet radius being resolved with 18 grid points at the maximum resolution of the adaptive grid) traveling into a uniform environment of nenv D 100 cm3 and Tenv D 100 K. Some of the density stratifications resulting from these numerical simulations are shown in Fig. 3. From this Figure, we see that the structure of the internal working surfaces depends only weakly on the functional form of the ejection velocity variability. This is

Fig. 2 Sinusoidal (fc ), “skewed left” (fl ) and “skewed right” (fr ) forms for the ejection velocity variability (see Sect. 2.3)

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Fig. 3 Density stratifications (corresponding to a t D 800 yr integration time) resulting from simulations with the ejection velocity time-variabilities shown in Fig. 2 (also see Sect. 2.3). The displayed domains have an axial extent of 4  1017 cm. The density stratifications are displayed (in g cm3 ) with the colour scheme given by the bar on the bottom right. A color version of this plot can be found in the electronic version of this book and in Appendix A (Fig. A.20)

a result of the fact that the internal working surfaces rapidly approach the asymptotic regime of large distances from the source [52], in which the flow does not depend on the details of the ejection velocity history.

2.4 A 3D Simulation of a Variable, Precessing jet We finalize our presentation of numerical simulations with a 3D model of a jet with a two-mode sinusoidal variability (with periods of 30 and 230 yr and corresponding half-amplitudes of 20 and 100 km s1 ) and a mean velocity of 200 km s1 . The direction of ejection precesses with a period of 1000 yr on a cone with a half-opening angle of 5ı . The jet has an initial top hat cross section of density nj D 1000 cm3 , temperature Tj D 1000 K and rj D 1016 cm, moving into a uniform environment with nenv D 100 cm3 and Tenv D 100 K. The initial jet radius is resolved with 10 grid points at the maximum resolution of the adaptive grid. Figure 4 shows a time sequence of column density maps, computed by integrating the atom+ion number density along one of the axes perpendicular to the precession axis. This figure illustrates the complexity of the resulting flow, in which

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Fig. 4 Column density time-sequence computed from the variable+precessing jet model described in Sect. 2.4. The horizontal extent of the frames is of 4  1017 cm. The column densities are given (in cm2 ) by the bar on the top. A color version of this plot can be found in the electronic version of this book and in Appendix A (Fig. A.21)

fast moving condensations (i.e., working surfaces) travel within the cavity left behind by the leading bow shock, and eventually impact the dense cavity walls in the region close to the bow shock head, or along the inside of the bow shock wings.

3 Conclusions We have described a set of variable ejection velocity, radiative jet simulations. These simulations explore the effects of having different forms of the ejection variability (single- and double-mode sinusoidal, and non-sinusoidal), the effects of having a non-top hat ejection velocity cross section, and the effect of introducing a precession of the ejection axis. The main ingredients with which one can play in order to obtain radically different jet flow configurations are of course the periods and amplitudes of the modes of the ejection velocity variability (aperiodic variabilities have been explored by [41], and by Adam Frank in these Proceedings). As illustrated by the simulations presented here, the precise functional form of the velocity variability (in between the “outflow events”) is only of secondary importance. Also, one finds that introducing a centre-to-edge, non-top hat injection velocity profile has a strong effect on the shape of the individual working surfaces. This is an interesting ingredient for attempting to model knots along specific HH jets.

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A precession in the ejection direction also can have important effects, particularly at larger distances from the outflow source. Finally, we should mention that magnetic fields can have a strong effect on the morphology of the individual working surfaces travelling down a variable jet. This point has not been explored in the present paper. Acknowledgements We acknowledge from the CONACyT grants 46828-F and 61547, and from the “Macroproyecto de Tecnolog´ıas para la Universidad de la Informaci´on y la Computaci´on” (Secretar´ıa de Desarrollo Institucional de la UNAM).

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Position-Velocity Analysis of HH 111: Physical Structure and Dust Content Linda Podio, Silvia Medves, Francesca Bacciotti, Jochen Eisl¨offel, and Tom Ray

Abstract Spectral diagnostics techniques have proved to be a powerful tool in deriving from the observations a wealth of information on the jet structure and dynamics and, hence, constraints for the theoretical models for the jet launching and propagation. From the analysis of the optical forbidden lines detected on the EFOSC2 spectra of HH 111 we derived the gas physical conditions and the depletion of calcium both as a function of the distance from the source and in different velocity channels. The performed analysis highlighted that the density, the temperature

L. Podio () and T. Ray School of Cosmic Physics, Dublin Institute for Advanced Studies, 31 Fitzwilliam Place, Dublin 2, Ireland e-mail: [email protected]; [email protected] S. Medves Universit`a di Pisa, Dipartimento di Fisica, Largo B. Pontecorvo 3, 56127 Pisa, Italy e-mail: [email protected] F. Bacciotti INAF - Osservatorio Astrofisico di Arcetri, L.go E. Fermi 5, 50125 Firenze, Italy e-mail: [email protected] J. Eisl¨offel Th¨uringer Landessternwarte Tautenburg, Sternwarte 5, D-07778 Tautenburg, Germany e-mail: [email protected] K. Tsinganos et al. (eds.), Protostellar Jets in Context, Astrophysics and Space Science Proceedings, DOI 10.1007/978-3-642-00576-3 36, c Springer-Verlag Berlin Heidelberg 2009 

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and the ionization fraction peak at the position of the strongest shocks and at high velocities, in agreement with theoretical predictions for the shock structure. Interestingly also the gas-phase abundance of Calcium is maximum in the strongest shocks confirming theoretical predictions for dust reprocessing in shocks.

1 Introduction Although it is widely accepted that jets play an essential role in star formation, there are still many open questions about how they are generated and they propagate in the interstellar medium. The development of spectral diagnostics techniques (e.g., [1]) and their application to different dataset allowed derivation of the physical and dynamical properties of the the jet thus addressing some of these questions (e.g., [14, 8]). Although it has been shown that there exists different velocity components in the jet [9] only a few studies investigated the variation of the jet properties with the gas velocity [2, 12, 3, 5]. To further investigate this aspect we apply the so-called BE technique [1] to partially velocity-resolved EFOSC2 spectra of the well known HH 111 jet [15]. In Sect. 2 we briefly explain the technique, we present the derived jet physical properties and we discuss their dependence on gas velocity. In Sect. 3 we use these results to estimate the presence of dust grains in the jet. As we will show, these kind of estimates are very important in order to understand dust reprocessing in shocks and to derive constraints for the jet launching mechanism.

2 The Jet Physical Structure in Position and Velocity In this section we present the results obtained by applying spectral diagnostics (i.e. the BE technique) to partially velocity-resolved spectra of HH 111. (The data were acquired at the 3.6 m ESO telescope equipped with the optical spectrograph EFOSC2). Our analysis relies on the dependence of the ratios between optical forbidden lines on the gas physical conditions in the jet, i.e. on the electron and total density, ne and nH , the ionization fraction, xe , and the temperature, Te . As a first step from the spectral images of the detected SC , O0 , and NC lines we computed the line ratios pixel by pixel. The obtained position-velocity (PV) diagrams (see Fig. 1 in the book colour section) give a qualitative idea of the variations of the gas physical conditions along the jet and with the gas velocity. The ratio between the sulphur doublet lines ([SII]6731/6716) is a tracer of the gas electron density and its PV shows that ne peaks at the positions of the brightest knots and at high velocities. The [NII]/[OI] and [OI]/[SII] ratios, on the other hand, increase mainly for increasing xe and Te . The ratio PVs indicate a peak of the excitation (xe , Te ) in the high velocity component of knot L. Interestingly, HST high-angular resolution images of the jet clearly show that HH 111 is a chain of shock working surfaces where knot L has a clear “bow” morphology [15]. Our PV diagrams confirm this picture highlighting the shock velocity structure.

Position-Velocity Analysis of HH 111

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Fig. 1 Position-Velocity diagrams of the forbidden lines ratios used in the diagnostics technique. From top to bottom panel: the [SII]6731 line intensity; the [SII]6731/6716 ratio increases with the electron density, ne ; the [NII]/[OI] ratio is mainly dependent on the ionization fraction, xe , and increases for increasing xe ; the [OI]/[SII] ratio increases for increasing temperature Te . The PV diagrams of the ratios show that ne peaks in the HVC of the most luminous knots (F, H, J and L) and a clear peak of xe and Te in the HVC of knot L. These are typical shock signatures. A color version of this plot can be found in the electronic version of this book and in Appendix A (Fig. A.22)

To obtain a quantitative estimate of the gas physical properties we compared the observed ratios with the theoretical ones, predicted by using the BE technique. Since the data angular/spectral resolution are much lower than the spatial/spectral sampling (FWHMseei ng 1.200 , v 300 km s1 ; spatial/spectral scale: 0.15700 /pix, ˚ 0.872 A/pix) we integrated the line fluxes in space and velocity before applying the diagnostic. The integration in space is over the seeing FWHM, and in velocity over two spectrally resolved intervals: the Low and the High Velocity Component (LVC and HVC). The variations of the physical parameters both along the jet and in the two velocity components is shown in Fig. 2. The derived parameters are in agreement with previous estimates [14]: ne varies between 102 and 2.2 103 cm3 , xe goes from a few percent to 0.4, Te 0.9–1.4 104 K, and nH ranges between 2 103 and 3 104 cm3 . With respect to previous analyses, the increased spatial sampling and the velocity information highlight several shocks features, such as ne maxima in the HVC of the brightest knots, higher values of Te in the HVC along the full jet length, and, finally, a strong and steep increase of xe and Te in the HVC of knot L, which shows a bow morphology [15] and a high shock velocity [7]. These results are an important observational validation of the shock velocity structure predicted by models [6].

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Fig. 2 From the forbidden line ratio PVs shown in Fig. 1 and by using the BE technique [1] we derived the physical parameters along the jet and in two velocity components: the High Velocity Component (HVC, solid line) and the Low Velocity Component (LVC, dotted line). From top to bottom panel: intensity profiles of the optical lines, the electron density, ne , in units of 103 cm3 , the ionisation fraction, xe , the temperature, Te , in units of 104 K, and the total density, nH , in units of 104 cm3 . Note that the HVC is denser and warmer along the full jet beam, with ne peaking at the position of the most brilliant knots. Moreover xe and Te increase steeply in the HVC of the strong, bow-shaped knot L [7, 15]

3 The Presence of Dust in the Jet Refractory species, such as calcium, are often depleted in the interstellar medium because their atoms are locked onto dust grains. On the other hand, the passage of shocks can destroy the dust releasing the refractory atoms into the gas cloud (e.g., [4]). Thus an estimate of the depletion of refractory species can be used to gauge the amount of dust grains in the jet beams. To this aim we compared observed and expected ratios between emission lines of refractory (Ca) and non-refractory (S) species. The expected ratios are calculated through the derived parameters and assuming solar abundances [13, 14]. Figure 3 shows that the observed [CaII]/[SII] ratios are always lower than the predicted ones, i.e. calcium is depleted with respect to its solar abundance (see bottom panel for a quantitative estimate of the Ca depletion). Interestingly, the Ca depletion, and hence

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Fig. 3 Top and middle panel: Comparison between observed (diamonds) and predicted (dashed lines) [CaII].7290C7324/= [SII].6716C6731/ ratios along the jet and in two velocity channel: the HVC (top panel) and the LVC (middle panel). The distance between observed and predicted [CaII]/[SII] ratios indicates a depletion of the Ca gas-phase abundance which is quantified in the bottom panel. Bottom panel: The Ca gas-phase abundance ([Ca]gas /[Ca]solar ) varies along the jet going from 0.25 to 0.5 and is higher in the HVC of the knots E, G and L. This means that the shocks along the jet are not strong enough to completely destroy dust grains and that the efficiency of this process is tightly related to the shock velocity

the dust content in the jet, is lower at higher velocities and is minimized in the HVC of the bow-shaped knot L (see HST images from [15]) which is also the one with the higher shock velocity [7]. This result shows that the “slow” shocks occurring along the jet beam (vs 20–50 km s1 , [7]) are only partially destroying the dust grains, and that the efficiency of dust reprocessing in jets shocks strongly depends on the shock velocity, in agreement with the theoretical models [4]. Moreover the depletion of calcium tends to decrease with the distance from the source in both the LVC and the HVC ([Ca]gas /[Ca]solar varies between 0.25 and 0.5 along the jet). This can only be explained if the dust in the beam comes from the disk and is gradually destroyed by the working surfaces as they propagate outwards. Note that we are not observing ambient dust since it is easily destroyed by the passage of the leading bow-shock which has large shock velocity (80 km s1 , [7]) and the dust reformation timescale is much larger than the dynamical time of the jet (103 –104 yr). This finding can put severe constraints on the size of the disk

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region from which the jet is launched. The accretion disk, in fact, is populated by dust grains only beyond the so-called dust evaporation radius, Revp [10]. Thus accurate estimates of Revp in jet sources will allow us to test the validity of the models proposed to explain the jet origin [11, 16]. Acknowledgements This work was partially supported by the European Community’s Marie Curie Research and Training Network JETSET under contract MRTN-CT-2004-005592.

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Application of Tomographic Techniques to Stellar Jets Fabio De Colle, Carlos del Burgo, and Alejandro C. Raga

Abstract We show how standard tomographic techniques may be applied to study the three dimensional structure of stellar jets. To achieve this, we invert observed emission maps, and determine the physical parameters from the line ratios. For the first time this technique is applied to the HH30 jet, showing that the jet is denser and more fragmented than what can be inferred when deducing the physical parameters directly from the observed emission maps.

1 Introduction Most of our understanding about the physical conditions of the stellar jets plasma comes from the use of emission line ratios. In particular, a standard approach to the problem is to determine the electron density ne from the ŒSIIœ6716=œ6731 ratio

F. De Colle () and C. del Burgo Dublin Institute for Advanced Studies, 31 Fitzwilliam Place, Dublin 2, Ireland e-mail: [email protected]; [email protected] A.C. Raga Instituto de Ciencias Nucleares, Universidad Nacional Aut´onoma de M´exico, Ap.P. 70543, 04510 DF e-mail: [email protected] K. Tsinganos et al. (eds.), Protostellar Jets in Context, Astrophysics and Space Science Proceedings, DOI 10.1007/978-3-642-00576-3 37, c Springer-Verlag Berlin Heidelberg 2009 

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(e.g. [6]) and the electron temperature Te from ratios of emission lines with different excitation temperatures (e.g. ŒOIœ5577=œ6300). De Colle and del Burgo [4] showed that the “usual” assumption of homogeneity (along the line of sight) of the jet plasma, leads to systematic errors in the determination of the physical parameters across the jet. This paper is devoted to show how is possible to determine the structure of a jet both along the line of sight and across the jet. Section 2 introduces the necessary techniques that can be used to determine the three dimensional jet structure. Section 3 presents an application of these techniques to the HH30 jet. Finally, Sect. 4 summarises the results obtained.

2 Method to Determine the Physical Parameters Our strategy for determining the jet structure from the observed emission line maps consists of two steps. 1. We invert the observed emission line maps to determine the volumetric emission coefficient maps as a function of the radial distance from the jet axis 2. From the emission coefficient ratios we determine the physical parameters ne , Te and xH in each radial position. In the following, we describe briefly the two steps. For a complete description see [4].

2.1 Emission Line Map Inversion The determination of the emission coefficient i.r/ from an observed emission profile I.x/ can be reduced to the inversion of the Abel transform [4]. Z

R

I.x/ D 2 x

i.r/rdr ; p r 2  x2

(1)

where R is the jet radius, r is the cylindrical distance from the jet axis and x is the distance from the projection of the jet axis on the plane of the sky. Several methods are available to invert the Abel transform (e.g., [3]). In this work, we use the multi-Gaussian approximation [2]. This method consists of a fit to the observed data with a sum of Gaussian: I.x/ D

n X

ai e x

2 = 2 i

;

(2)

iD1

The solution of the Abel inversion is given in an analytical form once the fitting parameters are known, and can be writte as

Application of Tomographic Techniques to Stellar Jets

i.r/ D

n X 1 ai r 2 = 2 i e p  i iD1

313

(3)

This method is very efficient in the parametrization of monotonic decreasing of noisy data. In addition, it is straightforward to include the point spread function (PSF) in the deconvolution, leading to an analytical expression for the inverted solution. The main hypothesis used in the emission line inversion is that the jet is assumed to be axisymmetric. Actually, this represents a generalisation of the “homogeneous” medium hypothesis, and is justified for some observed jets (e.g., HH30). Also, the images are not deconvolved with the PSF before inverting them. In addition, the multi-Gaussian method, as mentioned before, works very well for monotonic data sets, but in general it is not able to predict large variations in the solution, e.g., a cavity at the center of the jet as predicted by disk-wind models. Both aspects will be considered, using more general regularisation techniques, in future work (De Colle et al. in preparation). Figure 1 is an example of the application of the described technique. Figure 1 shows the observed data, the multigaussian fit (top) and the reconstructed data (bottom). The ratios calculated using the reconstructed data are different with respect to the original ratios (see for example the [SII] œ6716 curves with respect to the [NII] in the upper and lower panels of Fig. 1).

Fig. 1 Upper panel: Observed line emissivities (points) against multi-Gaussian fit as a function of the position across the jet. Lower panel: Observed line ratios, fit and reconstructed line ratios. Figure taken from [4]

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2.2 Jet Diagnostics To determine ne , Te and xH we use the BE method [1]. This method assumes that the nitrogen and the oxygen ionisation fractions are determined assuming charge exchange equilibrium with hydrogen, and that the sulphur is completely ionised. In this way, using four strong optical lines (the ŒSIIœœ6716; 6731 doublet, the ŒOIœ6300 and ŒNIIœ6548 lines), and assuming the relative population of the different species, it is possible to determine the physical parameters of the jet.

3 Application to the HH30 jet We apply the tomographic inversion technique described in Sect. 2 to the determination of the three-dimensional structure of the HH30 jet. For this purpose, we use HST data taken by [5] with the Space Telescope Imaging Spectrograph (STIS), and using the Slitless spectroscopy technique. The data consists of two-dimensional images of the jet for different emission lines: the ŒSIIœœ6716; 6731 doublet, the ŒOIœ6300 and ŒNIIœ6548. The results obtained for the physical parameters are shown in Fig. 2. The electron density, temperature and Hydrogen ionization fraction are shown in groups of two panels each (left to right). For each couple of images, the first panel corresponds to the physical parameters determined using the observed emission maps, while the

Fig. 2 Original and reconstructed values of the physical parameters for the HH30 jet. From the left Panels 1-2: original and reconstructed ne , Panels 3-4: original and reconstructed Te , Panels 5-6: original and reconstructed xH

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second panel corresponds to the physical parameters determined using the inverted emission maps (i.e., the volumetric emission coefficients). Several interesting features may be noted from Fig. 2. First at all, it is evident that stellar outflows are not homogeneous! Density, temperature and ionisation fraction present inhomogeneity in the reconstructed case compared to the original data. The electron density obtained from the inversion technique is much higher on the jet axis. This effect is also observed, to a lesser degree, also for the ionisation fraction. On the other side, the electron temperature decreases on the jet axis, consistently with a stronger cooling due to higher densities therein. Finally, the jet seems to be more fragmented in the reconstructed case.

4 Conclusions We have shown how it is possible to apply tomographic techniques to invert observed high resolution emission maps, in order to estimate the three dimensional structure of a stellar jet. From the inversion of the HH30 jet we can drive the following qualitative conclusions:  The jet is much denser on the axis.  Temperature is lower on the axis.  The knots are formed by fragmented, small-scale structures.

Interestingly, the results presented here could be compared with numerical simulations of radiative jets from variable sources (see [7]) to determine the ejection history of the jet from the central disk-star object, and to obtain information on the initial density, temperature and ionization fraction profiles, and finally compared with predictions of ejection models. In particular, the fragmented density structure present along the jet may possibly imply a chaotic ejection velocity variability. These results may be used potentially to extract information on the three dimensional jet structure, to have a better understanding of the jet physics, and to compare the determined structure with those predicted by jet ejection models. All these aspects, together with a discussion of more quantitative results, will be presented in a future paper. Acknowledgements The authors kindly acknowledge Pat Hartigan for sharing his HH30 data. A.C.R. acknowledges support from the CONACyT grants 46828-F and 61547. This work was supported in part by the European Community’s Marie Curie Actions - Human Resource and Mobility within the JETSET (Jet Simulations, Experiments and Theory) network under contract MRTN-CT2004 005592.

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References 1. F. Bacciotti, & J. Eisl¨offel, A&A, 342, 717 (1999) 2. O. Bendinelli, ApJ, 366, 599 (1991) 3. I.J.D. Craig, J.C. Brown, Inverse problems in astronomy - A guide to inversion strategies for remotely sensed data, (Adam Higler, Bristol, 1986) 4. F. De Colle, C. del Burgo, & A.C. Raga, A&A, 485, 765, (2008) 5. P. Hartigan, & J. Morse, ApJ, 660, 426, (2007) 6. D.E. Osterbrock, Astrophysics of gaseous nebulae and active galactic nuclei, University Science Books (1988) 7. A.C. Raga, J. Cant´o, F. De Colle, et al., this proceedings

Measurement of Magnetic Fields in Stellar Jets Patrick Hartigan

Abstract This article reviews magnetic field measurements in jets from young stars, focusing on the physics and application of the three main techniques, Zeeman splitting and polarization, gyrosynchrotron radiation, and the analysis of shocked cooling zones. Estimates of field strengths in stellar jets are rare, and do not refer to the axis of the beam close to the source, where knowledge of the field and its geometry is most critical for constraining launching mechanisms of jets. Nevertheless, the existing measurements demonstrate that magnetic fields in YSO jets are strong enough to be important in the dynamics of the cooling zones behind internal shock waves, even though the ram pressure in the bulk flow dominates the magnetic pressure at large distances from the source. Models of pulsed magnetic flow show that velocity perturbations sweep up the field into dense working surfaces within the jet, increasing the relative importance of magnetic pressure to the dynamics in these regions and reducing its importance in the rarefaction regions that lie between the dense knots.

P. Hartigan () Physics and Astronomy Department, Rice University, 6100 S. Main, Houston TX 77005, USA e-mail: [email protected]

K. Tsinganos et al. (eds.), Protostellar Jets in Context, Astrophysics and Space Science Proceedings, DOI 10.1007/978-3-642-00576-3 38, c Springer-Verlag Berlin Heidelberg 2009 

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1 Introduction Many places in the universe that have accretion disks also drive collimated polar jets. Several examples of this process have been discussed at this conference, including jets from young stellar objects (YSOs, [1]), AGNs [16], compact X-ray sources [15], and CVs [23]. YSO jets are a particularly good place to study the phenomenon because the objects are close enough to resolve the widths and collimation properties of jets and study their proper motions. YSO jets radiate emission lines behind shock waves in the flow, and spectra of these lines determine the radial velocities and internal line widths everywhere in these regions (see [20] for a review). In addition, standard nebular diagnostic line ratios provide plasma parameters such as electron densities and temperatures, and it is also possible to learn a great deal about the accretion disk by studying the excess continuum and the emission lines that form as a result of infalling material from the disk onto the star. How accretion disks drive jets is a major unsolved problem in astrophysics. The available images of YSO jets (e.g. HH 30 [2, 11]) show that jets are already collimated within 10 AU of the source, much smaller than the size of the disk, ruling out models that rely on the ambient material to focus a spherical wind into a jet [3]. Instead, a structured B field that threads through a rotating disk naturally provides the geometry of a bipolar outflow, and all modern models use magnetic fields in some way to collimate the flow. Despite the impressive set of observational data described above, one key component of the system is missing, and that is a set of reliable measurements of the magnetic field strength and geometry in the jet. In this paper I will discuss three techniques that various researchers have used to estimate magnetic field strengths in YSO jets (Zeeman, gyrosynchrotron, and shocks) and summarize typical results that emerge from each type of study. It is also possible to infer the existence of aligned grains by polarization measurements of scattered light [4], but inferring field strengths from such data requires models of the dust properties throughout the scattering region which are highly uncertain and beyond the scope of this article. Ideally we would like to measure field strengths along the axis of jets where the velocity is highest (several hundred km s1 ) and where optical emission lines are present that constrain the densities and temperatures in the jet. Unfortunately, none of the available techniques is well-suited for such a measurement. Magnetic fields in YSO jets are too weak to be detected easily with Zeeman splitting in the optical and near-IR, and synchrotron emission is also weak and difficult to interpret. One can use an understanding of radiative shocks to infer field strengths, but this technique has thus far proved effective in only a few cases, and the measurements refer to knots >1000 AU from the source, not the collimation region close to the star. In describing of the physical processes behind each of these techniques I have drawn in part from the excellent textbooks by Rybicki and Lightman [22], Cohen-Tannoudji et al. [5], and Shu [24] for Zeeman splitting and synchrotron radiation. There are several review articles on shock physics, including Draine and

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McKee [6] for general description, Ray et al. [20] for recent observations of YSO jets, and Hartigan [8] for physics of magnetic fields and cooling zones in jets, and a previous ‘jet-set’ contribution for physics of emission line diagnostics [10].

2 Zeeman Splitting External magnetic fields lift the degeneracy of atomic energy levels by interacting with the dipole moment of a bound electron. This interaction broadens the line by splitting the levels, and the lines have different polarization properties depending on the quantum numbers of the upper and lower levels. This section summarizes how Zeeman splitting works and how it applies to YSO jets.

2.1 Physics: Effect of Magnetic Fields on Energy Levels Classically, for an electron in a circular orbit, the magnetic dipole moment is N DIA/c, where I is the current, and A is the area of the orbit. With IDev/(2r)b z and AD r2 , we obtain N DeLz /2mc where Lz is the angular momentum of the electron perpendicular to its orbit. Substituting Lz D „ml for the quantum mechanical case we expect typical magnetic splitting to be on the order of B B, where B De„/2mc is the Bohr magneton, 9:26  1021 cgs. In non-relativistic quantum mechanics with no spin the Hamiltonian for an electron in an atom is .p  eA=c/2 C esterms (1) HD 2m where esterms are the electrostatic terms between the electron and the nucleus and other electrons. For a uniform external field Bı zO, the vector potential is AD  12 r B, so that H D Hı C H1 C H2 p2 Hı D esterms C 2m e L B H1 D  2mc  e2 e2  2 .r  B/2 D H2 D x C y2 2 2 8mc 8mc

(2) (3) (4) (5)

In the above we have used [p (r B) + (r B) p] = [B (p r) + (p r) B] = 2L B, with L = r p. The terms Hı , H1 and H2 are, respectively, the unperturbed Hamiltonian, the paramagnetic term, and the diamagnetic term. The diamagnetic term is negligible, so the perturbation becomes

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H1 D  N B D 

B L B: „

(6)

Including spin, for each electron i, the magnetic moment is given by N i D

B .li C gs si / : „

(7)

In L-S coupling, the total angular momentum operator J defines the state. Using LD

S J L J JI S D JI J2 J2

and N D  we obtain H1 D 

B .L C 2S/ „

B gJ J B I „

so that

 H1 D  BBı mj

(8)

gJ D

(9)

L J 2S J C 2 J J2

3j.j C 1/  l.l C 1/ C s.s C 1/ 2j.j C 1/

(10) (11)

where we have taken the gyromagnetic ratio gs D2.0023  2. A typical transition produces a pattern of lines with a characteristic energy splitting 2 B B. Converting this energy to a radial velocity v we obtain v D 5:9  104



B G



      B km s1 D 2:8 km s1 (12) 21 cm kG 1 m

Hence, Zeeman splitting is easiest to detect at longer wavelengths. The strongest fields we expect to see in jets are on the order of a Gauss, so any splitting at optical or infrared wavelengths will be overwhelmed by thermal motions in the gas, which is why no measurements of field strengths are possible from the strong optical and near-IR forbidden lines present along jets. Figure 1 shows how light emitted by atoms in a magnetic field is polarized. Observer #2 views along the direction of the field and sees the  components (mD ˙1) circularly polarized and sees no  (mD0) component. In contrast, observer #1 sees linear polarization in the directions indicated. Thus, by subtracting a left-circularly-polarized (LCP) spectrum from its right-handed counterpart (RCP), one can measure the component of the magnetic field along the line of sight. When a magnetic field is present, the difference between the LCP and RCP spectra produces a characteristic signature of a negative residual on one side of the line center and a positive residual on the other side (e.g. [25]). Subtracting the  and  components does not help observer #1 to measure B because these components have the same average energies and do not reveal the presence of a field until the field is strong enough to clearly separate all three components in the spectrum.

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Fig. 1 Polarization of line components in Zeeman splitting

2.2 Line Splitting and Polarization Examples Most measurements of magnetic fields in star formation regions that make use of Zeeman splitting are at radio wavelengths, where the effect is easiest to measure. Observations of circular polarization at 21 cm in Orion indicate Bjj D174˙20 G [25]. This field produces a very small velocity shift between the LCP and RCP, but is still detectable because 21 cm is a strong line, molecular gas is cold, and radio observations have excellent velocity resolution. Another common application of this technique are observations of masers, where the masing mechanism causes the line emission to peak sharply over a narrow velocity range so it is possible to separate small shifts between the RCP and LCP components. A recent measurement of RCP and LCP in the star formation region W3 by Fish et al. [7] recorded a typical field of 5 mG from the masing regions. While many masers have proper motions and therefore are associated with an outflow, they usually do not define a jet, and are often time-variable and difficult to interpret. A complicating factor is that they usually occur in regions of massive star formation where multiple flows are present and dense filamentary material abounds. Masers are absent in the best examples of YSO jets. Optical and near-IR measures of fields using Zeeman splitting exist, but because the wavelengths are shorter only strong fields can be detected, and these fields typically arise from the photosphere or in accretion streams that connect the disk to the star. A recent limit of 30 G for Bjj in Ae/Be stars using the VLT to measure circular polarization in the Ca II H and K lines illustrates the current limit of the technique for a bright source [14]. A field of 2.5˙0.1 kG in the He I emission line of BP Tau was observed by Johns-Krull et al. [12], but this line forms within the accretion columns from the disk onto the star and not within a jet. Observations of absorption line broadening in the near-IR caused by Zeeman splitting exist for young stars [13], and are typically 2.5 kG. These fields come from starspots on the stellar photosphere. The observations are based on line widths rather than on polarizations, so have the advantage of being independent of the field geometry, but are limited to strong fields because absorption lines are broadened by rotation.

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3 Synchrotron and Cyclotron Emission When electrons spiral around a magnetic field lines the electrons are accelerated by the Lorentz force and emit radiation. Measuring the flux, spectral energy distribution, and polarization properties of the emitted light constrains the strength and geometry of the magnetic field. The technique has been limited to date by the relative rarity of this emission from stellar jets, but could hold more promise in the future. This section summarizes the cyclotron, gyrosynchrotron, and synchrotron emission process as they apply to YSO jets.

3.1 Physics: Continuum Emission from Particles in a Magnetic Field An accelerated charge radiates linearly polarized light in a direction perpendicular to the acceleration. As described by Rybicki and Lightman [22], there is a straighftorward way to see why this must occur. A charged particle at rest at the origin will have electric field lines that point radially away from the origin. If the particle receives a sudden impulse in the x-direction so that it moves at a constant velocity, the information about the impulse travels outward from the origin at the speed of light. Within this light sphere the retarded potentials show that the electric field points to the current position of the particle. This remarkable result implies that except along the x-axis, where the field is always in the ˙b x direction, the electric field must bend suddenly at the light sphere. As a result, the electric field has a perpendicular component at the light sphere proportional to 1/r, and therefore must radiate because there is a non-zero Poynting flux at large distances. An observer at infinity sees linearly polarized light along the projected direction of the acceleration vector. An electron moving in an external magnetic field B feels a Lorentz force perpendicular to its direction of motion that causes it to spiral around the field line with an angular frequency !B D eB= mc. For the non-relativistic case (cyclotron,   1), an external observer sees a sinusoidal electric field in the emitted radiation, and the power spectrum of this is a single emission ‘line’ at the cyclotron frequency !B . However, as  increases, the electron will progressively beam radiation in the forward direction, so the observed electric field in still periodic, but no longer sinusoidal. The Fourier transform of such a function consists of the harmonics of the gyration frequency, ! D n!B , where n is an integer. The peak of the harmonics occurs at !C   3 !B sin ˛, where ˛ is the pitch angle of the particle’s motion in the field [22]. The resulting emission is highly linearly polarized perpendicular to the field direction. For a more realistic case of a power law energy distribution of electrons (n(E)Ep ), the spectrum is also a power law with index s D .p  1/=2, and the total emitted power is proportional to Bk with k D .p C 1/=2. When  1 (synchrotron), the emission is strongly beamed into a narrow cone along its direction of motion. Hence, an observer sees almost all the radiation as

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Fig. 2 Geometry of synchrotron and gyrosynchrotron emission. An electron spirals around a uniform magnetic field as shown in the side view. The radiation is emitted in the forward and reverse direction of the electron’s motion, perpendicular to the direction of acceleration. Top: For synchrotron emission the electron is highly relativistic, so most of the emission is beamed in a narrow cone aligned with the direction of motion. Hence, the electron’s velocity vector must point almost directly toward the observer at some point in its gyration about the field, so the observer sees the path of the electron as indicated at right, and the radiation is mainly linearly polarized. Bottom: For gyrosynchrotron  1, so the opening angle of the emission cone is much wider. Thus, an observer typically sees radiation along a curved path in the sky as shown in the bottom right, implying circular polarization

arising from a small angular section of the orbital spiral where the motion of the electron points almost directly at the observer (Fig. 2, top). The apparent motion of the electron over this section as seen by the observer is linear, so the emitted light is linearly polarized. Synchrotron sources have little circular polarization because LCP and RCP radiation occur in equal amounts from electrons whose cones of emission lie just a bit above and below our line of sight, respectively. However, in the intermediate case of gyrosynchrotron emission (1<  <3) the emission is only moderately beamed, so the cone of emission is wide enough to allow an observer to see the circular motion of the electron during the emission process, and the emitted light retains a circularly polarized component unless the viewing angle happens to be exactly perpendicular to the field lines (Fig. 2, bottom). The amount of this circular polarization depends on the power law index, the pitch angle, and  , but with reasonable bounds of these parameters one can at least obtain an order of magnitude estimate of the magnetic field strength.

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3.2 The Case of T Tauri-S Observations of circular polarization in YSO jets are rare. The best example is that of Ray et al. [19], who detected LCP and RCP emission in 6 cm observations of the companion to T Tauri. The polarized emission was offset on either side of the star, suggestive of an ordered field structure away from the photosphere. The field strength was on the order of a few Gauss, assuming a typical range for the unknown pitch angles,  , and power law index. While intriguing, this measurement is difficult to interpret because the spatial resolution of the observations is similar to the offset of the emission from the star, so a jet geometry is not necessarily implied by the observations. If the emission does arise in a jet, it probably comes from a strong shock close to the star, and then the average field in the jet will be much lower.

4 Magnetic Pressure in Cooling Zones of Shocks This section outlines how one can combine the large number of emission lines present behind shocks in jets with detailed kinematic information about velocities to facilitate measurements of magnetic fields.

4.1 Physics: Magnetic Pressure in Postshock Cooling Zones Estimating field strengths within radiative shocks is in principle a straightforward concept. If we know the shock velocity and the preshock density, then conserving mass, momentum and energy across the shock implies that the density increases by a factor of four in a strong atomic shock [6]. Since the magnetic field is tied to the fluid, the component of the field that lies parallel to the shock front (perpendicular to the normal) must also increase by a factor of four. As the gas cools it compresses, and the magnetic field rises in proportion to the density. Because the magnetic pressure Pm B2  2 where B is the field strength and

the density, the ratio Pm /Pgas  . This ratio increases in the postshock zone until the magnetic pressure dominates in the dense cooling zone, even if the initial field was dynamically unimportant (i.e., B2 =8 v2 . Hence, it is possible to infer a field strength by comparing the observed density in the cooling zone with that expected for a nonmagnetic shock [8].

4.2 Application of Cooling Zone Method to Stellar Jets While simple in principle, applying the above method requires the shock to have a relatively simple geometry, and have a well-defined shock velocity and preshock

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density. Shocks in YSO jets often exhibit complex morphologies, but examples also exist of well-defined bow shocks. Estimating a shock velocity is possible by combining the observed line ratios with the shape of the bow shock. It is easiest to do this if the bow shock is strong enough to emit [O III] lines, which only occur when the shock velocity exceeds  90 km s1 . Essentially the point where the shock velocity drops below this value is marked by the edge of the [O III] emission along the bow shock, and the shape of the bow then gives the shock velocity. The preshock density is then measured from the total line fluxes, which must be dereddened. Finally, one can measure the electron density from emission line ratios of [S II] (see e.g. [10]), which implies a total density from the ionization fraction estimated from other emission lines. While uncertainties exist in all these measurements, the difference between the observed density in the cooling zone and that expected for a nonmagnetic shock is typically an order of magnitude, so the presence of even a weak field is easy to detect. Morse et al. [17, 18] applied the above method to two YSO jets and found B  15  G for the preshock gas in front of major bow shocks in the flows, and about a factor of 10 higher in the postshock gas. These bow shocks are located 5  104 AU from the source, so to infer anything about the field strength close to the source we need to understand how the field strength varies with distance in a typical YSO jet. A recent study that focuses on this issue is described in the next section.

5 Implications for a Velocity-Variable Magnetic Flow The standard picture of a magnetized jet is that the field is anchored in the accretion disk and becomes increasingly toroidal as the jet becomes collimated. However, if the field is too strong, then it will inhibit the formation of shocks. We observe shock waves at 105 AU when velocity pulses are only 30 km s1 , so the fields there must be less than about a mG, consistent with those estimated from the cooling zone method described above. Closer to the sources, the 1 G fields found from the circular polarization measurements imply a stronger field. Are the weak fields observed at large distances (104 –105 AU) consistent with stronger fields inferred within the acceleration regions (10 AU) of YSO jets? Jets appear to expand at roughly a constant opening angle, so the average density declines with distance from the source approximately as r2 . For a toroidally dominated disk wind, Br1  n0:5 , so the magnetic signal speed, which is governed by longitudinal magnetosonic waves, should be roughly constant in such a flow. However, within a shock wave B  n, so we might expect B  r2 in these regions. This steeper dependence of B on r would help to reconcile strong fields close to the star with weaker fields at large distances. Recent models of magnetized pulsed outflows show that they behave in a manner intermediate between the disk wind and shock limits described above [9]. Essentially velocity pulses in the jet sweep up the field into a few dense knots, and create rarefactions between the knots. Within rarefactions the magnetic signal speed drops

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markedly, which makes shock waves possible in the flow. The models which best reproduce the observed shock velocities are ones where the terminal magnetosonic Mach number of the flow is  5. Lower terminal magnetosonic Mach numbers are unable to produce both the high velocities of the jets and the low shock velocities present in the velocity perturbations (small amplitude perturbations become magnetic waves rather than MHD shocks).

6 Summary Despite their probable importance in YSO jets, magnetic fields have proved elusive to measure in these objects. Nonetheless, field estimates do exist, and the available measurements appear to be generally consistent with the notion of a jet driven magnetically from the source, with velocity pulses playing an important role in sweeping the field into a few dense regions, and areas of much lower field prevailing between the knots. Dynamically the field is strong enough in the dense knots to dominate the pressure, but must be significantly lower than the overall ram pressure of the flow in order to allow the weak shocks observed, which have shock velocities only  10% of the flow speed. Probably the most exciting future prospect for measuring field strengths in YSO jets is Zeeman splitting of molecular lines with ALMA (see the article by Richer in this volume [21]). While this technique will only be applicable to the youngest jets such as HH 212 which have strong molecular emission along the axis of the flow, any additional quantitative measurements of field strengths in these jets will go a long way to clarifying how YSO jets become collimated and how they are driven by their accretion disks. Acknowledgements This research was supported by NASA Origins of Solar System Grant NNG05GH97G to PMH.

References 1. Bally, J. 2009, in Protostellar Jets in Context, K. Tsinganos, T. Ray & M. Stute eds., (New York:Springer) 2. Burrows, C. et al. 1996, ApJ, 473, 437 3. Canto, J. 1980, A&A, 86, 982 4. Chrysostomou, A., Lucas, P., & Hough, J. 2007, Nature, 450, 71 5. Cohen-Tannoudji, C., Diu, B., & Lalo¨e, F. 1977, Quantum Mechanics, (New York:John Wiley and Sons) 6. Draine, B., & McKee, C. 1993, ARA&A, 31, 373 7. Fish, V., Brisken, W., & Sjouwerman, L. 2006, ApJ, 647, 418 8. Hartigan, P. 2003, Ap&SS, 287, 111 9. Hartigan, P., Frank, A., Varniere, P., & Blackman, E. 2007, ApJ, 661, 910 10. Hartigan, P. 2008, Jets From Young Stars II, Lecture Notes in Physics, 742, 15 11. Hartigan, P., & Morse, J. 2007, ApJ, 660, 426

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12. Johns-Krull, C., Valenti, J., Hatzes, A., & Kanaan, A. 1999, ApJ, 510, L41 13. Johns-Krull, C., Valenti, J., & Saar, S. 1999, ApJ, 617, 1204 14. Hubrig, S., Yudin, R., Sch¨oller, M., & Pogodin, M. 2006, A&A, 446, 1089 15. Livio, M. 2009, in Protostellar Jets in Context, K. Tsinganos, T. Ray & M. Stute eds., (New York:Springer) 16. Massaglia, S. 2009, in Protostellar Jets in Context, K. Tsinganos, T. Ray & M. Stute eds., (New York:Springer) 17. Morse, J., Hartigan, P., Cecil, G., Raymond, J., & Heathcote, S. 1992, ApJ, 399, 231 18. Morse, J., Heathcote, S., Cecil, G., Hartigan, P., & Raymond, J. 1993, ApJ, 410, 764 19. Ray, T., Muxlow, T. W. B., Axon, D. J., Brown, A., Corcoran, D., Dyson, J., & Mundt, R. 1997, Nature, 385, 415 20. Ray, T., Dougados, C., Bacciotti, F., Eisl¨offel, J., & Chrysostomou, A. 2006, in Protostars and Planets V, B. Reipurth, D. Jewitt, & K. Keil eds., (Tucson:University of Arizona Press). 21. Richer, J. 2009, in Protostellar Jets in Context, K. Tsinganos, T. Ray & M. Stute eds., (New York:Springer) 22. Rybicki, G., & Lightman, A. 1979, Radiative Processes in Astrophysics, (New York: John Wiley and Sons) 23. Sokoloski, J. 2009, in Protostellar Jets in Context, K. Tsinganos, T. Ray & M. Stute eds., (New York:Springer) 24. Shu, F. 1991, The Physics of Astrophysics Volume 1: Radiation, (Mill Valley CA:University Science Books). 25. Troland, T., Heiles, C., & Goss, W. 1989, ApJ, 337, 342

Jet Kinematics Alessio Caratti o Garatti and Jochen Eisl¨offel

Abstract We review some of the most important kinematical studies on protostellar jets, from the pioneering work of Herbig and collaborators, to the high-spatial resolution Hubble Space Telescope observations of jets, concluding with the latest works on time-variability and velocity changes of knots along the flows.

1 Introduction Our knowledge of protostellar jets has tremendously increased during the past decades. In particular big steps forward have been made by studying their kinematics. In this contribution we briefly review some of the most important achievements in the field, from the pioneering work of Herbig and collaborators (e.g. [9, 31]), to the high-spatial resolution Hubble Space Telescope observations of jets (e.g. [42, 26]), concluding with the latest work on time-variability

A. Caratti o Garatti () and J. Eisl¨offel Th¨uringer Landessternwarte Tautenburg, Sternwarte 5, D-07778 Tautenburg, Germany e-mail: [email protected]; [email protected]

K. Tsinganos et al. (eds.), Protostellar Jets in Context, Astrophysics and Space Science Proceedings, DOI 10.1007/978-3-642-00576-3 39, c Springer-Verlag Berlin Heidelberg 2009 

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(e.g. [3,34,4]), and velocity changes of knots along the flows [4]. Due to the limited space, we can only review some of the works published so far. The observational strategy used in kinematical studies is usually based on: a) proper motion (P.M.) measurements (using two or more images at different epochs); b) radial velocities measurements (by means of high resolution spectroscopy); c) a combination of the two to obtain the 3-dimensional kinematics of the jet (i.e. the absolute flow speed and the inclination of the flow with respect to the plane of the sky). With the evolution of new technologies (large-format CCDs, IR arrays, IFU, adaptive optics, etc...), several new techniques have been introduced as, for example:  Multi-epoch and multi-wavelength narrow band imaging (centred on various jet

emission lines), to study the different kinematical components at various wavelengths (optical, infrared) (e.g. [25]).  Large field multi-epoch imaging to entirely map the kinematics of parsec-scale jets (e.g. [13, 36]).  Higher spatial and spectral resolution at optical and NIR wavelengths to study the internal structure of the jets and the bow shock structure (e.g. [11, 10]).  Multi-epoch slitless spectroscopy of jets provides emission-line images for each line in the spectrum [27]. P.M.s, radial velocities, and physical properties can be simultaneously derived . A similar analysis can be done using IFU data [23], but the typical FoV is still too small and just small portions of the jet can be covered on a single measurement. The first P.M. studies of Herbig-Haro (HH) objects (end of 1970s, beginning of 1980s [9, 31]) were based on photographic plates and wide-band filters, with long time baselines (25–35 yrs), and relatively low spatial resolution. These measurements could provide bulk motions for the flow, rather than the motion of single knots. The relatively high tangential velocities measured (100–300 km s1 ), the alignment of the flows, and the opposite directions of the jets (e.g. for HH 1 and 2), suggested a shock origin, the association to young stellar objects (YSOs), and the bipolar nature of the jets. This was promptly confirmed by further spectroscopic analysis (e.g. [14, 42] and references therein), that revealed blue- and red-shifted lobes in the HH 46/47 jet and other flows. Moreover, combining distance and velocities, the dynamical age of some flows was inferred for the first time (see e.g. [43] and references therein).

2 Internal jet Structure 2.1 Motions within the beam and the bow shock Structure With the advent of CCDs several P.M studies of HH objects (HH 1/2, 34, 46/47, 111 etc..) were performed (e.g. [16,25,30,17,20]). Narrow-band filters and a higher spatial resolution made it possible to investigate the motion of single knots within

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RELATIVE FLUX

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Fig. 1 Upper Panel: Model and observed line profile of the HH 32 A bow shock (adapted from [29]). Lower Panel: Internal flow pattern in the HH 1 bow shock (adapted from [20]). The P.M. at the stagnation point of the bow shock (i.e. at the apex) was subtracted to illustrate the relative motion of the condensations in the bow shock

the flow and to confirm the bow shock geometry (see Fig. 1). According to bowmodels, the shock velocity is higher at the apex of the bow and decreases along the wings. Subtracting the “mean” P.M. vector of the jet from the P.M. vectors of the bow shock condensations, the internal motions (in the shock reference frame) become visible (e.g. [16] and Fig. 1, lower panel). This kind of shock geometry produces double peaked emission line profiles as well, usually observed by means of high resolution spectroscopy at both optical (e.g. [29,40] and Fig. 1, upper panel) and NIR wavelengths (e.g. [11]).

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the sky

observed radial velocity knot pattern speed particle flow speed

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20 40 60 Distance from source [ ” ]

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Fig. 2 Upper Panel: Relations between the observed P.M. and radial velocity in a jet, the pattern speed of the knot, the flow speed in the jet and the inclination angle of the jet with respect to the plane of the sky. Lower Panel: Ratio  of the pattern speed to flow speed for knots in the HH 34 jet. The marked region represents the expected range of values for  (0.82–0.91) of the HH 34 knots for a IWS scenario (see text and [16]). The figures have been adapted from [16])

The spatial orientation of the jet and its space velocity can be determined combining tangential (vt , derived from the P.M.) and radial (vr ) velocities. Note, however, that the measured P.M. (see Fig. 2, upper panel) is the projected knot pattern speed (vp , i.e. the phase velocity), which can be smaller than the projected flow speed (vf , i.e. the group velocity) of the jet. On the other hand, vr is the radial component of the particle flow speed. As a consequence, vt derived from P.M.s can be erroneous (i.e. lower than the real value), and the inferred jet inclination angle ˛ can be overestimated. Defining the ratio  D vp =vf , the measured vt coincides with the projected vf only when  D 1. In principle, ˛ can be derived from the terminal bow shock of the jet (in the knot closest to the apex), under the following assumptions: (1) it moves along the symmetry axis of the bow shock, and (2) its motion is stationary relative to the physical apex of the bow shock. In this case  D 1, from the P.M. the real vt is inferred, and ˛ D arctan.vr =vt ).

2.2 Knot Formation An important issue about jets, not yet completely solved, concerns the origin of the knots inside the flow. The formation of these shocks can be caused by KelvinHelmholtz (K-H) instabilities (e.g. [38]) or by time variability in the flow ejection

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velocity (e.g. [39, 12] and references therein), producing internal working surfaces (IWS). The main observational differences between the two mechanisms are the following [38]: (a) For IWS, the line fluxes from knots show a sudden rise and then a regular decline with distance, while in the K-H mechanism they exhibit a smooth increase and slow decrease. (b) In the IWS scenario  is close to one for all the knots. The knots are moving slower than the jet gas at vp D vf -vshock . Then  D(vf -vshock )/vf . Thus the expected  value can be obtained if the vshock is known (for example from spectroscopy, see e.g. [29]). On the other hand, the K-H mechanism produces shocks whose P.M. velocities tend to increase with distance. Then  has low values close to the source (0.5) and increases with distance. In Fig. 2 (lower panel) the measured  values in the HH 34 knots and the expected range of  values for a IWS scenario are reported (see also [16]). Here, knots close to the source are more likely originating by K-H instabilities.

2.3 PMs from HST The limited spatial resolution of ground based P.M. leads to a need for longer time baselines to detect the motion. However, when the epoch differences are becoming comparable to the post shock cooling time of the knots (from some years to few decades, depending on the species), knots may show significant variability (see Sect. 3). As a consequence, P.M. measurements may become increasingly inaccurate. HST observations offer a possible solution to this problem increasing the spatial resolution and decreasing the length of the time baseline (usually 2–4 years, see e.g. [42, 26]). These observations showed that: (1) the jets in the observed HH objects usually show large flow velocities and small velocity dispersions; (2) the tangential velocities decrease with increasing distance from the jet axis; (3) in some complex flows, such as HH 2, both bow shock and reverse bow shocks are detected. These last are always slowly moving or stationary knots.

2.4 Velocity Components, Asymmetries, and jet Rotation High resolution spectra of jets in classical T Tauri stars (CTTS) often reveal the presence of several velocity components close to the source: the so-called highand low-velocity components (HVC, LVC). HVC is associated to the extended collimated jet, whereas the LVC is usually confined near to the source (up to 100– 200 AU) (see e.g. [15] and references therein). Both components are observed in forbidden emission lines and, more recently, in molecular lines [2]. These features have been predicted by MHD theoretical models.

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Several studies have also shown velocity asymmetries between the blue and red lobes of several CTT jets (e.g. [32]). This phenomenon could be associated to different accelerations in each lobe. Another exciting spectral feature in CTT jets (as DG Tau, RW Aur, Th 28 and CW Tau) is the presence of radial velocity asymmetries across the jet axis (e.g. [1,8]) interpreted as jet rotation. Such asymmetries have been detected at different wavelengths from UV to NIR (see Bacciotti and Coffey’s contributions, this volume). Recently, evidence of rotation has also been detected in Class I jets [6].

3 Large Scale jet Structure 3.1 Parsec-scale Flows Several HH objects and jets from Class 0 and I YSOs appear to be giants flows, extending up to several parsecs from the source, as was first shown for the HH 34 jet [13]. Dozens of such giant flows have been discovered so far (e.g. [19, 41, 18]). Recently, giant flows from Class II YSOs have also been detected [36]. An important characteristic of giant flows is their systematic decrease in radial velocities and P.M.s for increasing distances from the driving source (e.g. HH 34, HH 111, [13,41]). The deceleration should be mainly due to progressive momentum transfer from the jet to the surrounding medium. A comparison of the observed P.M.s or radial velocities with the shock velocities indicates that some HH objects can be internal working surfaces, which form in the interaction between old slower ejecta and young faster ejecta (e.g. [28, 4], and Sect. 4).

3.2 H2 Kinematics For historical reasons H2 kinematical studies of jets are relatively new in the literature (e.g. [11, 37, 10, 7, 33, 22]). Nevertheless they are fundamental to analyse the dynamics of those embedded jets (mostly from Class 0 sources) which are not visible in the optical. In addition, the H2 kinematics is usually complementary to the optical (i.e. with H˛, [SII]). Several P.M. studies indicate that H2 knots in jets and HH objects often move at high tangential velocities (100–300 km s1 ), often very similar to their optical counterparts (e.g. [37, 7, 22]). These values are much higher than the dissociation speed limit of H2 (up to 80 km s1 , for the most extreme case). This implies that the H2 features are comoving with the flow, and that the measured velocities do not represent shock velocities.

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Fig. 3 Left Panel: Precession model plotted over the H2 1-0 S(1) line contours of the Cep. E jet (from [21]). The model is for a precession angle of 4ı , and a precession length scale of 22”.8. Given the velocity and the distance to the jet, useful parameters as the precession period can be derived. In case of a binary system a rough estimate of the separation between the sources and their masses can be obtained as well. Right Panel: [SII] image of HH 110 along with P.M.s (from [35]). The jet deflection is likely caused by a jet/cloud collision

3.3 Jet Precession, Bending and Deflection Since shocks in jets trace ejecta that are progressively older with increasing distance from the source, they provide a “fossil record” of the YSO ejection history. In particular, their geometry provides us with information on the dynamics of the driving source and its environment. S- or Z-shapes usually stem from precession (see Fig. 3, left panel), generated by tidal forces in a multiple system (e.g. [21, 5]). From the jet geometry, several dynamical parameters can be inferred, as the precession period, an estimate of the separation between the sources and rough estimate of their masses (e.g. [44]). C-shaped bending jets usually indicate the presence of side-winds (see Ciardi et al., this volume) or the motion of the source itself (e.g. [24]). Finally, some jets appear to be deflected (e.g. HH 110 [35], HH 54 [4]), likely due to the jet collision with a denser cloud or clumps (see Fig. 3, right panel).

4 Jet Time-Variability Recent observations of protostellar jets show time variability in both emission (e.g. [20, 33, 3, 4] and velocity [34, 4]), on time scales of some years.

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Our latest study on time-variability in the Cha II flows [4], based on several multi-wavelength observations over 20 years, reveals that on average the analysed knot fluxes are variable with time. About 80% of the analysed knots show some variability. Such variations often (70%) do not have a particular trend, but rather appear as temporary fluctuations around a mean value. For these objects, the P.M.s, inside the error bar, are usually constant with time. For the remaining knots a clear trend is visible (see Fig. 4), and the majority shows velocity variations as well. The structures which exhibit the highest variability in morphology, flux, and velocity can be mostly identified with (1) working surfaces or (2) interacting knots, where the faster fluid particles ejected at later times eventually catch up with the slower flow ejected earlier from the source. Usually, the fast gas from behind is decelerated, while the slow gas in front is accelerated. The interaction between two knots is well visible in HH 54 G0 and G (see Fig. 5, upper left panel). The quadratic fits performed on measured shifts of the knots (see right panel) indicate both deceleration (in G0) and acceleration (in G). Both H˛ and [SII] quadratic fits give similar values for the accelerated and decelerated motions measured in the two knots (see bottom left panel). Flux time variability is observed as well (central left panel). Figure 6 shows another clear case of accelerated motion along the HH 54 flow. In total we observed 17 knots which show genuine accelerated motions. Among them 12 objects accelerate and 5 decelerate. For most of them both modules and P.A.s of proper and accelerated motions are identical (inside the errors) in both H˛ and [SII] filters.

Fig. 4 Examples of flux variability in various knots of HH 52 B5 in [SII] (circles), H˛ (triangles), and H2 (squares) filters. Note that different filters have different timings, in particular the rise of the H2 emission has a delay of about 10 years. HH 52 C1 (central panel) has a trend similar to B5 (top panel), but the collision with a fast overtaking shock after 1995 newly increases its luminosity. Knot HH 52 C2 (bottom panel) represents a case of [SII] variability (from [4])

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Fig. 5 Upper Left Panel: Interaction of knots in HH 54. The fast knot HH 54 G0 impacts against G. Right Panel: Linear and quadratic fits (straight and dashed lines) performed on measured H˛ shifts in HH 54 G0 and G. The quadric fits model the data better, indicating deceleration in G0 and acceleration in G. Central Left Panel: Changes of proper motions in HH 54 G (triangles) and G0 (circles) with time. Open and filled marks indicate measurements in H˛ and [SII], respectively. Bottom Left Panel: Flux variability in HH 54 G and G0. All figures adapted from [4]

5 Summary We reviewed some important topics on protostellar jet kinematics, from their internal structure to the large scale properties. Other worthwhile issues, as for example studies on mass flux rate, dynamical age, etc., had to be omitted for space limitations, although they are fundamental to understand the YSO evolution history and the jet/medium interaction. From this brief review it becomes clear that these objects are not just a fossil record of the YSO ejection history, but they are still evolving with time, providing us with information about MHD shock evolution, as well.

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Fig. 6 Linear and quadratic fits performed on measured H˛ shifts in HH 54 Y2 (from [4]). The best fit is the quadratic, indicating that the knot is accelerating

Acknowledgements The present work was supported in part by the European Community’s Marie Curie Actions-Human Resource and Mobility within the JETSET (Jet Simulations, Experiments and Theory) network under contract MRTN-CT-2004 005592.

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Synthetic Jets – from Models to Observations and Back Jos´e Gracia

Abstract The main purpose of numerical simulations is to study the nature of physical processes at work in astrophysical objects. In general it is very difficult to simulate and match a particular object pixel by pixel, both due to the complex microphysics involved and due to the only insufficiently known initial and boundary conditions as imposed by the environment. Still, careful modelling and inversion of observations can unveil the structure of particular sources. We discuss two examples: (1) constraining the velocity field for YSO jets through inversion of line profiles, and (2) inferring the magnetic field topology for AGN jets from synchrotron emission maps.

1 Introduction Numerical simulations and observations cannot be compared directly. While the former describe the plasma state in terms of quantities like density, pressure, magnetic field and velocities, the latter observes only photon flux as a function of frequency.

J. Gracia () Dublin Institute for Advanced Studies, School of Cosmic Physics, 31 Fitzwilliam Place, Dublin 2, Ireland e-mail: [email protected] K. Tsinganos et al. (eds.), Protostellar Jets in Context, Astrophysics and Space Science Proceedings, DOI 10.1007/978-3-642-00576-3 40, c Springer-Verlag Berlin Heidelberg 2009 

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This comparison can be facilitated by either doing forward modelling, i.e. calculating synthetic emission maps from a given model and comparing to observations, or by backward modelling (called inversion henceforth), i.e. inferring the plasma state from observations by diagnostic or inversion methods. Both methods however pose severe difficulties. Translating numerical simulations into synthetic maps is highly non-trivial; in general, the local emissivity is a complicated function of temperature, electron density and the density of the respective ion, as e.g., nOII for singly ionised oxygen. In addition, the emissivity may also depend on the history of the ionisation state, etc. The emissivities are integrated along a given line-of-sight and projected onto the plane of the sky producing an ideal synthetic emission map. Finally, real detectors introduce additional bias through their detector response, which needs to be taken into account. In summary forward modelling needs to cope with a huge parameter space for the initial numerical simulation as well as complicated microphysics in order to calculate synthetic observations. This approach has been chosen e.g. by [7, 6] (see also their contribution in this volume) in order to compare self-similar models to specific jets based on measured jet width. On the other hand, in some cases it seems possible to better exploit current observations by using inversion techniques. For optically thin sources, the photon flux recorded in the detector is a superposition of the emission along the line-of-sight (LoS). Most diagnostic methods will only reveal quantities integrated along the LoS or average quantities weighted by the local emissivity. For systems, where the emissivity is strongly stratified along the LoS, those average quantities may be severely misleading (e.g. De Colle et al., this volume). However, in principle it is possible to reconstruct the true 3-dimensional structure of a source, in particular of its velocity field or magnetic field as shown in the following, by making additional assumptions about its geometry as for example axial or spherical symmetry. The main problems for this approach are the deconvolution due to the spatial resolution (PSF), and more importantly, the deprojection along the LoS. While photons don’t carry an ‘origin’ tag, it may be guessed or reconstructed for highly symmetric sources.

2 Inversion of Position-Velocity Diagrams for Stellar Jets Spectral line profiles do contain information on the velocity of the radiating plasma along the LoS. For each volume element along the LoS, the natural line is Dopplershifted according to its radial velocity (RV), i.e. the projection of the velocity vector onto the LoS direction vector. Position-velocity diagrams (PVDs) across or along the jet are routinely used to estimate the kinematics in a jet. However, given a good estimate for the inclination angle and assuming local cylindrical symmetry, PVDs across the jet can be used to constrain the true 3D velocity field of the jet, since the radial velocity can be expressed as RV D v sin  sin   vZ cos ;

(1)

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where , , vZ , v , are the azimuthal angle, inclination angle, azimuthal velocity, and the velocity along the jet axis, respectively. The algorithm could be described by onion peeling and can be sketched in the following way: imagine the jet to be composed of concentric, uniform, cylindrical shells. A spectrum taken on the jet axis will contain emission from all shells, while a spectrum taken further out will miss contributions from the innermost shells. A spectrum taken at the outermost shell position will show a significant signal in only one velocity bin. However, the signal may be recorded in different velocity bins, RVl and RVr , on either side of the jet. The difference is given by the azimuthal velocity in the outermost shell RVr  RVl D 2v sin  , while the sum RVr C RVl D 2vZ cos  is given by the velocity along the axis. Knowing these values at the outermost shell, one can construct the spectrum of the outermost shell as it would be seen on the next inner shell (i.e. taking into account ) and compare to the observed spectrum. Any disagreement must be due to plasma emitting on the next inner shell, and can be used to reconstruct v and vZ there. This procedure can be applied to consecutive innermore shells much like peeling an onion. In reality one has to record at the same time the emissivity of the shells. Also, due to the finite resolution of the velocity bins, the calculated values for the velocity components are determined only up to the width of the bins. In order to improve these guesses one has to iterate the procedure until convergence is achieved. To test the algorithm, we have calculated synthetic PVDs for the simulations presented in [6] using a set of tools described in [2] with a realistic spectral resolution of RV D20 km s1 . Figure 1 shows profiles of the true (in this case simulated) velocity components and the reconstruction with the algorithm described above. For this test we used a very optimistic spatial resolution with more than a hundred resolution elements across the jet (actually the spatial resolution of the numerical simulation). However, the algorithm works also at much lower resolution and could be applied to current or near future observations as long as the jet width is resolved.

3 Inversion of Synchrotron Emission for AGN Jets In contrast to thermal emission processes in YSO jets, synchrotron emission in AGN jets is fairly simple to model [5] and depends largely only on the density of the electrons, the magnetic field strength, and the direction to the observer. The latter is particularly interesting and makes synchrotron emission a very useful diagnostic. The intensity of the emission is given by the component of the magnetic field perpendicular to the direction towards the observer, i.e. the component in the plane of the sky. Further, the linear polarisation angle coincides with the direction of the magnetic field in the plane of the sky. Therefore, measurement of all Stokes components in principle allows the complete reconstruction of the magnetic field in the plane of the sky, both intensity and direction, if we could disentangle the superposition of different radiating plasma volumes along the line-of-sight.

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Fig. 1 Inversion of PVDs for YSO jets. The left panel shows azimuthal velocity v across the jet, the right panel velocity along the jet vZ . In both panels the light-shaded curve corresponds to the true velocity (in this case simulated data), while the black curve is the reconstruction from the PVD

While the superposition of radiation along the line-of-sight may be viewed as a curse, in the particular case of synchrotron emission in conjunction with a quasiaxisymmetric source it may be seen as blessing. Axisymmetry correlates the plasma state along the line-of-sight with the plasma state across the jet. The same annulus is seen at varying depth along the line-of-sight at different positions across the jet. This additional information in principle allows us to reconstruct the three-dimensional structure by inversion techniques. Furthermore, since the magnetic field in a particular annulus, is seen at varying orientation towards the observer at different positions across the jet, it is possible to infer all magnetic field components and not only the component in the plane of the sky. The jet of M87 is the best studied extragalactic jet. Gracia et al. [3, 4] recently presented MHD models and synthetic synchrotron emission maps that not only reproduce the observed opening angle over several orders of magnitude in length scale, but also reproduce key morphological features of the M87 jet, as the pronounced limb-brightening and the position of the optically bright knot HST-1. This simulations suggest, that most of the emission originates from a thin shell at finite distance from the jet axis. However, the MHD models can by construction not predict the small scale variations of the magnetic field. We suggest to go one step further and invert the radio observations of the jet in order to extract the local structure of the magnetic field. In order to do so, we will fit ‘thin shell’ models (a relativistic generalisation of [5]) to the measured profiles of Stokes I and Stokes Q across the jet. We have tested the method with realistic synthetic emission maps for the jet of M87. Figure 2 shows a (synthetic) map of total intensity Stokes I and a map calculated from the inverted magnetic field. The inversion algorithm shows some glitches where the jet width is underestimated. However, in general the jet width and limbbrightening are well reproduced. The same method will be applied to real VLBA data for the inner jet of M87 in the near future and will yield the magnetic field

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Fig. 2 Inversion of synchrotron emission maps for AGN jets. The left panel shows the true emission (in this case actually a synthetic map from simulated data), the right panel map calculated from the inversion

topology along the jet of M87 in terms of the helical pitch angle, the Lorentz factor, and the product of electron density times magnetic field strength. In addition we will constrain the relative strength of the turbulent magnetic field component. These results will be used as input for subsequent MHD models.

4 Concluding Remarks Application of inversion techniques can potentially reveal the 3D kinematics or magnetic field topology in jets. The demands on data quality are stringent in terms of signal-to-noise and resolution, both spectral and spatial. However current facilities are capable of meeting these; certainly for the case of synchrotron emission from AGN jets (e.g. VLBA), and probably for spectral line inversion in YSO jets (e.g. VLT/SINFONI). Result from those inversions will constrain existing MHD models. The application of these techniques is certainly not limited to jets, but can be extended to any system with reasonably known geometry showing optically thin emission lines dominated by Doppler effects due to plasma motion.

References 1. Gracia, J., Bogovalov, S., Tsinganos, K.: MHD models and synthetic synchrotron maps for the jet of M87. ArXiv e-prints, arXiv:0712.2734, (2007) 2. Gracia, J., Stute, M. et al.: From jet simulations to synthetic observations: the JETSET pipeline. In preparation

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3. Gracia, J., Tsinganos, K., Bogovalov, S. V.: Magnetic collimation of the relativistic jet in M 87. A&A, 442, L7–L10, (2005) 4. Gracia, J., Vlahakis, N., Agudo, I., Tsinganos, K., Bogovalov, S. V.: Synthetic synchrotron emission maps from MHD models for the jet of M87. ApJ, 695, 503–510, (2009) 5. Laing, R. A.: Magnetic fields in extragalactic radio sources. ApJ, 248, 87–104, (1981) 6. Stute, M., Gracia, J., Tsinganos, K., Vlahakis, N.: Comparison of synthetic maps from truncated jet-formation models with YSO jet observations. To appear in A&A, (2009) 7. Stute, M., Tsinganos, K., Vlahakis, N., Matsakos, T., Gracia, J.: Stability and structure of analytical MHD jet formation models with a finite outer disk radius. A&A, 491, 339–351, (2008)

X-Ray Emission from Young Stellar Jets ¨ Manuel Gudel, Stephen L. Skinner, Sylvie Cabrit, Jochen Eisl¨offel, Catherine Dougados, Roland Gredel, and Kevin R. Briggs

M. G¨udel () and K.R. Briggs Institute of Astronomy, ETH Z¨urich, 8092 Z¨urich, Switzerland e-mail: [email protected] S.L. Skinner CASA, University of Colorado, Boulder, CO 80309, USA e-mail: [email protected] S. Cabrit L’Observatoire de Paris, 75014 Paris, France e-mail: [email protected] J. Eisl¨offel Th¨uringer Landessternwarte Tautenburg, Tautenburg, Germany e-mail: [email protected] C. Dougados LAOG, Universit´e Joseph Fourier-CNRS, 38041 Grenoble Cedex, France e-mail: [email protected] R. Gredel Max-Planck-Institute for Astronomy, 69117 Heidelberg, Germany e-mail: [email protected]

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Abstract Jets from protostars and T Tauri stars define a separate and new class of X-ray sources. New evidence suggests that X-ray emission originates in microjets of several classical T Tauri stars. We summarize the present status of our search for jet-related X-ray emission and discuss physical implications for the jet gas.

1 Introduction X-rays have been detected from a few Herbig–Haro objects [13, 6], but also from jet regions relatively close to a protostellar binary [4]. X-ray jets close to young stars may help efficiently ionize larger parts of the circumstellar environment than could the central star alone. In particular they ionize the disk surface, thus inducing disk accretion instabilities [1] and altering the disk and envelope chemistry [5]. Strong attenuation of soft X-rays close to protostars prohibits systematic verification. Recently, X-rays have been discovered from “microjets” of classical T Tauri stars (CTTS), both in high-resolution images [7, 10] and spectroscopically [7, 8].

2 A Bipolar X-Ray Jet from DG Tau Figure 1a shows a Chandra X-ray image of the CTTS DG Tau [10]. There is clear evidence for a jet-like extension to the SW of the stellar image, along a position angle of 225 deg, but we also find a significant excess of counts in the NE direction (PA 45 deg). This is coincident with the optical jet axis, which for the SW jet has been given as 217–237 deg [3]. We verified, using raytrace simulations, that the jet sources are extended: a faint point source would occupy only a few pixels. The temperature of the emitting plasma is of order 3–4 MK, with a total 0.1–10 keV X-ray luminosity of 1:2  1028 erg s1 per jet [10].

3 Two-Absorber X-Ray Spectra: Evidence for Jets? X-ray spectra of several very strongly accreting, jet driving CTTS exhibit an anomaly (Figs. 1b and 2, Table 1): These “Two-Absorber X-ray” (TAX) spectra show a cool component subject to very low absorption and a hot component subject to absorption about one order of magnitude higher. The cool component shows temperatures atypical for T Tau stars, ranging from 3–6 MK. The spectrum of the CTTS Sz 102 (Fig. 2b) does not explicitly show two absorbers but is very strongly dominated by an unusually soft component similar to that in TAX sources. The hard component of DG Tau point source (Fig. 1b) requires an absorbing hydrogen column density (NH  2  1022 cm2 ) higher by a factor of 5 than predicted from the visual extinction AV of 1.5–3 mag if standard gas-to-dust ratios

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Fig. 1 (a) (left) Chandra image of DG Tau with jets in NE–SW direction (pixel size 0.4900 ). (b) (right) XMM-Newton CCD spectrum of DG Tau, illustrating the two components

Fig. 2 (a) TAX spectrum of HN Tau. (b) spectrum of Sz 102 (fit with one absorber)

Table 1 Star Star DG Tau GV Tau DP Tau HN Tau CW Tau a

TAX sources: X-ray parameters and general propertiesa AV L([O I])b lg MP wc NH;s Ts LdX;s 30 22 2 (mag) (10 ) (Mˇ /yr) (10 /cm ) (MK) (1029 ) 1.5–3 24  6:5 0.11 3.7 0.96 3–5 1.2  6:5 0.12 5.8 0.54 1.2–1.5 0.06  7:4  0 3.2 0.04 1.3  8:1 0.15 2.0, 6.6 1.46 2–3 14  7:1 ... ... ...

NH;h (1022 /cm2 ) 1.8 4.1 3.8 1.1 ...

Th (MK) 69 80 61 62 ...

LdX;h (1029 ) 5.1 10.2 1.1 3.5 ...

X-ray data and AV from [8], [9] or [11] and references therein. Index ‘s’ for soft component, ‘h’ for hard component. b Luminosities in units of 1030 erg s1 from [11] or calculated from [O I]6300 equivalent width ([12] or [11], see [11] for method; for high-velocity component if available). c Mass loss rates from [11] or [8] and references therein. d LX for 0.1–10 keV range, in units of 1029 erg s1 .

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are assumed ([8] and references therein). Because the hard component occasionally flares [8], it is likely to be “coronal”. The excess absorption can be induced by accretion streams falling down along the magnetic fields and absorbing the X-rays from the underlying corona. The excess absorption-to-extinction (NH =AV ) ratio then is an indicator of dust sublimation: the accreting gas streams are dust-depleted. In contrast, NH D 1:1 .0:8–1:4/  1021 cm2 (90% error range) of the soft component is lower than suggested from the stellar AV , NH .AV /  .3–6/  1021 cm2 . A likely origin of these very soft X-rays is the base of the jet [7], suggested by (1) the unusually soft emission compatible with the jet spectrum, (2) the low NH , and (3) the explicit evidence of jets in the Chandra image (see above). Schneider & Schmitt (this volume) further show that the soft component is very slightly offset from the harder, coronal component in the direction of the optical jet.

4 Interpretation and Discussion We estimate physical properties of the X-ray emitting jet plasma following G08 for DG Tau, considering two simple geometric models (see Fig. 3). First, we assume a conical jet opening to a cross-section radius of 0.500 (70 AU) at a distance of 500 from the star. We define the “inner” jet for distances within 0.500 of the star (accessible only spectroscopically), and the “outer” jet for distances of  100 –500 (imaged by Chandra). For the inner jet, we also consider a cylindrical model with a constant cross section radius of R D 1 AU. From this geometry, the projected jet expansion velocity of (0.2–0.3)00 yr1 [2], and the gradient in LX from the inner to the outer jet, we estimate a cooling time t of 2 yrs (inner jet plasma) to 20 yrs (outer jet plasma). The X-ray volume emission measure is EM D ne nH f V  0:83n2e fV

(1)

as determined from the spectra [8, 10], where V is full model jet volume, f is the volume filling factor of the plasma within V , ne and nH are the electron and

Fig. 3 Sketch illustrating adopted jet geometry for “inner” and “outer” DG Tau jet

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Table 2 TAX sources: Interpretation Jet Outflow rate Heating rate (erg s1 ) (Mˇ yr1 ) Outer 3:6  1013 1:7  1028 Inner (conical) 3:0  1012 1:4  1029 Inner (cylindrical) 3:0  1012 1:4  1029

Electron density (cm3 ) 1:7  105 1:7  106 1:7  106

Volume filling factor 2:3  106 1:9  104 3:1  103

Cooling time (yr) 20 2 2

hydrogen number densities, respectively, and the approximation holds for a plasma with solar composition. The total plasma mass is M D 1:28mp .EMf V /1=2

(2)

where mp is the proton mass, and the thermal energy content of this mass is Eth D 3:15.EMf V /1=2 kT:

(3)

If cooling is radiative and the losses correspond to the total radiative loss rate, L, extrapolated from the soft X-ray spectrum (L  1:42LX;0:110keV ), we find f 1=2 D

0:45LX t .EMV /1=2 kT

(4)

and finally the electron density in the X-ray source, ne D

2:4EMkT : LX t

(5)

The resulting parameters are summarized in Table 2. We find very small filling factors for the hot plasma, requiring relatively high densities. Assuming typical densities of 103 –104 cm3 in the cooler gas for the outer jet (T  104 K) and 106 cm3 for the inner jet, we derive an overpressure of the hot plasma of 2–3 dex (inner jet) and 3–5 dex (outer jet). It is therefore likely that the hot plasma, if enclosed by cooler jet gas, contributes to the transverse expansion of the jet. We note that these estimates hold under the assumption of radiative cooling. Additional adiabatic cooling by expansion is likely (see [10]).

References 1. S. A. Balbus, J. F. Hawley, ApJ 376, 214 (1991) 2. C. Dougados, S. Cabrit, C. Lavalley, F. M´enard, A&A 357, L61 (2000) 3. J. Eisl¨offel, R. Mundt, AJ 115, 1554 (1998) 4. F. Favata, C. V. M. Fridlund, G. Micela, et al., A&A 386, 204 (2002) 5. A. E. Glassgold, J. Najita, J. Igea, ApJ 615, 972 (2004) 6. N. Grosso, E. D. Feigelson, K. V. Getman, et al., A&A 448, L29 (2004)

352 7. M. G¨udel, S. L. Skinner, K. R. Briggs, et al., ApJ 626, L53 (2005) 8. M. G¨udel, A. Telleschi, M. Audard, et al., A&A 468, 515 (2007a) 9. M. G¨udel, K. R. Briggs, K. Arzner, et al., A&A 468, 353 (2007b) 10. M. G¨udel, S. L. Skinner, M. Audard, et al., A&A 478, 797 (2008) 11. P. Hartigan, S. Edwards, L. Ghandour, ApJ 452, 736 (1995) 12. G. A. Hirth, R. Mundt, J. Solf, A&AS 126, 437 (1997) 13. S. H. Pravdo, E. D. Feigelson, G. Garmire, et al., Nature, 413, 708 (2001)

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The Complex Morphology of the X-Ray and Optical Emission from HH 154: The Pulsed Jet Scenario Rosaria Bonito, Salvatore Orlando, Giovanni Peres, Fabio Favata, and Jochen Eisl¨offel

Abstract We study the optical and X-ray emission from protostellar jets, focusing, in particular, on the case of HH 154. This project consists of two different and complementary approaches: the development of hydrodynamical models of the jet/ambient interaction, and the analysis of multi-wavelength observations. Comparing the results derived from the simulations with the observations we can infer the physical mechanisms leading to the complex morphology of the X-rays source observed at the base of HH 154.

R. Bonito () INAF-Osservatorio di Palermo-COMETA e-mail: [email protected] S. Orlando INAF-Osservatorio di Palermo G. Peres Universita’ di Palermo-INAF-Osservatorio di Palermo F. Favata ESA, Community Coordination and Planning Office Paris J. Eisl¨offel Th¨uringer Landessternwarte, Sternwarte 5, D-07778 Tautenburg, Germany

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1 X-Ray Emission from Herbig-Haro Objects The Herbig-Haro (HH) objects have been observed in the last 50 years in optical, radio, and IR bands. In the last years, since 2000, also X-ray emission has been detected from a total of six HH objects. The first two X-ray emitting jets are HH 2 [7] and HH 154 [5]. The X-ray emission from the HH jets has been predicted according to the relation between the post-shock temperature and shock velocity (see [2]): in fact for values of the velocity of few hundred km/s we expect temperatures of million degrees, thus leading to X-ray emission. Bonito et al. (2007) [2] in Table 1 summarize the main parameters of the X-ray emitting HH objects. We will focus on HH 154 because: (1) it has been observed with both XMM, for a large effective area, and Chandra, for a good spatial resolution; (2) the strong X-ray emission from the stellar corona suffers a strong absorption, while the X-ray emission from the jet itself suffers a smaller absorption, so there is no contamination of the X-ray source from the jet due to the stellar corona; (3) HH 154 is the nearest most luminous X-ray emitting jet (we can obtain, with a single exposure of less than 100 ks, more than 60 cnts, see [4]).

2 A Simple Model: X-Ray Emission of the Shock from a Continuous Jet Using the FLASH numerical code [6], we solved the equation of conservation of mass, momentum, and energy which describe the jet/ambient system and in which we have taken into account the thermal conduction in Spitzer and saturated regimes, and the optically thin radiative losses. Solving these equations, we derived the bidimensional maps of the temperature and density of the jet/ambient system, the emission measure distribution as a function of the temperature, the absorbed focal plane spectrum, and the X-ray emission maps as described in details in [2]. We performed a wide exploration of the parameters which mainly describe the interaction between a continuous jet and the surrounding medium: the Mach number and the density contrast. A first important result is that only a narrow range of the initial parameters of the jet/ambient system can reproduce the observations (spectral behavior, best-fit temperature, X-ray luminosity, and shock velocity). In the case of HH 154 in particular, the best case is a light jet with high Mach number, corresponding to high initial velocity [1, 2]. Our model predicts a proper motion of the X-ray source of about 500 km/s which is detectable with Chandra/ACIS-I in few years. We also studied the spectral behaviour and the X-ray luminosity of the source and we derived that the best-fit parameters are in nice agreement with the observed data of HH 154. We can conclude that the model of a shock propagating into the surrounding medium reproduces in a natural way the observed properties of the X-ray emitting jets and therefore it is a natural candidate to explain the physical mechanism leading to the observed emission.

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3 Complex Morphology of the X-Ray source from HH 154; Joint Optical/X-Ray Observations Prompted by the results discussed in [1,2], in order to verify our model’s predictions, we performed a joint X-ray (with Chandra) and optical (with Hubble) observation of HH 154 in 2005. Comparing the X-ray source as observed in HH 154 with Chandra in 2001 and in 2005 (see [4]) the first important result is that the X-ray source shows a complex morphology, in fact it consists of two components: a bright, point-like, stationary (over 4 years) source at the base of the jet and an elongated source which shows proper motion, as predicted by our model. This is an important result both because the proper motion of the X-ray source from an HH jet has been observed for the first and only time to date in this case (HH 154) and because the value of the velocity is consistent with the model’s prediction (about 500 km/s, see [2]). Figure 1 shows the joint optical observations with Hubble of HH 154, in particular the difference image in the 2 filters, H˛ and [SII]. We can recognize the terminal working surface, the D knot at the head of the jet, and the so called F-complex at the base of the jet. We suggest the formation of a new shock at the base of the jet, with the same stratification of the H˛ and [SII] emission as detected in the terminal working surface, so we argue that we can identify this as an internal working surface. Most important this is coincident with the position of the X-ray source, as discussed in details in [3]. In conclusion, the 2005 joint optical and X-ray observations show new features of the X-ray source from protostellar jets: (1) the X-ray emission is located at the base of the optical jet; (2) it shows a complex morphology; (3) there is evidence of a variability of the morphology on the small time-scale of few years.

4 The Pulsed Jet Scenario Explaining the X-Ray Morphology Our simple model discussed in Sect.2 explains the spectral behaviour of the X-ray source, its best-fit temperature and emission measure, its luminosity and velocity,

N

new shock? E

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however it fails in explaining some new features observed in 2005. Therefore we need a new model in order to explain the location of the X-ray source at the base of the jet (which appears as a common feature in the HH jets), the complex morphology observed, and its variability in few years. Up to now, these features have been observed for the first and only time in HH 154. This is the only jet which allows us to perform spatially resolved study of the structure of its X-ray emission, being the nearest most luminous X-ray emitting jet. We use the same basic physics of the simple continuous model, because it explains in a natural way the X-ray emission from HH jets but assuming a time variable initial jet velocity. The variable velocity may be related to episodic phenomena in the accretion process. From the model results we synthesized the X-ray emission. We are performing the exploration of the parameter space, the Mach number, the density contrast, but also the distribution of the jet’s velocity: sinusoidal, exponential, power law. Here we discuss the case of an exponential distribution with few blobs ejected at high speed and most of the blobs ejected at low speed thus leading to self-interactions due to collisions between blobs ejected with different velocity in different epochs. We discuss in detail one case, a light jet with high maximum initial velocity. We can recognize the self-interaction phenomena leading to the formation of hot and dense blobs within the jet in the temperature and density maps of the model of a pulsed jet (see Fig. 2). Figure 3 shows the evolution in 10 years of the X-ray emission synthesized from our model. Two consecutive images correspond to a time lapse of 1 year. Although there is still X-ray emission from the head of the jet, it becomes fainter and fainter with time. This image shows that the pulsed model allows to have X-ray emission also from the base of the jet, with a complex morphology and variability in few years, as observed. We integrated the X-ray emission over the radial direction and the time obtaining the cnts along the axis of the jet and we conclude that the X-ray emission is located mainly at the base of the jet. Figure 4 shows the X-ray maps as we predict should be observed with ACIS-I, with the same resolution, derived from our model: the X-ray source modeled is complex and variable in few years, and the dimension of the source is consistent with the X-ray source observed with Chandra (on the left). We conclude that our model appears to be promising in order to explain the mechanism leading to the complex morphology of the fast-variable X-ray source detected at the base of the protostellar jets. Acknowledgements The software used in this work was in part developed by the DOE-supported ASC/Alliances Center for Astrophysical Thermonuclear Flashes at the University of Chicago, using modules for thermal conduction and optically thin radiation constructed at the Osservatorio Astronomico di Palermo. This work makes use of results produced by the PI2S2 Project managed by the Consorzio COMETA, co-funded by the MIUR. This study was supported in part by the European Community’s Marie Curie Research and Training Network JETSET (Jet Simulation, Experiments and Theory) under contract MRTN-CT-2004 005592. S.O., and G.P. acknowledge support from the Marie Curie Fellowship Contract No. MKT-CT-2005-029768.

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Fig. 2 Bi-dimensional cuts of the temperature (left semi-panels) and density (right semi-panels) maps of the pulsed jet model for three stage of evolution: the self-interaction between different pulses are highlighted by the ellipses. A color version of this plot can be found in the electronic version of this book and in Appendix A (Fig. A.24)

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Fig. 4 Comparison between the X-ray emission maps derived from the pulsed jet model and the ACIS observation of HH 154 in 2005. A color version of this plot can be found in the electronic version of this book and in Appendix A (Fig. A.26)

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References 1. Bonito, R., Orlando, S., Peres, G., Favata, F., and Rosner, R. 2004, A&A, 424, L1 2. Bonito, R., Orlando, S., Peres, G., Favata, F., and Rosner, R. 2007, A&A, 462, 645 3. Bonito, R., Fridlund C. V. M., Favata F., Micela G., Peres G., Djupvik A. A., and Liseau R. 2008, A&A 484, 389 4. Favata, F., Bonito, R., Micela, G., Fridlund, M., Orlando, S., Sciortino, S., and Peres, G. 2006, A&A, 450, L17 5. Favata, F., Fridlund, M., Micela, G., Sciortino, S., and Kaas, A. 2002, A&A, 386, 204 6. Fryxell, B., Olson, K., Ricker, P., Timmes, F. X., Zingale, M., Lamb, D. Q., MacNeice, P., Rosner, R., Truran, J. W., and Tufo, H. 2000, A&A, 131, 273 7. Pravdo, S. H., Feigelson, E. D., Garmire, G., Maeda, Y., Tsuboi, Y., and Bally, J. 1989, Nature, 413, 708

Radiative Shocks in the Context of Young Stellar Objects: A Combined Analysis from Experiments and Simulations Chantal Stehl´e, Matthias Gonz´alez, Edouard Audit, and Thierry Lanz

Abstract Hypersonic flows occurring during stellar formation are structured by radiation. When radiation is reabsorbed in different locations of the accretion flows or in the bow shocks of highly supersonic jets, the coupling between hydrodynamics and radiation becomes an important feature that significantly affects the hydrodynamical structure as well as spectroscopic signatures of these shocks.

C. Stehl´e () LERMA, Observatoire de Paris, CNRS and UPMC, 5 Place Jules Janssen 92195 Meudon, France e-mail: [email protected] M. Gonz´alez Instituto de Fusi´on Nuclear Universidad Polit´ecnica de Madrid - ETSII, calle Jos´e Guti´errez Abascal 2 28006 Madrid, Spain e-mail: [email protected] E. Audit, Service d’Astrophysique, DSM/IRFU/SAp, CEA/Saclay, 91191 gif-sur-yvette Cedex, France e-mail: [email protected] T. Lanz Department of Astronomy, University of Maryland, College Park, MD 20742 and LERMA, Observatoire de Paris, CNRS and UPMC, 5 Place Jules Janssen 92195 Meudon, France e-mail: [email protected]

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Radiative shocks can now be studied in the laboratory using high-energy lasers. Using recent laboratory experiment and state-of-the-art multi-dimensional radiative hydrodynamics simulations, we present an up to date description of the physical and hydrodynamical properties of radiative shocks, with an emphasis on the aspects that are important for stellar hypersonic flows.

1 Introduction Shocks are common events in astrophysical context. We are here interested in strong hydrodynamical shocks. For such shocks, propagating with a velocity v (in an hydrogen plasma) the temperature of the shocked gas is equal to TH D 1:5 105 

 v 2 .K; km=s/ 100

(1)

Shocks with typical velocities of several hundreds km/s are frequent in stellar jets of Young Stellar Objects (YSOs) and in the accretion flows connecting their disk to the stellar photosphere. These shocks radiate, in particular in X-rays [13]. The effect of the radiation on the flow is strongly dependent on the opacity, leading to distinguish between two classes of shocks: optically-thin shocks where the radiation flux induces an energy sink in the hydrodynamical equations, and optically-thick shocks characterized by the presence of a radiative precursor. In this latter case, the photons emerging from the hot compressed flow ionize the initially cool medium in which the shock propagates, generating a radiative precursor in front of the shock discontinuity. The structure of radiative shocks is imprinted by NLTE effects (excitation, desexcitation or chemical reactions) and their different associated time scales [4]. Depending on the velocity, the shocks models are classified as subcritical or supercritical. In this later case, the temperature remains constant through the shock front. The critical velocity that separates the two regimes is obtained when the radiative flux equals the internal energy flux entering the front, T 4 D 0 u. 0 ; T /, yielding (2) ucrit;H ' 500 1=5 Πkm=s; g=cm3  A typical structure of such supercritical shock is given in Fig. 1. Modeling hypersonic shocks is a difficult task that requires to solve simultaneously the hydrodynamics equations with radiation flux, energy and pressure terms, the radiative transfer equations, and population equations for the different species. Simplifications are thus often made, though their validity can hardly be assessed because of the lack of detailed observations of these strong shocks (insufficient resolution or embedded objects). In the context of YSOs, radiative shocks have been investigated by Raga [14] who showed that jets with velocities larger than 200 km/s present radiative precursors at the head of the jet. In their simulations, the jet (n  225 cm3 ) was propagating in an underdense medium (n  25 cm3 ). The ionisation effects induce not only

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Fig. 1 Typical temperature and radiative flux in a supercritical experimental radiative shock. The radiative flux is null in the optically-thick shocked gas

the presence of a radiative precursor, but also a reduction by a factor of about 2 of the maximum temperature (4 105 K instead of 8 105 K), with noticeable changes in the X-ray emission of jets faster than 200 km/s. The presence of a radiative precursor is also of interest in accretion flows connecting the accretion disk to the photosphere of young stars [17]. Connected to the opacity, the location of the accretion shock [9] is still under debate for YSOs. Another open question concerns the topology and the stationarity of accretions flows [10]. Experimental studies on high-energy lasers and Z-pinches therefore present a novel approach to investigate these hypersonic flows that may shed new light on the complex case of YSOs.

2 Experimental Radiative Shocks In radiative shock experiments on high-energy laser installations, the laser pulse with a typical duration of 1 ns or less is focalized on a foil closing a shock tube of millimetric scale. The interaction of the laser with the foil allows to drive a shock wave in the gas. Shock velocities range from 60 km/s up to 150 km/s depending on the laser energy (60 J to 1500 J). For a gas of atomic number A, the temperature of the shocked gas, TA D A  TH , where TH is given by (1); the critical velocity, ucrit;A D ucrit;A  A4=5 , where ucrit;H is obtained from (2). It is therefore easier to create strong supercritical shocks in high-Z gas and at low density. Most experiments have thus been performed in xenon (A D 131:29) at a pressure of a fraction of a bar ( ' 105 g/cm3 at room temperature). The experiments have initially been focussed on the shock launching phase (0–10 ns) where the precursor is faster than the shock front [1, 2, 5]. The shock tube is chosen as simple as possible, with a square section of less than 1 mm2 and a length of a few millimeters (Fig. 2). Time resolved visible interferometry demonstrated the existence of a radiative precursor, with electron density in qualitative agreement with 1D simulations [1]. Visible Thomson scattering also showed that

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Fig. 2 Cell used in a radiative shock experiment [7]. The dimensions of the glass tube are 0:7  0:7  3 mm. The piston is made from a 10 microns gilted plastic. The laser is coming from the top

the electron temperature in the shock front is lower than the ion temperature, as predicted by simulations [16, 18]. Moreover,  the shock front is slightly curved (X-ray radiography, [15]), and  the precursor velocity is smaller than expected from 1D simulations and deceler-

ates more quickly with time than predicted (visible interferometry, [3, 7]). Hydrodynamically, the tube may be considered as a 1D medium. Therefore, these last results have been attributed to 2D radiation effects. More specifically, depending on the transverse optical depth compared to unity, the photons emitted in the shock front may be lost in the tube walls and thus do not heat the gas [7, 11]. The latest experiments performed with the PALS laser in Prague were devoted to study the dynamics of radiative shocks in xenon over longer durations (50 ns instead of 10 ns). The new experiments showed that the dynamics of the radiative shocks is strongly reduced compared to 1D simulations. The combined effects of the radial losses and walls albedo [3, 7], defined as the proportion of radiative flux that is reflected on the tube surface, explain the differences with 1D simulations and are described in the next Section.

3 Modeling Laboratory Radiative Shocks We used HERACLES [6], a 2D radiation hydrodynamics code to model these radiative shocks. The shock propagates in xenon gas (initial conditions 0 =1.03 103 g.cm3 and room temperature) in a 3D cell .0:7 mm  0:7 mm  4 mm/. The piston velocity is chosen to reproduce the experimental velocity. Simulations reproduce well the dynamics of the precursor over 40 ns using an albedo of 40%, which corresponds to typical values expected for glass windows [7]. The effect of the albedo on the dynamics of the position of the precursor front is shown in Fig. 3. At early times, the precursor develops before the shock and its velocity is larger than the front velocity. After this launching phase, the precursor slows down and its extension reaches a maximum at about 5 ns for a D 0%, 10 ns for a D 50% and at times of more than 200 ns for a D 100% [8]. The velocity decreases until it becomes equal to the shock front velocity, which corresponds to the stationary limit. This results is important for astrophysical situations: it shows that the transverse optical depth has an important influence on the stationarisation time.

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Fig. 3 Position versus time of the precursor front relative to the position of the shock front for different values of the albedo : 0 % (dashed line), 50 % (dotted line) and 100 % (full line). The conditions are close to experimental shocks in xenon

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The value of the wall albedo will also strongly affect the angular distribution of the radiative flux. We performed 2D calculations of the energy flux L that emerges from the shock tube for different albedos, as a function of the angle ˛ relative to the direction of the shock velocity. The calculation is performed at a fixed time and the flux is integrated from the shock front up to the end of the shock tube. We found an anisotropic flux distribution, with a prefered direction almost colinear with the shock propagation for large values of a, and perpendicular for a D 0% (transparent windows) [8] (Fig. 4-a). This anisotropy may therefore induce specific variability for CTT accretion shocks (mimicked by a D 0%) in magnetospheric funnels correlated to the variable angle of observation of the shock during the stellar rotation. The laboratory case is close to a ' 50%, and one expects that the maximum of the smooth flux distribution L.˛/ to be peaked around ˛ ' =4, with still a noticeable contribution at ˛ D 0 (observation from the rear face of the tube). However, these computations have been performed in the grey approximation for the radiation. As we argued that the lateral losses depend on the transverse opacity, we expect that the flux distribution L.˛/ will be frequency-dependent, as illustrated in Fig. 4b that shows the predicted shock spectrum in the rear face of the tube. Strong

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absorption is seen around 10–40 eV in the spectrum where the black body spectrum would be near maximum. This absorption is the physical mechanism that is responsible for the development of the precursor.

4 Conclusions Radiative shocks are complex phenomena that require multidimensional radiative hydrodynamical simulations. Experiments have shown that the time needed to reach the stationary regime is affected by transverse radiation transport. The net frequency averaged radiative flux angular distribution is anisotropic and frequency-dependent. All these conclusions provide crucial insight for a realistic modeling of the stellar variability associated to accretion processes, especially in the magnetospheric funnels of T Tauri stars. Acknowledgements We acknowledge financial supports from the Access to Research Infrastructures activity in the Sixth Framework Programme of the EU (contract RII3-CT-2003-506350, Laserlab Europe), from the RTN JETSET (contract MRTN-CT-2004 005592), the ANR grant SiNERGHy (ANR-06-CIS6-009-01), and by french CNRS program PNPS. M.G. acknowledges the financial support provided by the European Commission TUIXS project and the French Ministry of Foreign Affairs through the Lavoisier grant.

References 1. Bouquet S., Stehl´e C., Koenig M. et al.: High power laser radiative Shock for laboratory Astrophysics, PRL 92, 225001–225004 (2004) 2. Bozier J.C., Thiell G., Le Breton J.P. et al.: Experimental Observation of a radiative Wave Generated in Xenon by a Laser-Driven Superciritcal Shock, PRL, 57, 1304–1307 (1986) 3. Busquet M., Audit E., Gonz´alez M. et al.: Effect of lateral radiative losses on radiative shock propagation, HEDP, 3, 8–11 (2007) 4. Draine B.T, McKee C.: Theory of intersellar shocks, Ann. Rev. Astron. Astrophys., 31, 373– 432 (1993) 5. Fleury X., Bouquet S., Stehl´e C. et al.: A laser experiment for studying radiative shocks in astrophysics, Las. Part. Beams, 20, 263–268 (2002) 6. Gonz´alez M., Audit E., Huynh P.: HERACLES: a three-dimensional radiation hydrodynamics code, A&A, 464, 429–435 (2007) 7. Gonz´alez M., Stehl´e C., Audit E. et al.: Astrophysical radiative shocks: From modeling to laboratory experiments, Las. Part. Beams, 24, 535–540 (2006) 8. Gonz´alez M., Audit E., Stehl´e C.: , submitted (2008) 9. Guenther H.M., Schmitt J.H.M.M., Robrade et al.: X-ray emission from claissical T tauri stars: accretion and coronae ?, A&A 466, 1111–1121 (2007) 10. Kulkarni A.K., Romanova M.M.: Accretion to magnetized stars through the Rayleigh-Taylor instability: global 3D simulations, Mon. Not. R. Astron. Soc., 386, 673–687 (2008) 11. Leygnac S., Boireau L., Michaut C. et al.: Modeling multidimensional effects in the propagation of radiative shocks, Phys. Plasmas, 13, 113301-113301-10 (2006) 12. Michaut C., Stehl´e C., Leygnac S. et Lanz T.: jump conditions in hypersonic shocks: Quantitative effects of ionic excitation and radiation, Eur. Phys. J. D 28, 381–392 (2004)

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13. Pravdo S.H., Tsuboi Y., Maeda M., X-rays from HH80, HH81 and the central region, APJ, 605, 259–271 (2004) 14. Raga A.C., Mellema G., Arthur S.J. et al.: 3D Transfer of the Diffuse Ionizing Radiation in ISM Flows and the Preionization of a Herbig-Haro Working Surface, Rev. Mex. Astron. Astrofis., 35, 123–133 (1999) 15. Reighard A.B., Drake R.P., Donajkowski T. et al.: Thomson scattering from a shock front, Rev. Sc. Inst., 77, 10E504-10E504-3 (2006) 16. Reighard A.B., Drake R.P., Dannenberg K.K. et al.: Observation of collapsing radiative shocks in laboratory experiments, Phys. Plasmas, 13, 082901-082901-5 (2006) 17. Stahler S.W., Shu F.H., Taam R.E.: The evolution of protostars. II - The hydrostatic core, ApJ, 242, 226–241 (1980) 18. Stehl´e C., Chi`eze J.P.: Radiative Shocks in Astrophysics: from the experiment to the modelisation, in the proceeding of SF2A, Eds.: F. Combes and D. Barret, EdP-Sciences (Editions de Physique), Conference Series, p. 493–494 (2002)

X-Ray Imaging Spectroscopy of Planetary Nebulae in the Chandra/XMM Era: New Insight into Stellar Jets Joel H. Kastner

Abstract Imaging spectroscopy with the Chandra and XMM-Newton X-ray Observatories can provide unique insight into the role of collimated outflows and jets as the shaping agents of planetary nebulae (PNe). Diffuse X-ray sources within PNe serve as probes of wind-collision-generated shocks that determine PN structure and drive PN evolution. In certain cases, high-resolution X-ray imaging of a PN can reveal collimated flows whose presence would otherwise only be inferred indirectly. Meanwhile, detections of X-ray point sources at the cores of PNe serve to constrain mass outflow launching and collimation models invoking accretion disks and/or magnetic fields. Results obtained for the young, rapidly evolving PN NGC 7027 and the bipolar nebulae Mz 3 and Hb 5 are particularly illustrative of the potential of X-ray imaging spectroscopy to elucidate the role of disks and jets in shaping PNe.

J.H. Kastner () Laboratoire d’Astrophysique de Grenoble, Universit´e Joseph Fourier — CNRS, BP 53, 38041 Grenoble Cedex, France e-mail: [email protected]

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1 Introduction Planetary nebulae (PNe), the near-endpoints of stellar evolution for intermediatemass (1–8 Mˇ ) stars, present a dizzying variety of optical and near-infrared morphologies: round; elliptical; profoundly bipolar; highly point-symmetric; chaotic and clumpy. The physical mechanisms responsible for this morphological menagerie have been the subject of hot debate among PN researchers over the past two decades (e.g., [3,5,8]). A coherent explanation for the transformation from a quasi-spherical wind during the progenitor star asymptotic giant branch (AGB) phase to highly collimated outflow during the PN phase has been particularly elusive. One hypothesis for such a transition is that just before or during PN formation (when the newly exposed core ionizes the ejected AGB envelope), a disk is present at the PN core. In the realm of young stellar objects (YSOs), of course, such a hypothesis (i.e., invoking a circumstellar disk to explain a jet) would not raise a theorist’s (or even observer’s) eyebrows — even though, as this conference readily demonstrated, there remains considerable disagreement concerning the details. In contrast, both the theoretical motivation and observational evidence for the mere presence of disks and jets within PNe is less obvious. Although the frequently observed bipolar and point-symmetric PN morphologies and accompanying kinematic signatures directly implicate jets as the origin of PN outflow collimation, there is a priori no reason to expect an accretion disk to form around a single, mass-losing AGB or post-AGB star. Hence, models of PN structural evolution and outflow collimation typically invoke either the presence of an accreting binary companion to the mass-losing primary (e.g., [14] and references therein), magnetic fields at the mass-losing star, or some combination of the two (e.g., [4] and references therein). However, it is exceedingly difficult to test such models given the large distances, central star luminosities, and visual extinctions typical of rapidly evolving PNe. The combination of high spatial resolution and moderate spectral resolution X-ray imaging made possible by the Chandra and XMM-Newton X-ray Observatories is providing such much-needed contsraints on PN shaping models (for a recent review, see [11]). I describe here specific examples of Chandra and XMM X-ray imaging spectroscopy (including a newly discovered, serendipitous PN X-ray source) that have shed new light on dynamical processes within PNe whose structure and gas kinematics are indicative of high-velocity, collimated flows.

2 Tracing a Transformation: NGC 7027 The well-studied young PN NGC 7027 (Fig. 1) is evidently undergoing a rapid transition from spherical to collimated outflow. Hubble Space Telescope (HST) optical/IR images of this PN show what appear to be multiple, well-collimated flows or “shrapnel” superimposed on a system of concentric rings surrounding a bright, central, elliptical shell (Fig. 1, left). The concentric rings most likely represent vestiges

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Fig. 1 HST optical (left) and Chandra (middle) X-ray images of NGC 7027. At upper right is the 0.3–2.5 keV spectral energy distribution extracted from the Chandra X-ray CCD data; at lower right, contours of high-velocity Br emission are superimposed on the Chandra X-ray image. Adapted from [9] and [6]. A color version of this plot can be found in the electronic version of this book and in Appendix A (Fig. A.27)

of “slow” (15  30 km s1 ), episodic, spherical mass loss by the AGB progenitor, while the elliptical shell (with its apparent equatorial density enhancement) may be “last-gasp” AGB or post-AGB star material that has been swept up by a rapidly accelerating stellar wind emanating from newly exposed hot layers within the central star. The nature of the apparent high-velocity shrapnel is somewhat less obvious. However, Chandra imaging revealed a soft (0.3–2.0 keV) X-ray-emitting region within NGC 7027 whose morphology closely matches that of the shrapnel-like “blowouts” detected in HST optical/IR imaging of this PN (Fig. 1, left and middle; [9]). This morphological correspondence, combined with the temperature of the superheated plasma inferred from X-ray CCD spectroscopy (3  106 K; Fig. 1, upper right), indicates that the X-rays are produced via shocks as high-velocity (300 km s1 ) outflows — jets or bullets — puncture and disrupt the elliptical shell and spherical rings. This hypothesis is supported by velocity-resolved IR imaging of Br emission, which reveals a close correspondence between the morphology of the highest-velocity ejecta and that of the brightest X-ray emission (Fig. 1, lower right; [6]). Hence, in NGC 7027, X-ray imaging provides direct evidence for the dramatic structural transformation of a PN, presumably brought on by the catastrophic disintegration of its central star(s).

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3 Revealing Shaping Agents: Menzel 3 and Hubble 5 Observations of extreme bipolar PNe — i.e., nebulae with “pinched waists” and well-defined polar lobes — should provide particularly stringent tests of models of PN shaping processes that invoke the influence of binary companions, magnetic fields, and/or accretion disks (e.g., [3]). However, the detections of wind-shockheated, diffuse plasmas within PNe obtained thus far via targeted observations by Chandra and XMM include only one example of an extreme bipolar nebula: Menzel 3 (Mz 3; [10]). This object is, in fact, a suspected symbiotic binary system (see review by Sokoloski in these proceedings). Chandra observations of Mz 3 revealed soft (0.3–1.5 keV) X-rays from apparent jets and relatively hard (1.5–3.0 keV) emission from a point-like central source (Fig. 2, left). This X-ray imaging spectroscopy thereby unveiled the collimated, high-velocity flow that is presently “inflating” the twin bipolar bubbles of Mz 3 (see also [2]) and, furthermore, yielded strong support for a model in which the nebula’s central collimating agent is an accretion disk around a heretofore undetected, compact secondary star. The PN Hubble 5 (Hb 5) is another photogenic example of a classical pinchedwaist bipolar nebula (Fig. 2, right). Its bipolar lobe morphology is similar to that of Mz 3, but Hb 5 also displays striking point-symmetry, circularly symmetric rings (a la NGC 7027), and a “skirt” of bubble-like protrusions around its waist that is reminiscent of the  Car nebula. Little information is available concerning the central star(s) of Hb 5 apart from the high excitation state of the nebula, which implies that the PN harbors a very hot ( 105 K) white dwarf [7]. However, as noted, the bipolar structure of this PN — like that of Mz 3 — appears to require that the PN core consists of a binary system in which an accretion disk has formed. As part of a larger program to analyze serendipitous X-ray observations of PNe (see [12]), we have determined that an X-ray source is spatially coincident with Hb 5 (Fig. 2, right; [13]). Although these archival XMM data are compromised by several factors — in particular, the limited spatial resolution of XMM relative to Chandra combined with the large off-axis angle of the Hb 5 X-ray source — the serendipitous

Fig. 2 Color composites of HST and X-ray (Chandra and XMM, respectively) images of the bipolar nebulae Menzel 3 (left) and Hubble 5 (right). The HST images are color-coded blue-green, and the X-ray images are color-coded red. North is up and east is to the left in both images; the fields of view are approximately 2000  2000 (Mz 3) and 7500  7500 (Hb 5). Adapted from [10] and [13]. A color version of this plot can be found in the electronic version of this book and in Appendix A (Fig. A.28)

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XMM image and spectrum of Hb 5 are indicative of diffuse emission from hot ( 3  106 K) plasma surrounding a compact, central X-ray source at the nebular waist. Targeted (on-axis) Chandra imaging spectroscopy of Hb 5 is required to establish whether the X-rays emanate from collimated flows (as opposed to more widely distributed, superheated plasma) and whether Hb 5 harbors an X-ray-luminous central engine similar to that within (the suspected symbiotic system) Mz 3.

4 Conclusions The cases of NGC 7027, Mz 3, and Hb 5 illustrate the potential of Chandra and XMM imaging spectroscopy to elucidate the processes responsible for the structural evolution of PNe and, in particular, to constrain models describing PN-shaping jets (e.g., [1]) as well as to provide diagnostics concerning jet driving sources. Additional X-ray observations targeting the youngest PNe as well as extreme, pinched-waist bipolar nebulae are needed, if we are to ascertain the frequency of appearance of jets within rapidly evolving PNe and to better understand the role of binaries, accretion disks, and jets in shaping PNe. RIT imaging science students Geoff Franz and Rudy Montez kindly generated the color composite images of Mz 3 and Hb 5. Support was provided by the Laboratoire d’Astrophysique de Grenoble and NASA Astrophysics Data Analysis award NNX08AJ65G to RIT.

References 1. Akashi, M., Meiron, Y., & Soker, N. 2008a, New Astron., 13, 563 2. Akashi, M., & Soker, N. 2008b, MNRAS, submitted (astro-ph/0805.2332) 3. Balick, B. & Frank, A. 2002, ARAA, 40, 439 4. Blackman, E.G., & Nordhaus, J.T. 2007, in Asymmetrical PNe IV, in press (astroph/0708.4199) 5. Corradi, R.L.M., Manchado, A., & Soker, N., eds. 2008, Asymm. Planetary Nebulae IV, in press 6. Cox, P., Huggins, P.J., Maillard, J.-P., et al. 2002, A&A, 384, 603 7. Gesicki, K. & Zijlstra, A.A. 2007, A&A, 467, L29 8. Kastner, J.H., Soker, N., & Rappaport, S.A., eds. 2000, Asymm. Planetary Nebulae II, ASP Conf. Ser., Vol. 199 9. Kastner, J.H., Vrtilek, S.D., Soker, N., 2001, ApJ, 550, L189 10. Kastner, J.H., Balick, B., Blackman, E.G., et al. 2003, ApJ, 591, L37 11. Kastner, J.H. 2008, in Asymmetrical PNe IV, in press (astro-ph/0709.4136) 12. Kastner, J.H., Montez, R., Balick, B., & de Marco, O. 2008, ApJ, 672, 957 13. Montez, R., Kastner, J.H., Balick, B., & Frank, A. 2008, ApJ, submitted 14. Soker, N. & Rappaport, S.A. 2000, ApJ, 538, 241

Part VI

Molecular Outflows and Turbulence Injection by Jets

3D Modeling of the 2006 Nova Outburst of RS Ophiuchi: Collimated Outflows and Jet-Like Ejections Salvatore Orlando, Jeremy J. Drake, and J. Martin Laming

Abstract Chandra/HETG observations of the recurrent nova RS Ophiuchi at day 13.9 of its 2006 outburst reveal a spectrum covering a large range in plasma temperature and characterized by asymmetric and blue-shifted emission lines, suggesting a jet-like ejection, collimated by the central binary. We investigate the origin of line asymmetries by performing 3-D hydrodynamic simulations. We found that our model reproduces the observed X-ray emission in a natural way if an equatorial density enhancement is taken into account. The asymmetric nature of the circumstellar medium into which the early blast wave is driven leads to the shock collimation in the plane of the sky. Most of the early X-ray emission originates from a small region

S. Orlando () INAF - Osservatorio Astronomico di Palermo “G.S. Vaiana”, Piazza del Parlamento 1, 90134 Palermo, Italy; Consorzio COMETA, via Santa Sofia 64, 95123 Catania, Italy e-mail: [email protected] J.J. Drake Harvard-Smithsonian Center for Astrophysics, 60 Garden Street, Cambridge, MA 02138, USA e-mail: [email protected] J.M. Laming Space Science Division, Naval Research Laboratory, Code 7674L, Washington DC 20375, USA e-mail: [email protected]

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propagating in the direction perpendicular to the line-of-sight and localized at the interaction front between the blast wave and the equatorial density enhancement. The model predicts asymmetric and blue-shifted line profiles remarkably similar to those observed and explains the asymmetries as due to substantial X-ray absorption of red-shifted emission by ejecta material.

1 Introduction RS Ophiuchi (RS Oph) is a symbiotic recurrent nova that went into its latest outburst on 2006 February 12.83 UT [9]; previous outbursts were recorded in 1898, 1933, 1958, 1967, and 1985. RS Oph is thought to be a binary system, comprising a red giant star that does not fill its Roche lobe, and a white dwarf of mass near the Chandrasekhar limit [3, 5, 16]. The outbursts occur on the white dwarf due to thermonuclear runaway of hydrogen-rich material transferred from the companion red giant onto the surface of the white dwarf (e.g. [7]). During the 2006 outburst an intensive international observing campaign was organized, incorporating observations ranging from radio to X-ray wavelengths, and monitoring the outburst since the early phases of its evolution. Chandra/HETG observations at day 13.9 revealed a rich spectrum of emission lines indicative of emitting plasma with temperatures ranging between 3 and 60 MK [4, 10]. Drake et al. [4] noted that the lines are too strongly peaked to be explained by a sphericallysymmetric shock, suggesting a collimation mechanism of the X-ray emitting plasma in the direction perpendicular to the line-of-sight. The lines also appear asymmetric and slightly blue-shifted, while the red wings of the line profiles become weaker with increasing wavelength. Drake et al. [4] suggested that the asymmetric nature of the circumstellar medium (CSM) in which the explosion occurred can be responsible for both the broad range in plasma temperature and the shock collimation observed. Here we investigate the origin of the line asymmetries, broadening and blueshifts observed with Chandra/HETG (see [13] for more details). We aim at exploring possible diagnostics of the early blast wave and of the inhomogeneous CSM in which the explosion occurred. To this end, we model the expansion of the blast wave from the 2006 outburst of the recurrent nova RS Oph through the extended outer atmosphere of the companion red giant, using detailed 3-D hydrodynamic simulations. From the simulations we synthesize the X-ray emission and derive the spectra as they would be observed with Chandra/HETG.

2 The Model We model a blast wave occurring on the white dwarf and propagating through the extended outer atmosphere (the wind) of the companion red giant (see also [13]). The blast is off-set from the origin of the wind density distribution by 1.5 AU (i.e. the

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system orbital separation; [3]). In addition to the r 2 wind density distribution, we include also an equatorial density enhancement (hereafter EDE) in the red giant wind, as suggested by VLBA radio synchrotron observations [11] and by HST observations [1] of the 2006 blast wave. The blast wave is modeled by numerically solving the time-dependent fluid equations of mass, momentum, and energy conservation in a 3-D Cartesian coordinate system .x; y; z/. The model takes into account the radiative losses from an optically thin plasma (e.g. [15, 8]) and the thermal conduction [17], including the effects of heat flux saturation (e.g. [2]). The calculations described in this paper were performed using FLASH, an adaptive mesh refinement multiphysics code [6]. The code has been extended with additional computational modules to handle radiative losses and thermal conduction (see [14] for the details of the implementation). From the simulations, we synthesize the X-ray emission in the Chandra/HETG band originating from the blast wave, applying a methodology analogous to the one described by Orlando et al. [12] in the context of the study of supernova remnants. The synthesis takes into account the thermal broadening of emission lines, the Doppler shift of lines due to velocity components along the line-of-sight (LoS), and the absorption due to shocked CSM and ejecta (assuming that the metallicity of ejecta is 10 times higher than that of CSM). We also assume that the plane of the orbit of the system is inclined by 35ı to the LoS according to the value estimated from observations before the 2006 outburst [3].

3 Results We explored the parameter space of the model, considering two sets of simulations, either with or without the EDE. In all the models examined, we found the typical evolution of radiative shocks propagating through an inhomogeneous medium: the fast expansion of the shock front with temperatures of several millions degrees and the development of dense and cold regions dominated by radiative cooling, as the forward and reverse shocks progress through the CSM and ejecta, respectively. As an example, Fig. 1a shows a 2-D section in the .x; z/ plane of the distribution of mass density for the most promising case (our “best-fit” model), including the EDE, at day 13.9. We found that the shock morphology is quite complex, aspherical, and originates from the propagation of the shock through both the off-set red giant wind and the density enhancement at the equatorial plane. The EDE determines the collimation of the blast wave perpendicularly to the plane of the orbit of the central binary system and leads to a bipolar shock morphology, reminiscent of the morphology observed in the radio and optical bands, distorted (by the off-set red giant wind) and converging on the side away from the red giant. From the simulations, we synthesized the X-ray emission in the Œ0:6  12:4 keV band to match the observing range of Chandra. We found that most of the X-ray emission originates from a jet-like structure with a size of 10 AU propagating upward in the plane of the orbit (see Fig. 1b). The X-ray emission is due to the

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interaction between the blast wave and the EDE and the X-ray source appears to propagate perpendicularly to the LoS. We derived the X-ray spectrum as it would be observed with Chandra/HETG and compared the observed line profiles with the synthetic ones. The synthetic spectrum shows emission lines from different elements, forming over a wide range of plasma temperatures, and reflecting the broad nature of the plasma temperature distribution. We analyzed the line profiles of the most prominent spectral lines to investigate the origin of the broadening and asymmetries revealed in the Chandra/HETG observations (see Fig. 2). We found that the synthetic line profiles are more peaked than expected for a spherically symmetric shock, as found for observed lines [4]. In fact, our model predicts that most of the X-ray emission originates in a compact region propagating away from the red giant in the direction perpendicular to the LoS. We also found that the line profiles are asymmetric and slightly blue-shifted in nice agreement with the observations. Our model shows that the asymmetries are due to substantial X-ray absorption of red-shifted emission by ejecta material. Finally, we found that both shocked ejecta material (light gray lines) and shocked CSM (gray lines) contribute to the observed emission with relative weight depending on the wavelength and with the ejecta component more blue-shifted than the CSM component.

4 Summary and Conclusions We have investigated through detailed hydrodynamic modeling the origin of asymmetries and broadening of the emission lines observed with Chandra/HETG during the 2006 outburst of RS Oph (see [13] for more details). We found that the simulated

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nova remnant is highly aspherical: the blast wave is efficiently collimated by the inhomogeneous CSM. The model reproduces the observed X-ray emission in a natural way if the CSM is characterized by an equatorial density enhancement, as suggested by radio and optical observations. Most of the early X-ray emission originates from a small region localized at the interaction front between the blast wave and the EDE and propagating in the direction perpendicular to the line-of-sight. As a result, the line profiles are peaked as those observed. Our best-fit model predicts asymmetric and blue-shifted line profiles remarkably similar to those observed. We found that the observed asymmetries are due to substantial X-ray absorption of red-shifted emission by ejecta material. Acknowledgements S.O. acknowledges support by the Marie Curie Fellowship Contract No. MTKD-CT-2005-029768 of the project “Young stellar objects, their surroundings and jets: Advanced observational and MHD studies”. This work was supported in part by the Italian Ministry of University and Research (MIUR) and by Istituto Nazionale di Astrofisica (INAF). The software used in this work was in part developed by the DOE-supported ASC/Alliance Center for Astrophysical Thermonuclear Flashes at the University of Chicago, using modules for thermal conduction and optically thin radiation built at the Osservatorio Astronomico di Palermo. The simulations were executed on the Grid infrastructure of the Consorzio COMETA. This work makes use of results produced by the PI2S2 Project managed by the Consorzio COMETA, a project cofunded by the Italian Ministry of University and Research (MIUR) within the Piano Operativo Nazionale “Ricerca Scientifica, Sviluppo Tecnologico, Alta Formazione” (PON 2000-2006). More information is available at http://www.pi2s2.it and http://www.consorzio-cometa.it.

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Molecular Outflows: Observations Rafael Bachiller

Abstract The wealth of mm-wave observations being provided by large radiotelescopes and interferometers on bipolar molecular outflows are reviewed. The physical outflow parameters are mainly obtained from CO observations, but we also pay attention to more rare and complex molecules such as CH3 OH, H2 CO, SiO, HCOC , etc. Recent results from chemical surveys of young protostellar outflows show that the peculiarities of the chemical behavior together with other structural parameters (such as the presence of molecular bullets, the flow opening angle, and the mechanical power efficiency) can be used to produce a classification of the observed outflows, and we suggest that this represents a rough time evolutionary sequence. We finally discuss the properties of the underlying primary wind which is assumed to drive the molecular outflows. Recent observations and theoretical work suggest that a primary wind with two components (a highly collimated jet and a wide-angle wind) is needed to explain the observations.

R. Bachiller () Observatorio Astronomico Nacional (IGN), Calle Alfonso XII, 3. 28014 Madrid, Spain e-mail: [email protected]

K. Tsinganos et al. (eds.), Protostellar Jets in Context, Astrophysics and Space Science Proceedings, DOI 10.1007/978-3-642-00576-3 46, c Springer-Verlag Berlin Heidelberg 2009 

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1 Molecular Outflows in Context Nearly thirty years after its discovery [50], the bipolar molecular outflows around young stars and protostars still remain as one of the most puzzling phenomena involved in star formation. In fact, although such outflows are one of the most spectacular manifestations of the birth of every new star (much more spectacular that the purely accretion processes), many basic questions about the origin and propagation of protostellar outflows still remain unanswered. The presence of winds emerging from young stars was first observed in the optical, and the properties of the spectacular optical jets are reviewed in the articles by J. Bally and by J. Eisl¨offel in this volume. Evidence for stellar winds is also provided by the free-free emission at cm wavelengths [19, 42]. When the stellar winds interact with the surrounding molecular cloud, they give rise to high-velocity outflows of molecular gas. Bipolar molecular outflows are then mainly made of swept-up material. Such outflows can be observed in lines of many molecules, most notably in the H2 infrared lines (e.g. [11, 25]) and in the rotational lines of carbon monoxide (CO) at millimetre wavelengths (e.g. [4, 3], see Fig. 1 as an example). Systematic searches for molecular outflows around young stellar objects have revealed that this is a extraordinarily common phenomenon: several hundreds

Fig. 1 L 1157 is a text-book example of bipolar outflow molecular outflow. This composition shows the Spitzer 8 m emission (gray-scale) overlaid with the CO 21 emision (contours). The upper left inset is the same gray scale without the contours. The lower right inset is a closeup (˙75 arcsec) of the absorption structure (presumably a disk) with contours of (14)  0.042 MJy sr‹1 . Figure from Looney et al. [31]. CO 21 emission from Bachiller et al. [9]

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outflows have been identified today (e.g. [56]). So it is now believed that every star, no matter its mass, goes through an outflow phase during its early evolution, and this is a clear indication that outflows are an essential ingredient in the process of star formation.

2 Current Challenges in Molecular Outflow Research Although the assumption that those molecular outflows are mainly made of ambient material which has been swept-up and put into motion by a “primary wind”, the properties of such a primary agent are not well understood. The optical jets are clearly related to the molecular outflows since when both phenomena are observed both coincide in direction and sense. However, the momentum and power of the ionized jet component is usually not enough to accelerate the associated molecular outflows. The detection of highly-collimated jets of extremely-high-velocity (EHV) gas around Class 0 protostars (e.g. L 1448, [6]; IRAS 03282, [7]; HH 211, [24]; IRAS 04166 + 2706, [53]; HH 212, [18]) provides a link between molecular outflows with relatively wide opening angles and the highly-collimated optical jets (Fig. 2). The properties of such outflows are reviewed by B. Nisini in this volume. A wealth of observations, in particular of Class 0 objects, clearly indicate that

Fig. 2 L 1448 provides an example of jet-like molecular outflow. The left panel shows the EHV jet-like component as seen in CO 21 emission (solid and dashed contours represent blue-shifted and red-shifted emission, respectively). Note the presence of high-velocity molecular bullets along the jet axis. The spectra obtained toward the positions of bullets B2 and R2 are shown in the right panels. Figure from Bachiller et al. [6]

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jets are needed to explain the properties of molecular outflows in general [37, 4]. However, the suggestion made by Shu et al. [49] that wide-angle winds are also important in the generation of outflows can not be ruled out, as none of these models (pure jets or pure wide-angle winds) seems able alone to explain the morphologies of all kind of outflows ([29]; T.P. Downes, this volume). The study of the nature of the primary wind is making rapid progress recently thanks to (1) the recent increases in sensitivity and resolution of mm-wave interferometers (IRAM, SMA, CARMA) and (2) the production of more refined theoretical models of outflows that include both jet and wide-angle components (e.g. M.N. Machida, this volume). However, many pieces are still missing in the building of a theory that should unify all the mass-loss phenomena observed in young stars and protostars.

3 CO Observations of Molecular Outflows Molecular outflows provide a time integrated picture of the wind activity around a young stellar object. Such outflows are best traced by their emission in lines of Carbon Monoxide (CO). This is because (1) CO is a very abundant molecule, (2) CO has both simple energy level structure and low dipole moment (0.1 Debye), and (3) different isotopic species can be observed at mm wavelengths. A measurement of the CO excitation temperature provides a good estimate of the gas kinetic temperature, and by observing lines from different isotopic species (e.g 12 CO and 13 CO) it is possible to estimate the CO optical depth. From the temperature and the optical depth, it is then possible to estimate the CO column density, and, by assuming a CO/H2 ratio, the H2 column density, so the mass of the outflowing gas (see e.g. [32]). Observations at mm wavelengths usually provide very good velocity resolution, so from CO observations it is thus possible to study the distribution of mass with the line-of-sight velocity, and estimate in this way the momentum, kinetic energy, and mechanical power of the outflowing gas. Although CO spectroscopy provides indeed the best method to obtain the physical parameters of molecular outflows, these estimates are subject to important uncertainties. First, the orientation angle of the outflow to the line-of-sight is needed to convert the observed radial velocities in actual space velocities. Second, there is an important ambiguity separating the outflow emission from the ambient cloud component in the CO line profiles. Third, the CO/H2 ratio is not well known in regions that can be significantly affected by shocks. As first pointed out by Cabrit and Bertout [14], the usual accuracies in the estimate of outflows parameters are within a factor of 2 in the mass, a factor of 10 in the momentum, and a factor of 60 in the mechanical power. The wealth of CO observations accumulated during nearly 30 years has revealed that molecular outflows have terminal velocities in the range of a few km s1 to about a few 102 km s1 , sizes ranging from less than 0.1 pc to several pc,

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kinematical time-scales from about 103 yr to about 105 yr, and masses from a few 104 Mˇ to several hundreds Mˇ . Interestingly enough, protostellar outflows have kinetic energies comparable to the gravitational binding energy of the dense core where the stars have been formed, and this seems to be true both for low-mass and high-mass young stellar objects. So the outflow driven by a new star has the potential to disperse the entire core where the star has been formed (e.g. [51, 48]). In the case of intermediate (Herbig Ae/Be) stars, outflows are able to sweep out about 90% of the core by the end of the premain-sequence phase [21]. It is worth underlining the recent efforts to study outflows around very low mass objects (eg. L 1014, [13]; IC 348-MMS 2, [52]). Although these outflows only contain about a few 104 Mˇ , the bipolarity and general characteristics are very similar to those around more massive objects. Moreover, recent work made by E. Whelan and collaborators (see their article in this volume, see also [41]) provides evidence for stellar winds around brown dwarfs, this is a work that clearly deserves continuation by searching for CO outflows around such very low mass objects.

4 Rare and Complex Molecules in Molecular Outflows The propagation of the supersonic outflows through the surrounding medium happens via shock waves. The rapid heating and compression of the gas trigger different microscopic processes (such as molecular dissociation, endothermic chemical reactions, sublimation of grain mantle ices, and disruption of the grains) which do not operate usually. Outflows can then contribute to the chemical enrichment of the vicinity of young stars, and the corresponding chemical composition of the medium (taken together with other phenomena such as the gradual clearing of the outflow path) can be considered as an indication of the age of the outflow (e.g. [9]). The chemical anomalies triggered by outflows are well illustrated in the molecular line surveys of the so-called chemically active outflows (L 1157: [8]; BHR71 [22]). More recent mapping observations have shown the presence of significant spatial chemical segregation in a number of sources, for example: L 1157 [9], NGC 1333-IRAS 4 (Choi et al. 2004), and Cep A [16]. With enhancements up to six orders of magnitude, SiO is probably the molecule that best distinguish the regions subject to strong shocks (e.g. [4], and references therein). The SiO line profiles usually present wings as broad as those of the CO lines. Sputtering of atomic Si from dust grains [45, 15] is believed to be the main mechanism for producing the observed high abundances of SiO. Nevertheless, recent observations by Santiago-Garc´ıa et al. [43], when compared to the chemical model of Glassgold et al. [23], also suggest that the usual gas-phase chemistry in a dense protostellar wind could produce high abundances of SiO (see Sect. 6). Narrow lines of SiO (line widths < 1 km s1 ) have also been detected around a number of low-mass young stellar objects [30, 17, 26]. Such narrow lines have been interpreted as the signature of a shock precursor component [26, 27]. An alternative

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explanation provided by Codella et al. [17] is that the SiO is indeed produced in the high-velocity gas which is subsequently slowed down in timescales of the order of 104 yr. CH3 OH and H2 CO have been observed to be enhanced in several outflows by factors of the order of 100 (e.g. [9, 28, 34]). These two species are believed to be evaporated directly from the icy dust mantles. Other simple species like HCOC , SO, SO2 , etc, which have been also observed to be enhanced, could be produced in the gas phase. The chemistry of Sulfur-bearing molecules is particularly interesting, as some of these molecules could provide quantitative indicators of the chemical evolution that could be used as rough clocks to date outflows (e.g. [54, 55, 16]). Very recently, Arce et al. [2] have reported the detection of complex organic molecules (HCOOCH3 , CH3 CN, HCOOH, and C2 H5 OH) toward a hot spot in the young protostellar outflow L 1157. Since the time scales associated with the outflow gas are too short (<2103 yr) to allow the formation in the gas phase, it is concluded that these species are directly ejected from grain mantles. The relative abundances of these molecules with respect to methanol are found to be similar to those measured in hot cores and in the molecular clouds near the Galactic Center, which suggests that the composition of the dust grain mantles is very similar in many different regions of the Galaxy.

5 The Evolution of Young Protostellar Outflows Significant observational effort has been devoted during recent years to study the abundance of some key molecules other than CO in bipolar outflows. SantiagoGarc´ıa et al. [43] have carried out a systematic study in lines of SiO, CH3 OH, and H2 CO to characterize a sample of 16 outflows around (Class 0 to Class I) protostars. These three molecules are traditionally assumed to be good tracers of shocks: SiO is considered to be a result of the partial disruption of dust grain cores, and CH3 OH and H2 CO are assumed to be directly evaporated from dust grain mantles. This survey can be used to classify young outflows according to the emission characteristics of these molecules, and to empirically define a rough time sequence (see also [10]). As we progress in this sequence, the opening angle of the outflow increases, the terminal velocity decreases, the outflow efficiency (ratio of the outflow mechanical power to the luminosity of the central object) decreases, and the bolometric temperature of the central object (a measure of its age, according to Myers et al. [38]) increases. The four groups identified in this sequence are the following: Group 1.- The first stage in the evolution sequence is represented by molecular outflows with highly-collimated EHV components (such as L 1448, IRAS 03282, and IRAS 04166). The abundance of SiO in the EHV jet presents enhancements by several orders of magnitude with respect to the ambient medium. Group 2.- The second stage corresponds to chemically active outflows (such as L 1157, BHR 71, and IRAS 16293). These outflows are less collimated and have lower terminal velocities than in the phase before. They have hot spots exhibiting

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strong chemical anomalies. The spectral lines of SiO, CH3 OH, and H2 CO present profiles with strong wings (e.g. [8, 9]). Group 3.- The third stage, which can be considered as the transition from Class 0 to Class I, is represented by outflows like L 483 and L 1527. No sign of wings is seen in the profiles of the SiO and CH3 OH lines, and only weak wings are perceived in the profiles of the H2 CO lines. Group 4.- The last stage corresponds to outflows powered by Class I objects. These outflows (represented by L 1551 IRS 5) display clear shell structures and are usually associated to optical HH objects. No chemical anomalies are observed. The dense cores seem to be strongly disrupted by the action of the outflows. From a comprehensive survey of outflows around Class 0, I, and II objects, Arce & Sargent [1] confirmed the increase of opening angle as the outflow evolves, they measured that the opening angle at the base of the outflow increases from about 10ı to 100ı , as the central object evolves from Class 0 to Class II. Futher evidence that cavities expand with time, beginning as a thin jet-like channel that widens to reveal the central protostar/disk system, has been recently provided by Seale & Looney [46].

6 The Chemistry of Molecular EHV Jets An interesting result of the work by Santiago-Garc´ıa, et al. [43] concerns the comparison of SiO, CH3 OH, and H2 CO abundances in the low-velocity outflows and in the EHV jets. The CH3 OH/SiO ratio presents a nearly constant value of 30 in the low velocity outflows, independent of the source and of the SiO intensity (and the behavior of H2 CO is very similar); such a ratio is well explained with models of shock chemistry, confirming that the low-velocity outflow emission arises from shocked ambient gas. On the other hand, the EHV gas has no detectable CH3 OH emission (no matter how bright the SiO lines are), so the corresponding CH3 OH/SiO upper limits in the EHV gas are one or two order of magnitude lower than the ratio in the wing component (see Fig. 3). This indicates that the chemistry of the EHV gas is radically different to that of the shocked ambient material. An interesting possibility is that the EHV emission is directly tracing a neutral protostellar wind. The chemistry of this kind of wind was modelled by Glassgold et al. [23] who showed that CO and SiO are the two most abundant molecular species after H2 . This is in general agreement with the observations of the EHV bullets which, in spite of deep searches in several molecules, have only been detected in CO and SiO. The lack of molecular species in the EHV gas is also reminiscent of the Oxygen-rich envelopes around evolved stars in which CH3 OH and H2 CO have extremely low abundances (e.g. [35]). In conclusion, the EHV jet-like components observed in many Class 0 outflows are clearly very closely linked to the primary wind emerging from the central protostar. Moreover, if the interpretation of the chemical differences of the EHV gas and

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Fig. 3 The H2 CO/SiO and CH3 OH/SiO abundance ratios in the low-velocity wing component of molecular outflows (black squares) present nearly contant values independently of the SiO abundance. However, these ratios exhibit much lower values (by one or two orders of magnitude) in the jet-like EHV component (open symbols, all upper limits). From Santiago-Garc´ıa et al. [43]

the low velocity outflow is correct, the EHV component could actually and directly represent the primary wind emerging from the root star/disk system.

7 A Case of Study: The Outflow from IRAS 04166+2706 IRAS 04166+2706 (IRAS 04166 hereafter) is a particularly interesting object since it is one of the most nearby Class 0 protostars (in Taurus, at only 140 pc of distance) whose molecular outflow contains a highly-collimated extremely-high velocity (EHV) component. The outflow was first detected by Bontemps et al. [12], but the jet-like nature of the flow was revealed by the higher sensitivity observations of Tafalla et al. [53]. IRAS 04166, the driving object, is a protostar of 0.4 Lˇ of luminosity surrounded by an envelope (most likely a circumstellar disk) of 0.014 Mˇ . From observations of this outflow made with the IRAM interferometer, which provide very high linear resolution (down to 210 AU), Santiago-Garc´ıa et al. [44] are able to separate the distribution of the EHV gas from that of the lower velocity component providing important clues on the nature of the primary wind which is driving the outflow. The observations show that the outflow, which is highly symmetric, consists of two components: (1) at low velocities (<10 km s1 , unprojected) two opposed limb-brightened conical shells with and opening angle of 32ı , (2) at high velocities (>30 km s1 , unprojected) two opposed highly-collimated jets placed at the shell axes. The highly collimated jet is very “clean” and shows no evidence for precession. This bipolar jet is made of a symmetric collection of at least 7 pairs of high-velocity

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Fig. 4 Left:Map of the CO 21 integrated emission in the EHV velocity range (50> jV  V 0j > 30 km s1 ) for the blue-shifted (north) and red-shifted (south) lobes of the IRAS 04166 outflow. First level and intervals are 1.5 by 1.0 Jy beamk1 km s1 (blue) and 1.25 by 1.0 Jy beamk1 km s1 (red). The star marks the position of the protostar traced by the 1.3 and 3.5 mm continuum. Right: Position-velocity diagram of the CO 21 emission for the blue- and redshifted lobes along the IRAS 04166 outflow. First level and intervals are 150 by 75 mJy beam1 , respectively. The dotted lines marks the velocity range of the EHV jet emission, and the dashed vertical line marks the position of the protostar. From Santiago-Garc´ıa et al. [44]

clumps (“molecular bullets”) which broaden with distance to IRAS 04166 (see Fig. 4). These molecular bullets are similar to those observed in other Class 0 objects ([4], and references therein). The morphology of the outermost peaks is reminiscent of bow shocks, similar to those observed in L 1448 [20]. The full opening angle of the IRAS 04166 jet is less than 10ı . Since there is no sign of precession (like for instance in L 1157, [9]), and the width of the bow shocks (<10ı ) is insufficient to explain the larger opening angle (32ı ) of the low velocity shells, we are led to conclude that the shells are created by a wind component of relatively wide opening angle. In summary, in addition to the prominent EHV jet, a wide-angle wind component of the primary wind is needed to explain the observations. A primary wind consisting of both a highly-collimated EHV jet and a wide-angle wind is well in the line of the unified models of Shang et al. [47] and Machida et al. [33] (see also the paper by M.N. Machida in this volume). The IRAS 04166 outflow presents a conspicuous saw-tooth pattern in the position-velocity (PV) diagrams of the jet (see Fig. 4). Such a peculiar pattern results from the systematic internal velocity gradients within the bullets: the gas in the bullet heads moves systematically slower than the gas in the tails. The most natural explanation of such saw-tooth pattern is that each bullet represents a dense gas compression travelling inside the nearly constant high velocity jet. This behavior

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has been modelled in the context of Herbig-Haro “pulsed” jets ([40, 36]; A. Raga, this volume). In these models, the ejection velocity suffers time variations, so that the fast-moving material runs into slower previously-ejected material at different positions along the jet. The internal working surfaces which are produced at these positions, where the gas is compressed and heated, are expected to emit bright optical lines [39]. Time variations in the ejection velocity is thus the most plausible explanation for the saw-tooth pattern in the PV diagram of IRAS 04166. The physical cause of such variations is not well known, but since the ejection is assumed to be directly related to accretion, it seems likely that variations in the accretion rate are translated into variations in both the mass and the ejection velocity.

8 Conclusions and Future Prospects As we have emphasized in this paper, the recent high-angular resolution observations of outflows are often done in many molecular lines (not just in CO lines), and this is allowing the characterization of the outflows with increasing detail. Molecular spectroscopy is not only able to provide the general physical parameters of the outflows, but it is also able to provide the chemical composition of the jets and of the surrounding molecular gas. The physical and chemical impact of outflows can then be used as an indication of the evolutionary stage, so empirical models of the outflow evolution can be established. The nature of the primary wind which is the actual driving agent of outflows is one of the most enigmatic problems which needs to be solved. As we have pointed out in this paper, the recent observations by Santiago-Garc´ıa et al. [44] suggest that the protostellar primary wind consists of two components: a jet-like component made of neutral (atomic and molecular) gas surrounded by a wider angle component. This result is well in line with recent theoretical work ([47, 33]; Machida, this volume). There is also a wealth of observations showing the episodic nature of the high-velocity jets (e.g. [4]). The variations in the velocity of the ejecta (probably produced by time variations in the accretion) can produce internal working surfaces along the jet which take the form of symmetrical “molecular bullets” in the observations. This primary jet could be mainly done of molecular gas, according to the comparison of chemical observations to the only available model for the chemistry in such a kind of wind [23]. On the other hand, a wide-angle wind is needed to explain the wide shells observed in a molecular outflow like this around IRAS 04166, where there is no sign of precession and the terminal bow shocks of the jet are significantly narrower than the shells. Indeed it seems difficult today to directly probe (observationally) this wide-angle component of the primary wind. Further progress in outflow research is expected to result from the improvements of current telescopes and from the construction of new large facilities. In particular, the IRAM, CARMA, and SMA interferometers, which are already providing excellent data, are expected to provide increasingly high angular resolution and sensitivity which should allow the mapping of many outflows with even finer detail.

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The superb capabilities of the Atacama Large Millimeter Array, ALMA (expected to be fully operational by 2013) will clearly shed light on many phenomena related to outflows and with star formation in general (Sepherd 2008; J. Richer, this volume). The Herschel Space Observatory (HSO, to be launched in 2009) will allow the measurement of some key shock tracers, in particular water, which can not be observed from the ground. The observation of the neutral atomic component in outflows is also of great interest. This wealth of very high quality observations should be accompanied by model simulations of increasing detail and versatility. Such numerical work should provide model maps of the observable parameters (densities, temperatures, and also velocities, and line intensities) that are directly comparable to the observed maps. Acknowledgements The author it indebted to Dr. Mario Tafalla and Joaquin Santiago-Garc´ıa for stimulating and very fruitful research collaboration on outflows over several years; such collaborative work has provided the main basis for this review paper. Interesting discussions with Drs. A. Fuente, C. Codella, and with many members of the JETSET network are also gratefully acknowledged.

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21. Fuente, A., Mart´ın-Pintado, J., Bachiller, R., Rodr´ıguez-Franco, A., Palla, F.: Astron. Astrophys. 387, 977 (2002) 22. Garay, G., K¨ohnenkamp, I., Bourke, T.L., Rodr´ıguez, L.F., Lehtinen, K.K.: Astrophys. J. 509, 768 (1998) 23. Glasssgold, A.E., Mamon, G.A., Huggins, P.J.: Astrophys. J. 373, 254 (1991) 24. Gueth, F., Guilloteau, S.: Astron. Astrophys. 343, 571 (1999) 25. Hodapp, K.W., Ladd, E.F.: Astrophys. J. 453, 715 (1995) 26. Jim´enez-Serra, I., Mart´ın-Pintado, J., Rodr´ıguez-Franco, A., Marcelino, N.: Astrophys. J. Lett. 603, L49 (2004) 27. Jim´enez-Serra, I., Mart´ın-Pintado, J., Rodr´ıguez-Franco, A., Mar´ın, S.: Astrophys. J. Lett. 627, L121 (2005) 28. Jørgensen, J.K., Hogerheijde, M.R., Blake, G.A., van Dishoeck, E.F., Mundy, L.G., Sch¨oier, F.L.: Astron. Astrophys. 415, 1021 (2004) 29. Lee, C.-F., Mundy, L.G., Stone J.M., Ostriker, E.C.: Astrophys. J. 576, 294 (2002) 30. Lefloch, B., Castets, A., Cernicharo, J., Loinard, L.: Astrophys. J. Lett. 504, L109 (1998) 31. Looney, L.W., Tobin, J.J., Kwon, W.: Astrophys. J. Lett. 670, L131 (2007) 32. Levreault, R.M.: Astrophys. J. Suppl. Series 67, 283 (1988) 33. Machida, M.N., Inutsuka, S.I., Matsumoto, T.: Astrophys. J. 676,1088 (2008) 34. Maret, S., Ceccarelli, C., Tielens, A.G.G.M., Caux, E., Lefloch, B., Faure, A., Castets, A., Flower, D.R.: Astron. Astrophys. 442, 527 (2005) 35. Marvel, K.B.: Astron. J. 130, 261(2005) 36. Masciadri, E., Vel´azquez, P.F., Raga, A.C., Cant´o, J., Noriega-Crespo, A.: Astrophys. J. 573, 260 (2002) 37. Masson, C.R., Chernin, L.M.: Astrophys. J. 414, 230 (1993) 38. Myers, P.C., Heyer, M., Snell, Ronald L., Goldsmith, P.F.: Astrophys. J. 324, 907 (1998) 39. Raga, A.C., Binette, L., Cant´o, J., Calvet, N.: Astrophys. J. 364, 601 (1990) 40. Raga, A.C., & Noriega-Crespo, A.: Astron. J. 116, 2943 (1998) 41. Ray, T., Dougados, C., Bacciotti, F., Eisl¨offel, J., and Chrysostomou, A.: Protostars and Planets V, 231 (2007) 42. Rodr´ıguez, L.F.: Rev. Mex. Astron. Astrofis. (Ser. de Conf.) 1, 1 (1995) 43. Santiago-Garc´ıa, J., Tafalla, M., Bachiller, R.: in preparation (2008a) 44. Santiago-Garc´ıa, J., Tafalla, M., Johnstone, D., Bachiller, R.: Astron. Astrophys., in print (2008b) 45. Schilke, P., Walmsley, C.M., Pineau des Forˆets, G., Flower, D.R.: Astron. Astrophys. 321, 293 (1997) 46. Seale, J.P., Looney, L.W.: Astrophys. J. 675, 427 (2008) 47. Shang, H., Allen, A., Li, Z.Y., Liu, C.F., Chou, M.Y., Anderson, J.: Astrophys. J. 649, 845 (2006) 48. Shepherd, D.S., Borders, T., Claussen, M., Shirley, Y., Kurtz, S.: Astrophys. J. 614, 211 (2004) 49. Shu, F.H., Ruden, S.P., Lada, C.J., Lizano, S.: Astrophys. J. Lett. 370, L31 (1991) 50. Snell, R.L., Loren, R.B., Plambeck, R.L.: Astrophys. J. Lett. 239, L17 (1980) 51. Tafalla, M., Bachiller, R., Wright, M.C.H., Welch, W.J.: Astrophys. J. 474, 329 (1997) 52. Tafalla, M., Kumar, M.S.N., Bachiller, R., Astron. Astrophys. 456, 179 (2006) 53. Tafalla, M., Santiago, J., Johnstone, D., Bachiller, R.: Astron. Astrophys. 423, L21 (2004) 54. Wakelam, V., Caselli, P., Ceccarelli, C., Herbst, E., and Castets, A.: Astron. Astrophys. 422, 159 (2004) 55. Wakelam, V., Ceccarelli, C., Castets, A., Lefloch, B., Loinard, L., Faure, A., Schneider, N., and Benayoun, J.J.: Astron. Astrophys. 437, 149 (2005) 56. Wu,Y., Huang, M., He, J.: Astron. Astrophys. Supp. 115, 283 (1996)

Molecular Outflows: Observations

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Driving Mechanisms for Molecular Outflows Turlough P. Downes

1 Introduction Molecular outflows are observed to be closely associated with star formation. The cumulative momentum and the momentum injection rate in these outflows are important parameters in theories of star formation. The cumulative momentum in an outflow is a measure of the feed-back from star formation on molecular cloud turbulence. The level of turbulence in a cloud also effects the formation of further stars and, indeed, the survival of the cloud itself (e.g. [14]). In addition the rate of injection of momentum is an important constraint for theoretical models of outflows from young stars [9, 17]. Hence, while these outflows are interesting in themselves, it is also critical to understand their origin and behaviour as part of the general study of how stars themselves form.

T.P. Downes () Dublin City University & Dublin Institute for Advanced Studies e-mail: [email protected]

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In Sect. 2 we first detail the fundamental properties of molecular outflows in general. We restrict our discussion to those properties which are common to all outflows, or at least a substantial fraction of them. These are the properties which must be reproduced by our models of outflow generation before we can claim to understand molecular outflows. Section 3 details models which have been presented in the literature to date while Sect. 4 focuses on an in-depth comparison of the jetand wind-driven models of molecular outflows. Finally, Sect. 5 contains some conclusions on our discussion and suggestions for further study in this area.

2 Fundamental Properties of Molecular Outflows In order to model molecular outflows we first need to know some of the fundamental properties of these flows. The reader is referred to the excellent review [17] for an in-depth review of the observational properties of molecular outflows. Molecular outflows are massive phenomena with masses ranging from 0.1 – 100 Mˇ [9]. This fact alone imposes a strong constraint on models of molecular outflows: they cannot be launched directly from material processed through an accretion disk (e.g. [10]). They must consist of ambient molecular material accelerated by outflows of some kind from the nascent star. The lifetime of molecular outflows is difficult to determine. The most common method to date has involved defining some characteristic velocity (e.g. the intensityweighted mean) and some characteristic length associated with the outflow (e.g. the maximum length [22, 20]). The age is then given by T D

L ; v

(1)

where T is the estimated age (the “dynamical age”) of the system and L and v are the characteristic length and velocities of the flow. Generally speaking the dynamical age of molecular outflows is thought to be in the range 104 – 105 yrs (e.g. [17, 10]). There is considerable inaccuracy associated with these age estimates and different ways of calculating L and v can have a significant impact on the calculation with possible under- or over-estimates of up to an order of magnitude [6, 15]. Molecular outflows are, by definition, rather poorly collimated with length-towidth ratios ranging from five up to as high as 30 or even more (e.g. [17]). It appears that younger outflows are better collimated than their older counterparts, or perhaps that older outflows (of age 105 yrs) have a wide-angled component which the younger outflows do not possess [17, 19]. Images of two typical outflows are shown in Fig. 1. Two different morphologies are observed for PV diagrams of molecular outflows: “Hubble-like” diagrams, and concave diagrams. An example of a Hubble-like PV diagram is that of HH240/241 in Fig. 2. It is clear that the maximum velocity appears to scale roughly linearly with the distance from the driving source. This feature is

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Fig. 1 CO emission (contours) overlayed on H2 emission (grey scale) for two typical outflows: HH240/241 (left) and VLA 05487 (right). Adapted from [12]

Fig. 2 PV diagrams for HH240/241 (left) and VLA 05487 (right). In each the H2 image of the jet is shown to the left of the plot. Adapted from [12]

quite common in observations of molecular outflows ([10, 11, 24, 8, 7, 13, 5, 1]). An example of a flow possessing a concave PV diagram is VLA 05487 (Fig. 2, right panel). In this type of diagram the maximum velocity does not occur at the maximum distance from the source and, indeed, there is a clear concave appearance to the PV diagram. Another important property of molecular outflows is the line profiles we observe from them. One of the most robust properties of all observational results on molecular outflows is that, in the wings of emission lines (particularly the J D 1  0 and J D 2  1 lines of CO), the intensity is a power-law in velocity: I.v/ / v ;

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where  is in the range 1.5 – 2 at low velocities to as large as eight at higher velocities [10, 24, 12, 7, 5, 6]. This result was extended to the S(1) 1 – 0 line of H2 by [18]. Many authors (e.g. [10, 27]) have made the assumption that the intensity is directly proportional to the emitting mass and so this power-law relation has become known as the “mass-velocity relation” in the literature. However, [5] have shown that this assumption is not valid except at low velocities.

3 Models of Outflows Any model of molecular outflows must explain the properties outlined in Sect. 2. We now detail the three most widely discussed models.

3.1 Steady-State Entrainment The first of our models relies on direct entrainment of ambient molecular cloud material by a stellar jet as it propagates [4, 25, 26]. In principle all models must have some sort of entrainment occurring as it is highly unlikely that the amount of material typically observed in molecular outflows could be processed through a central launch engine. A schematic diagram of this model can be seen in Fig. 3. As the jet propagates it is postulated that fluid instabilities such as the Kelvin-Helmholtz instability give rise to entrainment and acceleration of ambient material along the walls of the jet. A nice feature of this model is that it will naturally give rise to the “Hubble law” observed in PV diagrams (Sect. 2): close to the source the entrained material will not be moving very fast, while further away it will have accelerated more. However, in order for this model to work we must find an instability which will disrupt the jet boundaries in such a way that momentum can be transferred to the material surrounding the jet. The most obvious of these is the Kelvin-Helmholtz instability. There are a number of difficulties with this: 1. Stellar jets have a very high Mach number. This suppresses the growth of the KH instability. 2. The material surrounding the jet is, in fact, “cocoon” material: i.e. material which has largely been processed through strong shocks at the end of the jet and which Fig. 3 Schematic diagram of the steady-state entrainment model. Ambient material is entrained (light grey) along the walls of the stellar jet (dark grey) as it propagates through the molecular cloud. See text

Bowshock

Molecular Outflow Jet

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has subsequently expanded to form the tenuous medium between the jet and the bowshock. There is thus a large density contrast between the jet and the medium surrounding it – this fact also suppresses the growth of the KH instability. 3. Finally, and perhaps most serious, any suitable instability will have to be strong enough to transfer around 50% of the jet’s momentum to the ambient material without decollimating the jet itself. It is hard to envisage how this could happen, particularly given the large extent of stellar jets which have been observed in more recent times. There is one further problem with this model: as noted above, the material surrounding the jet has been processed through strong shocks and is low density. Both of these factors will lead to the medium being largely atomic: the strong shocks will dissociate a large fraction of any molecular material and the low density will lead to very long reformation times. This means that even if material were entrained via some instability that material would be largely atomic and would not, as a result, form a molecular outflow.

3.2 Jet-Driven Outflows: “Prompt Entrainment” Another model of molecular outflows is that of so-called “prompt entrainment”. In this model, as the stellar jet propagates into the molecular cloud it drives a bowshock into it. This bowshock sweeps up ambient molecular material and it is this material which is identified as the molecular outflow (e.g. [4, 16, 27, 24, 7, 5, 6] and many others). Figure 4 contains a schematic diagram of this model. Clearly, the bowshock will be significantly less well collimated than the jet itself and may indeed resemble a molecular outflow. It has been shown that a momentumconserving bowshock will give rise to the “Hubble law” seen in the PV diagrams of many molecular outflows [11, 7]. Figure 5 contains a PV diagram for the J D 2  1 CO line of a simulated jet-driven bowshock integrated over the whole flow. It is clear that there is an approximately linear increase of velocity with distance from the source (cf Fig. 2). In addition, both analytic calculations and numerical simulations have shown that the intensity in CO line wings is related to the velocity by a power-law with the appropriate exponent [11, 24, 7]. Particularly good fits between the model and observations have been demonstrated in [5]. This work showed that the break in Fig. 4 Schematic diagram of the prompt entrainment model. Ambient material is swept up by the bowshock of the stellar jet as it propagates through the molecular cloud. See text

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Velocity (km/s)

100

0

–100

–200 0

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Fig. 5 PV diagram for the J D 2  1 CO line of a simulated jet propagating at 30ı to the plane of the sky. Note the approximately linear increase in velocity with distance from the source of the jet. The units of distance are 1014 cm and those of velocity are km s1

the power-law to a steeper relation at higher velocities seen in observations is explained by the T1 dependence of emissivity on temperature once the temperature becomes significantly higher than the excitation temperature of the upper level of the emitting line. In addition, [5] extended this result from numerical simulations to the intensity-velocity relation for the H2 S(1) 1–0 line. In several respects, then, the prompt entrainment model satisfies the observations. There are two main challenges from observations for the prompt entrainment model. The length-to-width ratios of jet-driven bowshocks tend to be very large: stellar jets have very high Mach numbers (of order 102 with respect to the ambient sound speed) and hence they will be rather long and narrow. However, many molecular outflows are very poorly collimated. Secondly, [10] noted that there tends to be very little blue-shifted emission in the red lobe of a molecular outflow and vice versa for the blue lobe. This suggests that the momentum in a molecular outflow must be predominantly directed along its axis of symmetry. The bowshock, on the other hand, expands perpendicular to this axis and so will introduce some level of contamination. The first difficulty can be alleviated if the driving jet is precessing with a nonnegligible precession angle [23]. The second is more challenging and is still an open question.

3.3 Wind-Driven Outflows A different model of molecular outflows relies on wide-angle winds from the driving source to power the outflow. A suitable theoretical candidate for such a wind is the X-wind model [21]. The scenario for the wide-angled wind model is as follows.

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Wide angle wind Molecular Outflow

i = 30

0

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Fig. 6 Schematic diagram of the wind-driven model. Ambient material is swept up by the shock driven into the molecular cloud by the wind. See text

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Fig. 7 Simulated PV diagram for a wide-angled wind model of molecular outflows assuming the flow to be at an inclination angle of 30ı to the plane of the sky. Note the concave nature of the diagram (from [12])

An X-wind process drives a density-stratified wind into the ambient medium. This ambient medium has an r 2 profile since it is part of the envelope of the nascent star. The coincidence of this density profile of the ambient medium and the stratified nature of the X-wind results in a concave PV diagram for the resulting thin shell of swept-up ambient material [21]. Figure 6 contains a schematic diagram of this model. Some numerical simulations appear to confirm this result [13] (see Fig. 7). Such a scenario will naturally give rise to poorly collimated outflows of the type sometimes observed. The simulations of [13] assumed that the system is isothermal following from the theoretical assumptions of [21]. However, this assumption is unlikely to be valid for a wide range of outflows. Subsequent, higher resolution simulations of an isothermal flow do broadly support these early simulations, but once the assumption of isothermality is dropped the behaviour of the system changes dramatically [3]. Unfortunately no PV diagrams were simulated in the latter work. Figure 8 contains a simulated PV diagram for a wide-angled wind propagating into an ambient medium the density of which behaves as r 2 . Note the lack of concave structure for this PV diagram. Lee et al. [13] were able to generate broken power-laws for the mass-velocity relation from their simulations of wide-angled winds. However, these broken powerlaws were only evident for small inclination angles. In addition, the behaviour of the intensity-velocity relation was not possible to derive from the simulations as a result of the assumption of isothermality. Hence it is not yet known whether the wide-angled wind model reproduces the necessary line profiles.

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Fig. 8 Simulated PV diagram for the J D 2  1 CO line of a non-isothermal wide-angled wind with an inclination angle of 30ı to the plane of the sky propagating into an r 2 ambient medium (see text). Note the lack of a concave structure in this diagram

4 Jet-Driven versus Wind-Driven Outflows In this section we draw together the results discussed in Sect. 3 so that an overview of the relative strengths and weaknesses of the jet-driven and wind-driven models of molecular outflows is gained. The jet-driven model has the following strengths:  It reproduces the morphology of many molecular outflows (and possibly all out-

flows if precession of the driving jet is invoked).  It reproduces the “Hubble law” observed in many PV diagrams of molecular

outflows.  It is very successful at reproducing the intensity-velocity relations observed from

low J lines of CO. The wind-driven model is strong in the following respects:  It reproduces the morphology of many molecular outflows.  It appears to be able to reproduce the PV diagrams of some molecular outflows.  It may be able to reproduce the observed line profiles.

We now briefly list the weaknesses of each model. For the jet-driven model the current weaknesses are as follows:  It has difficulty reproducing the observational result that red-shifted lobes are

relatively uncontaminated by blue-shifted emission and vice versa.  Without invoking precession of the driving jet it is unclear whether a jet-driven

outflow can be as poorly collimated as some of those observed. Note, however, that variations in the density of the ambient medium can alleviate this situation [2]. The wind-driven model’s current weaknesses are listed below:  It is unclear whether it reproduces the observed line profiles.  It has not yet been shown to reproduce the observed PV diagrams without the

assumption that mass is proportional to intensity – an assumption which has been shown to be invalid in the case of molecular outflows.  It appears to require an r 2 density profile for the ambient medium.  It may require an unphysical assumption of isothermality.

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5 Conclusions There are still two viable models for the driving sources of moleuclar outflows. The jet-driven model is more developed in terms of comparison with observations and exploration of its parameter space with numerical simulations. It is remarkably successful at explaining many of the characteristics of outflows. Its main challenges lie in the generation of very poorly collimated outflows and the forward-directed momentum issue. Neither of these are impossible to explain with this model, but more work needs to be done to determine whether it can be successful in explaining these particular characteristics without resorting to extreme or unlikely physical conditions. The wind-driven model has not yet been as rigorously tested against observations as the jet-driven model. It is therefore more difficult to determine its main challenges. What is accepted is that it is certainly unable to explain the generation of well-collimated outflows. It will also have difficulties with the forward-directed momentum issue, although perhaps these difficulties may not be as severe as for the jet-driven model. We have yet to discover whether the assumption of an isothermal flow is necessary for this model, and also whether observed line profiles and PV diagrams are successfully reproduced by it. In short, more work is needed on this model.

References 1. H.G. Arce, D. Shepherd, F. Gueth, C.-F. Lee, R. Bachiller, A. Rosen, H. Beuther: In Protostars and Planets V, ed. Reipurth, B., Jewitt, D. and Keil, K., University of Arizona Press, 245 (2007) 2. S. Corkery: “Propagation of molecular outflows into inhomogeneous media”, MSc thesis, Dublin City University (2008) 3. A. Cunningham, A. Frank, P. Varni´ere, A. Poludnenko, S. Mitran, L. Hartmann: ApSS 298, 317 (2005) 4. D.S. De Young: ApJ 307, 62 (1986) 5. T.P. Downes, S. Cabrit: A&A 403, 135 (2003) 6. T.P. Downes, S. Cabrit: A&A 471, 873 (2007) 7. T.P. Downes, T.P. Ray: A&A 345, 977 (1999) 8. F. Gueth, S. Guilloteau: A&A 343, 571 (1999) 9. C.J. Lada: ARA&A 23, 267 (1985) 10. C.J. Lada, M. Fich: ApJ 459, 638 (1996) 11. C.R. Masson, L.M. Chernin: ApJ 414, 230 (1993) 12. C.-F. Lee, L.G. Mundy, B. Reipurth, E.C. Ostriker, J.M. Stone: ApJ 542, 925 (2000) 13. C.-F. Lee, J.M. Stone, E.C. Ostriker, L.G. Mundy: ApJ 557, 429 (2001) 14. C. Norman, J. Silk: ApJ 238, 158 (1980) 15. N.D. Parker, R. Padman, P.F. Scott: MNRAS 252, 442 (1991) 16. A. Raga, S. Cabrit: A&A 278, 267 (1993) 17. J.S. Richer, D.S. Shepherd, S. Cabrit, R. Bachiller, E. Chuchwell: In: Protostars and Planets IV, ed. Mannings, V., Boss, A.P., Russell, S.S. University of Arizona Press, 867 (2000) 18. L. Salas, I. Cr´uz-Gonz´alez: ApJ 572, 227 (2002) 19. D. Shepherd: ApSS 313, 41 (2007)

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20. D.S. Shepherd, K.C. Yu, J. Bally, L. Testi: ApJ 535, 833 (2000) 21. F.H. Shu, J. Najita, E.C. Ostriker, H. Shang: ApJ 455, 155 (1995) 22. R.L. Snell, N.Z. Scoville, D.B. Sanders, N.R. Erickson: ApJ 284, 176 (1984) 23. M.D. Smith, A. Rosen: MNRAS 357, 579 (1995) 24. M.D. Smith, G. Suttner, H.W. Yorke: A&A 323, 223 (1997) 25. S.W. Stahler: In Astrophysical Jets ed. Burgarella, D., Livio, M., O’Dea, C. Cambridge University Press 183 (1993) 26. S.W. Stahler: ApJ 422, 616 (1994) 27. Q. Zhang, X. Zheng: ApJ 474, 719 (1997)

Protostellar Jet and Outflow in the Collapsing Cloud Core Masahiro N. Machida, Shu-ichiro Inutsuka, and Tomoaki Matsumoto

Abstract Using three-dimensional resistive MHD nested grid simulations, we investigate the driving mechanism of outflows and jets in star formation process. Starting with a Bonnor-Ebert isothermal cloud rotating in a uniform magnetic field, we calculated cloud evolution from the molecular cloud core (nc D 104 cm3 , r D 4:6  104 AU) to the stellar core (nc D 1022 cm3 , r  1Rˇ ), where nc and r denote the central density and radius of each object, respectively. In the collapsing cloud core, we found two distinct flows: Low-velocity outflows (5 km s1 ) with a wide opening angle, driven from the adiabatic core, and high-velocity jets (30 km s1 ) with good collimation, driven from the protostar. High-velocity jets are enclosed by low-velocity outflow. The difference in the degree of collimation between the two flows is caused by the strength of the magnetic field and configuration

M.N. Machida () and S. Inutsuka Department of Physics, Graduate School of Science, Kyoto University, Sakyo-ku, Kyoto 606-8502, Japan e-mail: [email protected]; [email protected] T. Matsumoto Faculty of Humanity and Environment, Hosei University, Fujimi, Chiyoda-ku, Tokyo 102-8160, Japan e-mail: [email protected]

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of the magnetic field lines. The magnetic field around an adiabatic core is strong and has an hourglass configuration; therefore, the low-velocity outflow from the adiabatic core are driven mainly by the magnetocentrifugal mechanism and guided by the hourglass-like field lines. In contrast, the magnetic field around the protostar is weak and has a straight configuration owing to Ohmic dissipation in the high-density gas region. Therefore, high-velocity jet from the protostar are driven mainly by the magnetic pressure gradient force and guided by straight field lines. Differing depth of the gravitational potential between the adiabatic core and the protostar cause the difference of the flow speed. Low-velocity outflows correspond to the observed molecular outflows, while high-velocity jets correspond to the observed optical jets. We suggest that the protostellar outflow and the jet are driven by different cores, rather than that the outflow being entrained by the jet.

1 Introduction The observations indicate that outflow is ubiquitous in the star formation process, and flows from the protostars have two or more distinct velocity components [1, 2]. Typically, a flow from the protostar is composed of a low-velocity component (LVC) with 10–50 km s1 and a high-velocity component (HVC) with 100 km s1 . There is a clear trend toward higher collimation at higher flow velocity [3]. Since flows from the protostar have various morphological and kinematical properties, they cannot be explained by a single-class model. The flows that originated from the protostar are typically classified into two types: molecular outflow observed mainly with CO molecules [3], and optical jet observed by optical emission [4]. Molecular outflows observed by CO line emission exhibit a wide opening angle [5] and slower velocity of 10–50 km s1 [3], while optical jets observed by optical emission exhibit good collimation and higher velocity of 100–500 km s1 [6]. Observations indicate that around each protostar, high-speed jets with a narrow opening angle are enclosed by a low-velocity outflow with a wide opening angle [7]. However, the driving mechanism of these flows are still unknown. In this study, we calculate cloud evolution from the molecular cloud core (nc D 104 cm3 , rc D 4:6  104 AU) to stellar core formation (nc ' 1022 cm3 , rc ' 1 Rˇ ) using three-dimensional resistive MHD nested grid method, study the formation process of jets and outflows, and show the driving mechanisms of these flows.

2 Model Our initial settings are almost the same as those of [8, 9, 10, 11, 12]. We solve the resistive MHD equations including self-gravity (see, 1-5 of [10]). We adopt a spherical cloud with critical Bonnor-Ebert density profile having BE D 3:8411020 g cm3 (nBE D 104 cm3 ) of the central (number) density as the initial condition. The critical radius for a Bonnor–Ebert sphere Rc D 6:45 cs Œ4G BE;0 1=2 corresponds to

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Rc D 4:58  104 AU for our settings. Initially, the cloud rotates rigidly (˝0 D 7  1015 s1 ) around the z-axis and has a uniform magnetic field (B0 D 17 G) parallel to the z-axis (or rotation axis). To promote contraction, we increase the density by 70% from the critical Bonnor-Ebert sphere. The initial central density is therefore

0 D 6:53  1020 g cm3 (n0 D 1:7  104 cm3 ). We adopt the nested grid method [8, 13, 14] to obtain high spatial resolution near the center. Each level of a rectangular grid has the same number of cells (64  64  32), although the cell width h.l/ depends on the grid level l. The highest level of a grid changes dynamically. The box size of the initial finest grid l D 1 is chosen to be 24 Rc , where Rc denotes the radius of the critical Bonnor–Ebert sphere. A new finer grid is generated whenever the minimum local Jeans length J becomes smaller than 8 h.lmax /. The maximum level of grids is restricted to lmax D 30.

3 Results The molecular gas obeys the isothermal equation of state with temperature of 10 K until nc ' 5  1010 cm3 (isothermal phase), then cloud collapses almost adiabatically (51010 cm3 < nc < 1016 cm3 ; adiabatic phase) and quasi-static core (i.e., first core) forms during the adiabatic phase [15,16]. In our calculations, the first core forms when the central density reaches nc ' 81012 cm3 . The magnetic flux is removed from the first core during the adiabatic phase by the Ohmic dissipation [17]. After central density reaches nc ' 1016 cm3 , the equation of state becomes soft reflecting the dissociation of hydrogen molecules at T ' 2  103 K, and collapses rapidly. By this epoch, the central temperature becomes so high that the thermal ionization of Alkali metals reduces the resistivity and so that Ohmic dissipation becomes ineffective. Thus, the magnetic field becomes strong again as central region collapses. The second core (or protostar) [15] forms at nc ' 1021 cm3 . The magnetic field strength increases rapidly after the second core formation epoch (n > 1021 cm3 ), because the shearing motion between the second core and ambient medium amplifies the toroidal magnetic field around the second core. Figure 1 shows the structure of the low- and high-velocity flows, and the configuration of the magnetic field lines. It also shows the shapes of the first core (left panel; the projected density contours on the wall) and the protostar (right panel; the red isosurface). The purple and blue surfaces in Fig. 1 indicate the iso-velocity surface of vz D 5 km s1 and vz D 0:5 km s1 , respectively. The flow inside the purple iso-velocity surface has a velocity of v > 0:5 km s1 (low-velocity component; LVC), while the flow inside the blue iso-velocity surface has a velocity of v > 5 km s1 (high-velocity component; HVC). The HVC is enclosed by the LVC. The LVC flow is mainly driven from the first core, while the HVC flow is mainly driven from the protostar. The LVC and HVC are strongly coiled by the magnetic field lines anchored to the first core and protostar, respectively. Our results show that the flow appearing around the first core has a wide opening angle and slow speed, while the flow appearing around the protostar has a

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Fig. 1 Bird’s-eye view. The structure of high-density region ( > 0:1 c ; red iso-density surface), and magnetic field lines (black-and-white streamlines) are plotted in each panel. The structures of the jet (v > 7 km s1 ) and outflow (v > 0:5 km s1 ) are shown by iso-velocity surfaces, respectively. The density contours (false color and contour lines), velocity vectors (thin arrows) on the mid-plane of x D 0, y D 0, and z D0 are, respectively, projected in each wall surface

Fig. 2 Schematic view of the jet and outflow driven from the protostar and the first core, respectively

well-collimated structure and high speed, as shown in Fig. 2. The speed difference is caused by the difference of the depth in the gravitational potential. The flow speed corresponds to the Kepler speed of each object. Because the first core has a shallow gravitational potential, its flow is slower. The flow driven from the protostar, which has a deeper gravitational potential, has a high speed. In our calculations, the low- and high-velocity flows have speeds of vLVF ' 3 km s1 and vHVF ' 30 km s1 , respectively. These speeds are slower than those of observations. Typically, observed molecular outflow and optical jet have speeds of vout;obs ' 30 km s1 , and vjet;obs ' 100 km s1 , respectively. However, since the first and second cores (protostar) have mass of Mfc D 0:01 Mˇ and Msc ' 103 Mˇ , re-

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spectively, at the end of the calculations, each core increases its mass in the gas accretion phase. The Kepler speed increases with the square root of the mass. When the mass of each core increases by 100 times, the Kepler speed increases 10 times. Thus, the speed of the low- and high-velocity flows may increase by 10 times, and reach vLVF ' 30 km s1 and vHVF D 300 km s1 , respectively, which correspond to typical observed values.

4 Discussion Observation shows that the molecular outflows have wide opening angles and low flow speeds, while the optical jets have good collimation and high flow speeds. Molecular outflow has been considered to be entrained by the optical jet driven from a circumstellar disk around the protostar. In this study, we calculated the cloud evolution from the molecular cloud core to protostar formation, and found that two distinct flows are driven from different objects, and the observed features of molecular outflow and optical jet were naturally reproduced. Thus, we expect that the low-velocity flow from the first core corresponds to the molecular outflow, while the high-velocity flow from the protostar corresponds to the optical jet. The different collimation of low- and high-velocity flow is caused both by the configuration of the magnetic field lines around the drivers and their driving mechanisms. The magnetic field lines around the first core have an hourglass configuration because they converge to the cloud center as the cloud collapses, and Ohmic dissipation is ineffective before the first core formation. In addition, the centrifugal force is more dominant than the Lorentz force in the low-velocity flow (molecular outflow). Thus, the flow appearing near the first core is mainly driven by the magnetocentrifugal wind mechanism (disk wind). On the other hand, near the protostar, the magnetic field lines are a straight, and the magnetic pressure gradient mechanism is more effective for driving the high-velocity flow (optical jet). The magnetic field lines straighten by the magnetic tension force near the protostar because the magnetic field is decoupled from the neutral gas. However, the magnetic field lines are strongly twisted in the region in close proximity to the protostar, where the magnetic field is coupled with the neutral gas again. Thus, the strong toroidal field generated around the protostar can drive the high-velocity flow (optical jet), which is guided by the straight configuration of the magnetic field. Our calculations do not completely reject the well-known concept that the observed molecular outflow is entrained by the optical jet, because we calculate the formation of the jet and outflow only in the early star-formation phase. Further long-term calculations are needed to understand the mechanism of the optical jet and molecular outflow in more detail.

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References 1. Hirth, G. A., Mundt, R., & Solf, J. 1997, A&AS, 126, 437 2. Pyo, T., et al. 2003, ApJ, 590, 340 3. Arce, H. G., Shepherd, D., Gueth, F., Lee, C.-F., Bachiller, R., Rosen, A., & Beuther, H. 2007, in Protostars and Planets V, ed. B. Reipurth, D. Jewitt, & K. Keil (Tucson: Univ. Arizona Press), 245 4. Pudritz, R. E., Ouyed, R., Fendt, C., & Brandenburg, A. 2007, in Protostars and Planets V, ed. B. Reipurth, D. Jweitt, & K. Keil (Tucson: Univ. Arizona Press), 277 5. Belloche, A., And´e, P., Despois, D., Blinder, S. 2002, A&A, 393, 927 6. Bally, J., Reipurth, B., & Davis, C. J. 2007, in Protostars & Planets V, ed. B. Reipurth, D. Jewitt, & K. Keil (Tucson: Univ. Arizona Press), 215 7. Mundt, R., & Fried, J. W. 1983, ApJ, 274, L83 8. Machida, M. N., Matsumoto, T., Hanawa, T., & Tomisaka, K. 2006a, ApJ, 645, 1227 9. Machida, M. N., Inutsuka, S., & Matsumoto, T., 2006b, ApJ, 647, 151 10. Machida, M. N., Inutsuka, S., & Matsumoto, T., 2007, ApJ, 670, 1198 11. Machida, M. N., Tomiska, K., Matsumoto, T., & Inutsuka, S., 2008c, ApJ, 677, 327 12. Machida, M. N., Inutsuka, S., & Matsumoto, T., 2008, ApJ, 676, 1088 13. Machida, M. N., Tomisaka, K., & Matsumoto, T., 2004, MNRAS, 348, L1 14. Machida, M. N., Matsumoto, T., Tomisaka, K., & Hanawa, T. 2005a, MNRAS, 362, 369 15. Larson, R. B., 1969, MNRAS, 145, 271 16. Masunaga, H., Miyama, S. M., & Inutsuka, S., 1998, ApJ, 495, 346 17. Nakano, T., Nishi, R., & Umebayashi, T. 2002, ApJ, 573, 199

Outflow Driven Turbulence in Star Forming Clouds Adam Frank

Abstract Setting young stellar object jets and outflows in their broadest context requires an understanding of outflows as “feedback” in the development of molecular cloud turbulence and the determination of star formation efficiencies. In this contribution I review our group’s recent studies exploring relationships between protostellar outflows and turbulence in molecular clouds. We first present studies of turbulence and fossil cavities driven by YSO outflows using numerical simulations which track the evolution of single transient jets driven into a turbulent medium. Our simulations show both the effect of turbulence on outflow structures and, conversely, the effect of outflows on the ambient turbulence. These studies demonstrate that individual transient outflows have the capacity to re-energize decaying turbulence. Next we present simulations of multiple interacting jets. We show that turbulence can readily be sustained by these interactions and show that it is possible to broadly characterize an effective driving scale of the outflows. Comparing the velocity spectrum obtained in our studies to that of an isotropically forced control we show that in outflow driven turbulence a power law of the form E.k/ / k ˇ is indeed achieved. However we find a steeper spectrum ˇ  3 is obtained in outflow driven turbulence

A. Frank () Dept. of Physics and Astronomy, University of Rochester, Rochester, NY 14627 e-mail: [email protected]

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models than in isotropically forced simulations ˇ  2:0. Taken together both studies provide broad support for the conclusion that fossil cavities driven by decaying jets can provide a source of turbulence and feedback which mediate star formation processes in molecular clouds. Whether this does obtain in real clouds remains a point which must be demonstrated

1 Introduction To set YSO jets in context one must first determine the proper scale at which the question is to be asked. Jets as a phenomena extend across an astonishing range of spatial (and temporal) scales. Different kinds of physics and different issues arise depending on the scale and, hence context, one seeks to understand. “Feedback” – jet driven changes in the environment which gave rise to the jet – is a fascinating and important problem for jet studies as its speaks directly to the issue of what role jets play in the ecology of star formation. If no feedback occurs on any scales then jets are, essentially, an epiphenomena of star formation and while lovely to look at may not be particularly important scientifically. If, however, jets play an important role in determining properties of the disks, envelopes and clouds from which they are born then their study becomes more than a kind of cosmic art history. When discussing issues of jet feedback one can distinguish between three scale lengths. First there is “micro-scale” feedback which concerns the effect of outflows on their own launch scale (L < 10 AU) meaning their effect on disks, stars and disk/stellar magnetospheres. There is a great deal of study in these domains looking at issues such as jet angular momentum extraction from disks [5, 16]. Next we have “mesoscale” feedback which relates the impact of jets on environments associated with infalling envelopes (L < 104 AU). Considerable work has focused on these domains: i.e. the way outflows can sculpt envelopes [12,10] or their influence on the accretion onto the star [23]. Finally “macro-scale” feedback, which is the concern of this paper, is associated with the effect of outflows on the scale of star forming clusters or the entire molecular cloud itself. Macro-scale feedback is relevant to two important, related issues facing modern theories of star formation. These are the nature of turbulence in clouds and the relative inefficiency of star formation. For excellent reviews see [14, 34]. Surprisingly low values of star formation efficiency SFR are observed (typically ranging from 0.01 to 0.1). The low values of SFR can be accounted for using some form of support, such as supersonic turbulence within the cloud to keep it from forming stars. But while turbulence can provide isotropic pressure, both hydrodynamic and MHD turbulence decay quickly [41, 29]. Thus if clouds are stable, long-lived structures supported against self-gravitational collapse by turbulence, those turbulent motions must be continually driven either internally via gravitational collapse and stellar feedback or externally via turbulence in the ISM. Feedback in the form of stellar outflows is one means of driving turbulence in molecular clouds [36]. Observational studies make it clear that combined outflow

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energy budgets are sufficient to account for cloud turbulent energy [2, 3, 20, 40, 45]. Analytical work by [31, 32, 33, 22, 39] has explored the role of feedback on clouds and recent simulations [27] and [35] have mapped out the complex interplay between star formation and turbulence concluding that outflows could re-energize turbulence. However, studies of single jets by [4] came to the opposite conclusion in the sense that single jets will not leave enough supersonic material in their wakes to act as a relevant source of internal forcing. While the potential for protostars to act as a means of internal forcing remains attractive, many outstanding questions remain open. The energetics of turbulent resupply must be understood in terms of a coupling efficiency between outflow and cloud. In addition most observational studies of protostellar turbulence have focused on active outflows. These may not, however, be the only means of coupling. This contribution builds on a series of works by our group beginning with Quillen et al. [40] who explored turbulent motions and outflow activity in NGC 1333. That study demonstrated that fossil cavities rather than active outflows are responsible for driving turbulence. The relation between turbulence, active and fossil outflows was further studied in Cunningham et al. [8, 9].

1.1 Single Jets and Turbulence In Cunningham et al. [7] we explored the ability of individual transient jets to return their energy and momenta to pre-existing turbulent backgrounds through direct numerical simulation. These simulations used the AstroBEAR AMR multi-physics code Cunningham et al. [11]. Each simulations was carried forward on a periodic domain initialized with a decaying turbulent environment. A jet was then driven into the grid for a timescale determined by toff . Table 1 presents the parameters of these simulations including our control runs which had no either jet or no turbulence.

1.2 Jet Morphology Morphological effects of the turbulent medium on the jet-driven bow-shock are readily apparent in our study (Fig. 1) despite the comparatively slow turbulent speed in the ambient medium, (vturb is only 10% of the jet propagation speed). In the jet Table 1 Parameters for single jet/turbulence runs Run toff Turbulence

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which is driven into a quiescent environment (run J) a thin shock-bounded Mach disk structure forms at the head of the outflow as expected when radiative cooling dominates. When the same jet is imposed into a turbulent environment (run 3), the Mach disk is observed to broaden significantly. As the 3-D rendering of the simulation (Fig. 2) shows, the bow shock eventually fragments entirely. Our result suggest an interpretation of jet evolution in turbulent media which can account for the eventual disruption of fossil cavities and a mechanism for transferring bulk, directed jet motions into randomized turbulent ones. It is well known thin, shock bounded layers can be disrupted via instabilities [37, 43, 42]. Thus turbulent eddies occurring across scales provide a space filling environment of multi-mode perturbation seeds. These seeds act to initiate instability growth in the bow shock (delineating the fossil cavity) and facilitate the coupling of energy and momentum from the original outflow to the turbulent environment. For example, the growth rate  of Non-Linear Thin Shell instabilities of an initial p displacement L and wavenumber k can be approximated as [18]   Cd1=2 ck kL here Cd is of order 1. Considering only the non-linear thin shell instability with L  h and k  1= h we see that for jet of velocity vj propagating across a cloud of size rc l in time tc l rcl c rcl  M 1 : (1)  tcl  vj h h

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Given typical conditions typical of YSO jets of M  100 and rcl  1 pc we have  tcl >> 1. Thus perturbation seeds provided by the turbulence will lead to the fragmentation of the bow shock by the time the cavity has reached is maximum extent. The fragmentation of the bow shock is the mechanism which randomizes its bulk momentum and turning a laminar flow into a turbulent one.

1.3 Turbulence and Kinetic Energy Power Spectra Plots of velocity power spectra, (left panels in Fig. 1), show the degree to which the outflows interact with the background turbulence. The jet/turbulence simulations show increasing power in the largest flow eddies (smallest wavenumber). This indicates that the outflow cavities in provide sufficient power at length scales comparable to the width of the simulation domain to support turbulent motions in the ambient flow against energy decay. The results can be interpreted as the net result of the decay of turbulent energy at large wavenumber balanced against with the resupply of turbulent energy by the injected outflow scales at smaller wavenumber. One important point to be taken from the plots is that the simulation allow us identify the driving scale of turbulent resupply associated with the scale of cavity propagation as kdrive  1=LBS . In addition, comparison of the turbulent power law slopes shows that the “no jet” simulation (ˇ D 2:09) and the “jet-only” model (ˇ D 2:46) bracket the range of inertial range fall-off. The outflow models with ambient turbulence fill in this range with a clear steepening appearing increasing outflow energy (Table 1). This can be interpreted as an outflow-induced suppression of sub-cavity scale flow eddies. The power spectra therefore indicate that outflow cavities act to support the ambient turbulence in two ways. First the outflows power eddies of comparable extent to the length of the cavity LBS and second, the opening of outflow cavities inhibit the cascade of energy to sub-outflow scales. Finally note that the outflow cavity driven by the short, nearly impulsive jet in run 1 has been completely subsumed into turbulent eddies. By the end of the simulation the flow field for the short jet pulse (run 1) is qualitatively similar to that of the turbulence-only control (run 0). This is the behavior that was expected by Quillen et al. [40] and Cunningham et al. [9] for long extinct outflow structures embedded in turbulence. Furthermore, the net mechanical energy of the disrupted outflow cavity in run 1 decays at a rate comparable to that of the turbulence-only control in run 0. Thus the disrupted outflow cavity becomes turbulent itself as it evolves under the influence of the turbulent environment.

2 Multiple Jet Driven Turbulence In this section we describe work by Carrol et al. [6] on the dynamics of cloud material set in motion by multiple transient collimated jets. The resolution of the outflows in our work is higher than those used in previous studies allowing us to tease out

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their dynamics and its role driving turbulence. Following the dimensional analysis of [33] we consider a cloud of mean density 0 , with outflows occurring at a rate per volume S with momentum I. This allows one to define characteristic outflow scales 4=7 3=7 3=7 I I 1=7 ; T D I 3=70S 4=7 Comof mass, length, and time: M D 0S 3=7 ; L D 1=7 1=7 0

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supersonic characteristic velocities which, in turn, indicates that outflows contain enough momentum to drive significant supersonic turbulence. Our simulations seek to explore a detailed realization of this idea. Note these relations are for spherical outflows and provide appropriate order of magnitude estimates for collimated outflows. Figure 3 shows density cross cuts from both the outflow driven simulation and an isotropically driven control run. Examination of the outflow driven case shows that outflow cavities are initially able to expand into the quiescent environment without being disrupted and easily grow to pc size structures. Since the outflows are bipolar, the total vector momentumR injected by each outflow is zero, however, the total scalar momentum defined as jPj d V grows steadily until the cavity interactions begin to dissipate momentum in the expanding shells. By 1:8T the cavities have entirely filled the domain. By 3T the system has reached a statistically quasisteady state. Comparison of the late time outflow-driven simulation image and the isotropically-forced simulation image in Fig. 3 show that both have reached states of highly disordered flows with structure present on a variety of scales. As we show below the visual impression of turbulence in both cases is supported by statistical measures of the flows as well as their time evolution. The comparison by visual inspection however is noteworthy because of the “holes” which appear in the outflow driven density distribution. These are created by fossil outflow cavities and exist until the cavities are subsumed by the turbulent motions [7] or interact with another cavity. As discussed above such shells have been observed in turbulent flows around star forming regions like NGC1333 [40] and point to an important morphological, rather than statistical, signature of outflow-driven turbulence.

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In Fig. 4 we plot the one-dimensional velocity power spectra of the outflowdriven turbulence simulation together with the isotropically forced turbulence simulation (with injection scales [2.1 – 2.4]L). Consideration of the right hand panel demonstrates that the spectrum of the outflow driven turbulence is quite different from the isotropically forced turbulence. At large scales, the outflow driven spectrum rises slowly from the box scale to around K before turning over and steepening to ˇ D 3:2 at  3K. The power law behavior continues all the way to the dissipation scale. The isotropically forced spectrum shows a sharp peak at the driving scale and then falls with shallower power law appropriate to a Burgers model with ˇ D 2 [21]. The difference between the two spectra lies in the nature of the driving. Since the shocks in the outflow driven turbulence are not randomly distributed it should not be expected to follow a Burgers model power law. To understand the dynamic driving the steeper spectrum in the outflow simulations note that at sub-outflow scales the slope (ˇ  3 corresponds to a velocity length scaling v.l/ / l) that is consistent with the Hubble type flow seen in the expanding cavities. Thus it is the individual cavities which change the powerlaw. The time evolution of the spectra (for many outflows) shown in the right hand panel of Fig. 4 provides some insight into the behavior. The steep slope of the spectrum at sub-outflow scales is present already by 0:70T before the outflows have begun to significantly interact. This is consistent with the interpretation that the steep slope of the spectrum at later times 3T arises from the volume swept-up by the multiple fossil outflow cavities rather than their interaction. In other words, we can interpret the steep spectra as due to expanding outflows that sweep up small-scale (high k) eddies which, in turn, effectively removes them from the flow. This effect reduces the observed power on these scales and steepen the slope of the energy spectrum. In spite of the differences in slope the development of a steady turbulent powerlaw spectrum, Fig. 4, along with the development of log normal density distributions (not shown), demonstrate one of the key conclusions of our study: transient outflow cavities can set the bulk of initially quiescent material into random but statistically steady supersonic motions. [44]

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3 Summary Our results add weight to the argument that protostellar outflows can play a significant role in altering their star-forming environments [35]. By acting to drive turbulence and, possibly, setting the observed values of the Star Formation efficiency outflows may be a kind of thermostat acting as the agent of self-regulation from cloud to disk and back again. Future work will need to take these studies further by moving both towards more realistic simulations at high resolution and towards a deeper understanding of the contribution individual physical processes play in self-regulation if it does occur. In addition it is important to find observational laboratories that are isolated enough to allow for clean tests of outflow-driven feedback theory. One such laboratory may be the L1551 cloud. Recent studies show that L1551 shows clear evidence for strong modification by its outflows [38] and we are currently engaged in an attempt to reproduce it properties (Yirak et al. 2008 in preparation).

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27. Li, Z. Y. & Nakamura, F. 2006, ApJ, 640, 187 28. Mac Low, M.-M., Klessen, R. S., Burkert, A., & Smith, M. D. 1998, Phys Rev Lett, 80, 2754 29. Mac Low, M.-M. 1999, ApJ, 524, 169 30. Mac Low, M.-M. 2000, Star Formation from the Small to the Large Scale, 445, 457 31. Matzner, C. D. 2000, ApJ, 545, 364 32. Matzner, C. D. 2001, AAS, 198, 9602 33. Matzner, C. D. 2007, ApJ, 659, 1394 34. McKee, C. F., & Ostriker, E. C. 2007, ARAA, 45, 565 35. Nakamura, F., & Li, Z. Y. 2007, ApJ, 662, 395 36. Norman, C., & Silk, J. 1980, ApJ, 238, 158 37. Sutherland, R. S., Bicknell, G. V., & Dopita, M. A. 2003, ApJ, 591, 238 38. Swift, J. J., & Welch, W. J. 2008, ApJS, 174, 202 39. Tan, J. C., Krumholz, M. R., & McKee, C. F. 2006, ApJ, 641, L121 40. Quillen, A. C., Thorndike, S. L., Cunningham, A. J., Frank, A., Gutermuth, R. A., Blackman, E. G., Pipher, J. L., & Ridge, N. 2005, ApJ, 632, 941 41. Stone, J. M., Ostriker, E. C., & Gammie, C. F. 1998, ApJ, 508 L99 42. Vishniac, E. T. 1994, ApJ, 428, 186 43. Vishniac, E. T., & Ryu, D. 1989, ApJ, 337, 917 44. Vazquez-Semadeni, E. 1994, ApJ, 423, 681 45. Warin, S., Castets, A., Langer, W., Wilson, R., & Pagani, L. 1996, AA, 306, 935

Jet Driven Turbulence? Robi Banerjee, Susanne Horn, and Ralf S. Klessen

Abstract Molecular clouds, the birth sites of stars, are permeated by supersonic gas motions. Here, we summarize our results from numerical simulations on individual jets interacting with their ambient medium. These single jet simulations show that the volume filling factor of supersonic turbulence excited by jets is very low. In general, supersonic motions, if driven at small scales, do not propagate far from their source and are damped quickly. Therefore, it is unlikely that the supersonic motions observed in molecular clouds are maintained by jets launched from protostellar objects. Our preliminary results from three dimensional simulations of collapsing cloud cores with a self-consistent description of mechanical feedback around protostellar objects point towards the same results.

R. Banerjee (), S. Horn, and R.S. Klessen Institute for Theoretical Astrophysics, Zentrum f¨ur Astronomie, University of Heidelberg, Albert-Ueberle-Str. 2, 69120 Heidelberg, Germany e-mail: [email protected]; [email protected]; [email protected]

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1 Introduction The interstellar medium (ISM) and star forming molecular clouds are permeated by turbulent, supersonic gas motions, e.g. [11,6,16,2] and references therein. However, it is known that supersonic turbulence decays quickly and has to be continuously driven to be maintained [17, 25, 22]. In principle, the energy input can be supplied from inside the molecular cloud (e.g. radiation from massive stars, jets, outflows) or from outside (e.g. supernovae, large scale gas streams). In particular, the birth of stars is in most cases, if not all, accompanied by outflows and high velocity jets. Protostellar jets propagate with velocities of about 300 km s1 as seen in the radial velocity shift of forbidden emission lines and proper motions of jet knots. Many of these jets remain highly collimated with opening angles less than 5ı over a distance up to several pc [19, 24]. Norman and Silk [21] proposed that, these Herbig-Haro (HH) outflows could power the turbulent energetics in molecular clouds and star formation could be a self-regulated process. The amount of turbulent energy controls the strength of gravitational collapse and subsequent star formation activity, e.g. [16]. Estimates of the individual power of jets show that about 100 protostellar objects release an amount of energy during their active phase which is comparable to the total kinetic energy of a star forming region with about 103 Mˇ of gas, see e.g. [3]. Once the energy is transfered to turbulent energy it decays on several rms crossing times L=vrms , which is considerably larger than the lifetime of individual jets. The energy inserted by individual jets and outflows therefore will remain in the cloud region during most of its star-formation period, which typically lasts a few 105 years [9, 16]. This indicates that protostellar jets and outflows are indeed able to sustain the gas turbulence. Whether the kinetic energy from the jets can be transferred to the ambient cloud material efficiently and with the right spatial and temporal characteristics is a debated question, see [15, 18, 23, 15, 14, 20] for different conclusions on this subject. Here, we summarise the results from the study of [3] where multidimensional simulations were used for a detailed investigation on the interactions of single jets with their ambient gas. Furthermore we discuss preliminary results from our global simulations of collapsing cloud cores with a self-consistent mechanical feedback from protostellar objects.

2 Turbulence from Single Jets The simulations of [3] are performed with the adaptive mesh refinement (AMR) code FLASH [7], in which the jets are modeled as a kinetic energy injection from the box boundary. This investigation is based on a detailed parameter study with varying jet speeds, density contrasts, and different magnetic field configurations. In Fig. 1 we show the time evolution of a typical jet configuration. Indeed, the jets interacting with the ambient medium excite instabilities and turbulent motions.

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Fig. 1 Density (top) and velocity (bottom) evolution of a Mach 5 jet (run M5c) at two different times, t D 3:0 (left) and t D 5:0 (right). The jet is continuously powered and runs into a homogeneous medium. The jet develops knots from reflections off the jet edge. This structure propagates also into the ambient media. Additionally, Kelvin-Helmholtz instabilities develop at the edge of the jet. The turbulent flow outside the jet is mainly subsonic [from [3]]. A color version of this plot can be found in the electronic version of this book and in Appendix A (Fig. A.30)

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Fig. 2 Shows the time evolution of the velocity PDFs from a continuously driven jet (left panel) and the time evolution of the kinetic energy in the case of a transient jet (right panel). The jet excites very little supersonic fluctuations (the peak at v=c D 5 is the injected jet itself). Furthermore, supersonic motions decay much faster than subsonic waves. [from [3]]

The instabilities, mainly from Kelvin-Helmholtz modes, develop at the edge of the jet and propagate into the surrounding gas. But, the associated velocities of these instabilities are sub-sonic. Even the bow shock region does not develop high velocity fluctuations. This can be seen from Fig. 2 where probability distribution functions (PDFs) of the excited velocity fluctuations are shown. The volume filling factor of

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excited supersonic motions is very small compared to the subsonic fluctuations. The only prominent supersonic feature in the velocity PDF is the jet itself. This is the result of the very nature of supersonic motions excited from a small region. Such motions cannot propagate far because they dissipate energy effectively in compression. The re-expansion of density fluctuations can only excite subsonic motions. What is left are normal (subsonic) turbulent waves. Furthermore, the supersonic fluctuations decay much faster than normal sound waves, as can be seen from the right panel of Fig. 2.

3 Global Collapse with Mechanical Feedback Although, detailed studies on single jet excited turbulence are very instructive, it is necessary to self-consistently embed such jets and outflows in simulations of global star formation to fully quantify their influence on this process. We are pursuing such 3D simulations with our version of the FLASH code, which includes now accreting sink particles to model collapsing (i.e. protostellar) regions. We use the properties of sink particles to launch collimated outflows, with strengths that are directly linked to the actual accretion rate onto the sinks ,see e.g. [5,26] for links between accretion rates and wind thrust. We compare the outcome of our collapse simulations with and without mechanical feedback in Fig. 3 and the left panel of Fig. 4. From the column density images and the computed velocity PDFs we conclude that the kinetic energy input from the outflows do not alter the global evolution of a collapsing cloud core that is forming

6.0221e+04 yr

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Fig. 3 Shows the comparison of the collapse of a turbulent cloud core without (left) and with (right) mechanical feedback (i.e. collimated outflows). The snapshots are taken at an early time into the evolution (t  0:6 tff , where the free-frall time tff  105 year). Shown is the column density of the cloud and the star formation sites (i.e. sink particles marked as black dots). The overall structure in both cases is the same, differences are only noticeable at small scales. A color version of this plot can be found in the electronic version of this book and in Appendix A (Fig. A.31)

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Fig. 4 The left panel shows the velocity PDFs in the case without (red) and with (blue) outflows from protostellar regions at the time t  603 years (see also Fig. 3). The fraction of outflow powered velocity fluctuations is small compared to the overall turbulent motions. In the right panel we show the further evolved cloud where outflows are launched from protostellar regions. A color version of this plot can be found in the electronic version of this book and in Appendix A (Fig. A.32)

star clusters. Even from these early results it is questionable whether jets and outflows regulate star formation. But, for a final conclusion, we have to investigate this topic further, in particular we will have to run these global simulations for longer times, to find out if a statistical equilibrium will be established.

4 Summary and Conclusion Based on our studies on single jet driven turbulence we conclude that collimated jets from young stellar objects are unlikely drivers of large-scale supersonic turbulence in molecular clouds. These conclusions are also supported by our early results of global, turbulent collapse simulations, where collimated outflows are launched selfconsistently from protostellar regions (modelled as sink particles). Alternatively the cloud’s turbulence can be powered by large scale flows which might be responsible for the formation of the cloud itself, e.g. [1, 8, 4]. Energy cascading down from the driving scale to the dissipation scale will then produce turbulent density and velocity structure in the inertial range in between [12]. If the large-scale dynamics of the interstellar medium is driven by gravity, as suggested, e.g., by [10, 13] gravitational contraction would also determine to a large extent the internal velocity structure of the cloud. Otherwise, blast waves and expanding shells from super novae are also viable candidates to power supersonic turbulence in molecular clouds, see e.g., [16]. Acknowledgements RB is funded by the Deutsche Forschungsgemeinschaft (DFG) with the Emmy-Noether grant BA 3706/1.

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References 1. Ballesteros-Paredes J, Hartmann L, V´azquez-Semadeni E (1999) Turbulent Flow-driven Molecular Cloud Formation: A Solution to the Post-T Tauri Problem? ApJ 527:285–297, DOI 10.1086/308076, arXiv:astro-ph/9907053 2. Ballesteros-Paredes J, Klessen RS, Mac Low MM, Vazquez-Semadeni E (2007) Molecular Cloud Turbulence and Star Formation. In: Reipurth B, Jewitt D, Keil K (eds) Protostars and Planets V, pp 63–80 3. Banerjee R, Klessen RS, Fendt C (2007) Can Protostellar Jets Drive Supersonic Turbulence in Molecular Clouds? ApJ 668:1028–1041, DOI 10.1086/521097, arXiv:0706.3640 4. Banerjee R, Vazquez-Semadeni E, Hennebelle P, Klessen R (2008) The early stages of molecular cloud evolution in the magnetised interstellar medium: clump morphology and evolution. ArXiv e-prints 0808.0986 5. Cabrit S, Bertout C (1992) CO line formation in bipolar flows. III - The energetics of molecular flows and ionized winds. A&A 261:274–284 6. Elmegreen BG, Scalo J (2004) Interstellar Turbulence I: Observations and Processes. ARA&A 42:211–273 7. Fryxell B, Olson K, Ricker P, Timmes FX, Zingale M, Lamb DQ, MacNeice P, Rosner R, Truran JW, Tufo H (2000) FLASH: An Adaptive Mesh Hydrodynamics Code for Modeling Astrophysical Thermonuclear Flashes. ApJS 131:273–334 8. Hennebelle P, Banerjee R, Vazquez-Semadeni E, Klessen R, Audit E (2008) From the warm magnetized atomic medium to molecular clouds. ArXiv e-prints, accepted for publication in A&A 0805.1366 9. Klessen RS (2003) Ludwig Biermann Award Lecture: Star Formation in Turbulent Interstellar Gas (With 11 Figures). In: Schielicke RE (ed) Reviews in Modern Astronomy, Reviews in Modern Astronomy, vol 16, pp 23 10. Li Y, Mac Low MM, Klessen RS (2005) Star Formation in Isolated Disk Galaxies. I. Models and Characteristics of Nonlinear Gravitational Collapse. ApJ 626:823–843, DOI 10.1086/430205, arXiv:astro-ph/0501022 11. Larson RB (1981) Turbulence and star formation in molecular clouds. MNRAS 194:809–826 12. Lesieur M (1997) Turbulence in Fluids. Kluwer Academic Publishers, Dordrecht 13. Li Y, Mac Low MM, Klessen RS (2006) Star Formation in Isolated Disk Galaxies. II. Schmidt Laws and Efficiency of Gravitational Collapse. ApJ 639:879–896, DOI 10.1086/499350, arXiv:astro-ph/0508054 14. Li ZY, Nakamura F (2006) Cluster Formation in Protostellar Outflow-driven Turbulence. ApJ 640:L187–L190, DOI 10.1086/503419, astro-ph/0512278 15. Mac Low MM (2000) Turbulence Driven by Stellar Outflows. In: Favata F, Kaas A, Wilson A (eds) ESA SP-445: Star Formation from the Small to the Large Scale, pp 457 16. Mac Low MM, Klessen RS (2004) Control of star formation by supersonic turbulence. Reviews of Modern Physics 76:125–194 17. Mac Low MM, Klessen RS, Burkert A, Smith MD (1998) Kinetic Energy Decay Rates of Supersonic and Super-Alfv´enic Turbulence in Star-Forming Clouds. Physical Review Letters 80:2754–2757, astro-ph/9712013 18. Matzner CD (2007) Protostellar outflow-driven turbulence. e-print: astro-ph/0701022, accepted by ApJ astro-ph/0701022 19. Mundt R, Buehrke T, Solf J, Ray TP, Raga AC (1990) Optical jets and outflows in the HL Tauri region. A&A 232:37–61 20. Nakamura F, Li ZY (2007) Protostellar Turbulence Driven by Collimated Outflows. ApJ 662:395–412, DOI 10.1086/517515, arXiv:astro-ph/0703152 21. Norman C, Silk J (1980) Clumpy molecular clouds - A dynamic model self-consistently regulated by T Tauri star formation. ApJ 238:158–174, DOI 10.1086/157969 ˚ (1999) Supersonic Turbulence in the 22. Padoan P, Bally J, Billawala Y, Juvela M, Nordlund A Perseus Molecular Cloud. ApJ 525:318–329, astro-ph/9905383

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23. Quillen AC, Thorndike SL, Cunningham A, Frank A, Gutermuth RA, Blackman EG, Pipher JL, Ridge N (2005) Turbulence Driven by Outflow-blown Cavities in the Molecular Cloud of NGC 1333. ApJ 632:941–955, DOI 10.1086/444410, astro-ph/0503167 24. Raga A, Cabrit S, Dougados C, Lavalley C (2001) A precessing, variable velocity jet model for DG Tauri. A&A 367:959–966, DOI 10.1051/0004-6361:20000415 25. Stone JM, Ostriker EC, Gammie CF (1998) Dissipation in Compressible Magnetohydrodynamic Turbulence. ApJ 508:L99–L102, DOI 10.1086/311718, astro-ph/9809357 26. Wu Y, Wei Y, Zhao M, Shi Y, Yu W, Qin S, Huang M (2004) A study of high velocity molecular outflows with an up-to-date sample. A&A 426:503–515, DOI 10.1051/0004-6361:20035767, arXiv:astro-ph/0410727

Part VII

JETSET Early Stage Researcher Presentations

Prospects for Outflow and Jet Science with ALMA John Richer

Abstract The study of the physics of protostellar jets and outflows is at present significantly limited by the angular resolution and sensitivity of our telescopes. ALMA, the Atacama Large Millimeter/submillimetre Array, is under construction in northern Chile. It will provide a dramatic increase in sensitivity, resolution and spectral bandwidth at millimetre and submillimetre wavelengths compared to current facilities. I discuss the potential breakthroughs in jet science that ALMA will enable.

1 Introduction Many recent breakthroughs in our understanding of protostellar jets and outflows have come from high-resolution images at optical and infrared wavelengths. We are now used to analysing images and spectra at up to 0.1 arcsec resolution, using ground-based AO techniques and the Hubble Space Telescope. At the nearest regions of star formation, at a distance of 140 pc, these data allow us to probe physical

J. Richer () Cavendish laboratory, J J Thomson Avenue, Cambridge e-mail: [email protected]

K. Tsinganos et al. (eds.), Protostellar Jets in Context, Astrophysics and Space Science Proceedings, DOI 10.1007/978-3-642-00576-3 51, c Springer-Verlag Berlin Heidelberg 2009 

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scales in jets and disks down to approximately 10–20 AU (for references, see the contribution by Bacciotti in this volume). However, such data limit us to the study of the warmer partially-ionised gas which emits in the optical and IR regime, and to the study of systems which are not deeply embedded in dust. To study the youngest systems, where the protostellar mass is small and vigorous accretion is ongoing, we need to study the deeply embedded Class 0 protostars at longer wavelengths. To date, studies of the molecular outflows and jets from such systems have produced some spectacular results (e.g. [6]). We know that outflows have a wide range of physical structures, and that they drive a rich chemistry. The energetics derived from the data shed vital clues as to the nature of the central engine. However, existing millimetre-wave interferometers such at the IRAM Plateau de Bure, the SMA and OVRO typically achieve an angular resolution of only 0.5–1.0 arcsec: higher angular resolution requires longer baselines, more collecting area and an ability to correct atmospheric seeing. In order to study the details of jet physics, including the jet launch point and collimation scales, we need to achieve one to two orders of magnitude better angular resolution.

2 Capabilities of ALMA 2.1 Design The Atacama Large Millimeter/submillimetre Array, ALMA [9], is a highfrequency radio interferometer comprising 66 precision antennas on a high-altitude site in northern Chile. Its use for studies of jets has recently been discussed by Shepherd [8]. ALMA is designed to perform diffraction-limited imaging at frequencies from 31 to 950 GHz, using baselines from 15 m to approximately 15 km. The design includes 50  12-m diameter antennas in the main array; in addition, an array of 12  7-m plus 4  12-m antennas forms the Atacama Compact Array (ACA) which improves the imaging of large angular size sources. The high-altitude site, Llano de Chajnantor at 5000-m above sea level, provides a large flat area for the antennas to be laid out, and a very dry atmosphere which permits operation with good sensitivity to frequencies of up to and beyond 1 THz. ALMAs 66 antennas provide very good imaging performance. The antennas can each be moved by a transporter vehicle to one of more than 200 antenna pads. Approximately 36 antenna configurations have been designed, varying in maximum baselines from 150 m to 15 km. The ALMA digital correlator can measure up to 2016 simultaneous cross correlations from the antennas, providing extremely good uv coverage even for snapshot imaging. The angular resolution achievable varies from roughly 3 to 0.03 arcsec at 100 GHz depending on the configuration chosen Fig. 1. The instantaneous field of view of ALMA is relatively small, roughly 60 arcsec 100 GHz. Larger sources can be imaged by mosaicking together separate fields,

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Fig. 1 The 66 antennas of ALMA on the Chajnantor plane in Chile. (Artist’s conception, courtesy of European Southern Observatory.)

and this will be a common mode of observation especially for galactic sources. In addition, the ACA’s antennas can be used to measure the zero and short spacing data required to fill the central hole in the uv plane, which has a radius of about 15 m. By adding the ACA data to those from the main array, high-fidelity imaging with complete uv coverage out to the maximum baseline can be achieved. ALMA’s receivers will cover, in 10 bands, the frequency range 31–950 GHz, although not all bands will be available from first light. All receivers offer 8 GHz bandwidth in two polarizations, and the correlator can measure all the correlations to allow full IQUV polarisation imaging. The bandwidth can be measured at high spectral resolution, with typically 8192 channel available for high resolution spectroscopy. ALMA uses specialised ‘adaptive optics’ techniques to correct atmospheric phase errors caused principally by water vapour fluctuations. The antennas can move very quickly to phase calibrators (typically quasars) to derive the atmospheric phase contribution. In addition, water vapour radiometers operating at 183 GHz on each antenna provide direct high-time-resolution measurements of the water vapour fluctuation which can be used to correct the short timescale phase errors due to water in the troposphere.

2.2 Scientific Capabilities The general capabilities of ALMA for galactic science are reviewed in Richer [7]. In Table 1, I present the predicted line sensitivity of ALMA after 8-hours integration on a single field, at five selected frequencies corresponding to the CO rotational transitions J=1–0, 2–1, 3–2, 4–3, and 6–5 (column 1). For each frequency, three array sizes which span the extremes of the possible configurations are shown (column 2). Bear

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J. Richer Table 1 Predicted line sensitivity of ALMA after 8 hours integration /GHz Bmax b b .140pc/ Vkep .1Mˇ / 115 150 m 3.600 500 AU 1.3 km/s 1500 m 0.3600 50 AU 4.2 km/s 15 km 0.03600 5 AU 13 km/s

T (1km/s) 4.6 mK 0.46 K 46 K

230

150 m 1500 m 15 km

1.800 0.1800 0.01800

250 AU 25 AU 2.5 AU

1.9km/s 6km/s 19km/s

6 mK 0.6 K 60 K

345

150 m 1500 m 15 km

1.200 0.1200 0.01200

170 AU 17 AU 1.7 AU

2.3km/s 7km/s 23km/s

9 mK 0.9 K 90 K

460

150 m 1500 m 15 km

0.900 0.0900 0.00900

130 AU 13 AU 1.3 AU

2.6km/s 8km/s 26km/s

32 mK 3K 300 K

690

150 m 1500 m 15 km

0.600 0.0600 0.00600

85 AU 8.5 AU 0.85 AU

3.3km/s 10km/s 33km/s

21 mK 2.1 K 200 K

in mind that there are at least 36 possible configurations in total allowing fine tuning of the trade off between angular resolution and sensitivity. For each frequency, the beam size is shown in angular units (column 3) and also in physical scale at 140 pc where the nearest low-mass stars are forming (column 4). In addition, column 5 shows the Keplerian rotation speed at one beamwidth from the central star, assuming a mass of 1 Mˇ . The final column shows the 1- brightness sensitivity in 8 hours on source, in a 1 kms1 wide spectral channel. ALMA is currently under construction. As I write, more than ten of the 12-m antennas have delivered to the site in Chile, and are currently being assembled and tested. The first 20 concrete antenna pads have been laid, and the first receiver system has been installed in an antenna. Scientific commissioning should begin in 2009, with the first fringes expected late that year. First open science opportunities with ALMA, using a restricted number of antennas, are anticipated for late 2011/early 2012, with the completion of the 66-antenna array planned for late 2013.

3 Potential Outflow and Jet Studies with ALMA ALMA will revolutionise our studies of protostellar jets, disks and outflows. Its large collecting area, sensitive receiver systems and very dry site give it up to 100 times the sensitivity of existing millimetre/submillimetre arrays. What can we do with this increased sensitivity? Essentially, it buys us resolution: in thermal brightness terms, an array’s sensitivity scales at the inverse square of the maximum baseline in the array. So this means that the two orders of magnitude increased sensitivity allow us to make observation width about 10 times better angular resolution than before.

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ALMA will therefore enable routine imaging at the 0.1 arcsec level. The table (Table 1) shows some of the capabilities. Consider for example the 12 CO J D 3  2 line at 345 GHz, which is an excellent tracer of warm outflowing molecular gas. In compact configuration, the angular resolution is 1.2 arcsec, and the sensitivity is an exquisite 0:009 K in a 1 kms1 wide spectral channel in 8 hours integration. This configuration will be useful for studying the large-scale structure of outflows: one is likely to integrate for much shorter times than 8 hours to image the high-velocity CO, and make mosaics to map out the extended outflow structures at good resolution. As we push out to a longer array, with a maximum baseline of 1500 m, the resolution increases by a factor of 10 to 0.12 arcsec, and the the sensitivity falls by a factor of 100 to 0.9 K. Such an observation would give a spectacular image of the high velocity molecular gas from a protostar, with very good velocity and spatial resolution. Pushing out to the largest 15-km array size, we would now achieve 0.012 arcsec resolution, corresponding to only 1.7 AU in a protostar in Taurus, although the sensitivity is now only 90 K. In such an observation, only the very hottest gas, presumably which has just been shocked, will be detectable, and it may be necessary to smooth spectral channels (or integrate for longer) to achieve higher signal-to-noise ratios: for example, by binning to 16 kms1 resolution the sensitivity will then be 22 K, sufficient to detect optically-thick 100 K molecular gas. Given that at this spatial scale the Keplerian speed is 23 kms1 , even this coarse velocity resolution should be enough to trace the jet dynamics close to the launch point. Pushing up to even higher frequencies, the table shows it is even possible to achieve sub-AU resolution using the CO 6-5 line at 690 GHz: such images would give unique insight into the velocity field very close to where the jet is launched and collimated.

3.1 Molecular Outflow Surveys with ALMA To study outflow propagation and the global properties of outflows, ALMA in its compact and intermediate configurations will be a powerful tool. For example, a recent survey of embedded YSOs (Jorgensen et al., 2008) used the SMA to map flows at  200 resolution. With ALMA, we could repeat and extend such surveys at 10 times better angular resolution i.e. 0.2 arcsec resolution, achieving T D 0:6 K in 8 hours using the CO 2-1 line (Table 1). Alternatively, one could observe a large sample of outflows in snapshot mode at low angular resolution: for example, in the compact configuration at 230 GHz one can achieve T D 0:1 K in only 4 minutes on source. Such surveys have the capability of providing a wealth of high-resolution data using mapping out the outflow velocity field in detail. In addition, by using the wide ALMA bandwidth and powerful correlator configurations available, it is possible to map simultaneously the velocity field of the disk and infalling material, allowing us to study directly the accretion/outflow link at high resolution.

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ALMA’s high angular resolution will permit many scientific issues to be addressed:  Multiplicity and clustering: because most stars form in clusters, and a good frac-

tion form multiple star systems, it is often difficult to trace flows back to their driving sources. With ALMA’s very high angular resolution, it should be possible in many more cases to disentangle outflow sources and resolve confused cluster forming regions.  Resolving distant high-mass outflows: high-mass molecular outflows are typically observed at low physical resolution because of their distances. ALMA will finally allow us to study such outflows in as much details as nearby low-mass outflows, and permit the study of outflows across the galaxy.  Testing scaling relations: the dependence of outflow momentum and energy on the YSO properties, such as mass and luminosity, can be studies with much less uncertainty because the individual outflows can be cleanly separated at high angular resolution.

3.2 Jet and Outflow Proper Motions with ALMA ALMA will enable for the first time routine studies of the proper motions in molecular gas. Such studies have become very useful in optical/IR studies, measuring the pattern speed in shocks and helping decipher the 3-dimensional geometry of outflows (see e.g. Bally’s contribution to this volume). With ALMA, we will be able to do the same using molecular line diagnostics. The angular proper motion expected is given by 1    t  v D :  D 0:200 100pc 100kms1 1year Because ALMA is an interferometer, it has excellent astrometric capabilities, so that the signal-to-noise ratio (SNR) will be the main limitation in measuring accurate proper motions. Assuming one can measure the location of a compact feature to an accuracy of b =SNR, the equation shows that ALMA should be able to detect transverse motions of 75 kms1 using observations separated by 1 year even assuming a modest SNR of 10 and a beamsize of 0.200 . Observations spaced over 5–10 years will show significant evolution of the outflow and place strong constraints on the outflow geometry and jet propagation. In addition, it is worth noting that assuming the disks are not azimuthally smooth, we can measure the disk proper motions as well with ALMA on similar timescales.

3.3 Jet/Outflow Shocks and Chemistry It is well-known that when molecular gas is shocked, very complex physical and chemical structures arise. ALMA will greatly advance such studies but only if we have advanced theoretical models of shocks including detailed chemistry with which

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to compare the data. In particular ALMA will produce very high resolution images of detailed shock physics, chemical enhancements in bowshocks, and is able to resolve typical molecular cooling lengths. In addition, there is the possibility of resolving the expected ion-neutral differential velocity predicted by C-shock models. The great advantage ALMA offers is high spatial resolution plus broad bandwidths allowing several molecular species to be studied. One particular issue that should be addressed now is what are the best tracers of jets and outflows? Although CO is the most ubiquitous tracer, the gas-phase emission of SiO, believed to arise from grain sputtering, is potentially the best diagnostic of young jets. The question remains: is the SiO tracing the jet itself through in situ formation, or does it merely trace shocks close to the jet where material is entrained? ALMA should finally resolve this issue. Perhaps atomic carbon, with its 492 and 809 GHz transitions, will finally be detected in jets, providing a potential new tracer of atomic jet material.

3.4 Calorimetry: Measuring the True Power of Outflows Although the energetics of outflows have been the subject of numerous studies, huge uncertainties remain in the measured outflow energies and jet powers (e.g. [1]). Such data are critical to any meaningful test of outflow generation theories. Detailed studies of a well-defined sample of outflows with ALMA should significantly reduce there observational uncertainties for the following reasons:  ALMA offers good coverage of CO rotational lines, including J D1–0, 2–1,

3–2, 4–3, 6–5, and 8–7. The last of these lines lies 200 K above ground state, so that multi-frequency imaging should allow accurate temperature estimates of the outflowing gas which are essential for accurate mass estimates.  ALMA has sufficient sensitivity to study the CO isotopologues such as 13 CO and C18 O in outflowing gas, which allows accurate optical depth estimation.  High-resolution data allow us to distinguish separate flows in regions where star clusters are forming.  As discussed above, disk observations with ALMA will allow accurate disk inclination angles to be derived, assuming the disk is orthogonal to the jet; in addition, the jet proper motions can also help establish the sin.i / correction. Finally, the protostellar mass can be inferred from the disk rotation speed. The relative abundance of CO in the outflowing gas will still need to be established somehow; nonetheless, we expect to be able to reduce significantly the scatter in the jet power-protostellar mass correlations, which provide strong constraints on disk and X-wind models.

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3.5 Towards the Central Engine: Launch, Collimation and Rotation We know from optical/IR studies that jet launch occurs on scales below 0.1 arcsec, i.e. 14 AU. With molecular diagnostics at millimetre and submillimetre wavelengths, the current best images have about 0.3–0.400 resolution. However, ALMA’s ability to image down to 0.01 arcsec, albeit at low surface brightness sensitivity, will allow us to probe down to AU scales in the disk. It is possible that this tenfold increase in angular resolution will finally enable us to see the jet launch point from the disk, and the scale on which it is collimated. However, this depends directly on the correct jet launching model (see the papers by Ferreira and Cai in this volume). In X-winds, the launch point is very close to the star, only a few stellar radii at most, and even ALMA cannot resolve this scale. However, extreme variants of the disk-wind models can in principle launch material from scales close to an AU and this could be directly tested by ALMA. Detecting the signature of jet rotation is of great interest as it provides a direct test of jet generation models. Several high resolution optical/IR spectroscopic studies (e.g. [5]) have detected evidence in the atomic jet for rotation, but the results are all unambiguous due to marginal angular resolution. ALMA’s ability to image down to 10 mas resolution should really help settle this debate. In addition, ALMA will measure cleanly the disk rotation, so a direct check can be made that any detected rotation see in the jet is aligned with the disk rotation. The main caveat is that millimetre and submillimetre emission lines from material present in the jet must be found. High-J CO lines, neutral carbon at 492 and 809 GHz, and SiO and water are all possible tracers depending on the chemistry in the jet/wind.

3.6 Probing the Magnetic Field in Outflow Sources The importance of magnetic fields in accretion disks and jets has long been recognised. Nearly all theoretical models rely on the magnetic fields in the disk or star to help accelerate outflowing material. But our ability to measure magnetic field structures and strengths is severely limited. ALMA will make significant advances to the state of the art in three areas: 1. Polarised dust continuum. We know that on protostellar envelope scales, polarised continuum flux can be used to trace magnetic field geometry, under the assumption that the emitting dust grains are rapidly aligned with their major axes orthogonal to the field direction. ALMA will for the first time enable such studies down to arcsec scales and below, allowing us to map out the projected magnetic field geometry on 100 AU scales. This will lead to strong tests of the models of magnetically-mediated star formation, and may provide clues as to the role of magnetism in the launch of jets. However, such observations do not lead to good magnetic field strength estimates, so there is usually a major ambiguity in interpreting the models.

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2. Linearly polarised emission lines: in the Goldreich-Kylafis effect [2,3], rotational emission lines can become linearly polarised in the presence of a magnetic field. ALMA can measure the line polarisations from the correlations of the crosspolar data, and hence should allow us to measure this effect in many outflow sources, including in the outflowing gas itself. However, the predicted patterns are highly model-dependent, with the magnitude and orientation of the effect varying sensitively with source model and optical depth, making such data hard to interpret. Better source models and radiative transfer tools would certainly help us plan for such observations. 3. Zeemann splitting: ALMA’s linearly polarised detectors make it well suited to precision measurements of circular polarisation, and its high sensitivity may help enable more Zeemann splitting detections in disk and outflow sources. Zeemann coefficients are very low in the submillimetre, making the observations challenging, but radicals such as CN offer promising emission lines in which to search for the effect.

4 Conclusions Protostellar jets are ideal ALMA science targets for high resolution imaging: there is exciting physics happening on very small angular scales, and we expect the gas emitting in those regions to be warm and compact, heated by shocks and/or the central protostar. ALMA will revolutionise our understanding of mass ejection from protostars. It has the ability to image molecular (and a few atomic) lines from gas deep into the potential well of protostars, to resolve collimation scale, search for jet rotation and study the disk/wind interface. In addition, ALMA has the ability to map disk, envelope and jet properties simultaneously using suitably chosen molecular emission lines, as well as the warm dust continuum emission, with very high velocity resolution. ALMA will provide full 3-dimensional data cubes, which will be much easier to model and interpret than the long slit spectra that typically are found in optical/IR data. The wide range of molecular lines available to ALMA, and its wide range of configurations, allows us to fine-tune observational programmes to pick out the jet physics on a precise range of physical and temperature scales, and constrain accurately the gas temperature and densities. Finally, ALMA’s excellent polarisation capabilities will allow us to make the most detailed images of magnetic field geometry on small angular scales.

References 1. Bontemps, S., Andre, P., Terebey, S., Cabrit, S. 1996. Evolution of outflow activity around lowmass embedded young stellar objects. Astronomy and Astrophysics 311, 858–872. 2. Goldreich, P., Kylafis, N. D. 1982. Linear polarization of radio frequency lines in molecular clouds and circumstellar envelopes. Astrophysical Journal 253, 606–621.

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3. Greaves, J. S., Holland, W. S., Ward-Thompson, D. 2001. Measurement of the Magnetic Field Direction in the NGC 2024 FIR 5 Protostellar Outflow. Astrophysical Journal 546, L53–L56. 4. Jørgensen, J. K., Bourke, T. L., Myers, P. C., Francesco, J., van Dishoeck, E. F., Lee, C.-F., Ohashi, N., Sch¨oier, F. L., Takakuwa, S., Wilner, D. J., Zhang, Q. 2007. PROSAC: A Submillimeter Array Survey of Low-Mass Protostars. I. Overview of Program: Envelopes, Disks, Outflows, and Hot Cores, Astrophysical Journal 659, 479–498. 5. Pesenti, N., Dougados, C., Cabrit, S., Ferreira, J., Casse, F., Garcia, P., O’Brien, D. 2004. Predicted rotation signatures in MHD disc winds and comparison to DG Tau observations. Astronomy and Astrophysics 416, L9–L12. 6. Richer, J. S., Shepherd, D. S., Cabrit, S., Bachiller, R., Churchwell, E. 2000. Molecular Outflows from Young Stellar Objects, In Protostars and Planets IV, p. 867. 7. Richer, J. 2005. Scientific requirements of ALMA, and its capabilities for key-projects: Galactic. ESA Special Publication 577, 33–38. 8. Shepherd, D. S. 2008. Molecular outflows observed with ALMA. Astrophysics and Space Science 313, 41–44. 9. Tarenghi, M., Wilson, T. L. 2005. The ALMA Project. EAS Publications Series 15, 423–430.

Two-component Jet Simulations: Combining Analytical and Numerical Approaches Titos Matsakos, Silvano Massaglia, Edo Trussoni, Kanaris Tsinganos, Nektarios Vlahakis, Christophe Sauty, and Andrea Mignone

Abstract Recent observations as well as theoretical studies of YSO jets suggest the presence of two steady components: a disk wind type outflow needed to explain the observed high mass loss rates and a stellar wind type outflow probably accounting for the observed stellar spin down. In this framework, we construct numerical twocomponent jet models by properly mixing an analytical disk wind solution with a complementary analytically derived stellar outflow. Their combination is controlled

T. Matsakos (), S. Massaglia, and A. Mignone DFG, University of Turin, via P. Giuria 1, 10125 Torino, Italy e-mail: [email protected]; [email protected]; [email protected] E. Trussoni INAF/Osservatorio Astronomico di Torino, via Osservatorio 20, 10025 Pino Torinese, Italy e-mail: [email protected] K. Tsinganos and N. Vlahakis IASA and Section of Astrophysics, Astronomy and Mechanics, Department of Physics, University of Athens, Panepistimiopolis, 15784 Zografos, Athens, Greece e-mail: [email protected]; [email protected] C. Sauty Observatoire de Paris, L.U.Th., 92190 Meudon, France e-mail: [email protected]; [email protected]

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by both spatial and temporal parameters, in order to address different physical conditions and time variable features. We study the temporal evolution and the interaction of the two jet components on both small and large scales. The simulations reach steady state configurations close to the initial solutions. Although time variability is not found to considerably affect the dynamics, flow fluctuations generate shocks, whose large scale structures have a strong resemblance to observed YSO jet knots.

1 Introduction In the last few years, a promising two-component jet scenario seems to emerge in order to explain Young Stellar Object (YSO) outflows. Observational data of Classical T Tauri Stars (CTTS) [1, 3] indicate the presence of two genres of winds: one being ejected radially with respect to the central object and the other being launched at a constant angle with respect to the equatorial plane (e.g. Tzeferacos et al., this volume). In turn, CTTS outflows may be associated with either a stellar origin, or a disk one or with both wind components having roughly equivalent contributions. In addition, such a scenario is supported by theoretical arguments as well (e.g. [2]). An extended disk wind is required for the explanation of the observed YSO mass loss rates, whereas a pressure driven stellar outflow is expected to propagate in the central region, being a strong candidate to address the protostellar spin down [6]. The goal of the present work is to study the two-component jet scenario, taking advantage of both analytical and numerical approaches. In particular, we construct numerical models by setting as initial conditions a mixture of two analytical YSO outflow solutions (each one describing a disk or a stellar jet), ensuring the dominance of the stellar component in the inner regions and of the disk wind in the outer. The combination is achieved with the introduction of few normalization and mixing parameters, along with enforced time variability of the stellar component. We investigate the evolutionary properties, steady states and the features of the final configurations of the dual component jets. Although the detailed launching mechanisms of each component are not taken into account, such models seem capable to capture the dynamics and describe a variety of interesting scenarios. The employed analytically derived MHD outflows, defined as ADO (Analytical Disk Outflow; denoted with subscript D) and ASO (Analytical Stellar Outflow; denoted with subscript S), have been derived in the context of self-similarity [9] and each one effectively describes a disk wind [10] or a stellar jet [8], respectively. In Matsakos et al. [4], we have addressed the topological stability, as well as several physical and numerical properties, separately for each solution. This article summarizes the numerical setup and reports the results of few significant cases of the dual component jet. A thorough study can be found in Matsakos et al. [5].

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2 Numerical Two-component Jet Models The two-component jet model parameters can be classified in two categories. The first one contains those associated to the relative normalization of the analytical solutions, i.e. the ratios of the characteristic scales of each model (denoted with subscript *, calculated on a specific fieldline at the Alfv´enic surface): `L D

R ; r

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where R is the Alfv´enic spherical radius of the ASO model, r is the cylindrical radius of the Alfv´enic surface of the ADO model (of a specific fieldline) and the subscripts L, V and B stand for length, velocity and magnetic field, respectively. We assume that the protostar has a solar mass and a radius of 0.01 AU. Since the disk wind launching region lies in the range 0.2–3 AU, we derive `L D 0:1 and `V D 5:96. On the other hand, we define `B D 2, which is the parameter controlling the relative dominance of each model. The second class of parameters concerns the mixing. In particular, we choose the combination to depend on the magnetic flux function A, which labels the fieldlines of each analytical model (AD or AS ). Therefore, we define a common trial magnetic flux At r D AD C AS and then all physical variables U are initialized with the help of the following mixing function: (

U2comp

"  "   #)  # At r d At r d D 1  exp  UD C exp  US ; qAm qAm

(2)

where Am is a constant corresponding to the matching surface rooted at 0.16 AU, q is a parameter that effectively moves this surface closer to the protostar and d sets the steepness of the transition from the inner ASO to the outer ADO solution. We choose q D 0:2 and d D 2, whereas a complete parameter study (including `B ) can be found in [5]. Essentially, (2) provides an exponential damping of each solution around a particular fieldline of the combined magnetic field. Moreover, since accretion and protostellar variability are expected to introduce fluctuations we multiply the inflow velocity with the following function: "     # 2 t r 2 1 exp  ; fS .r; t / D 1 C sin 2 Tvar 2rm

(3)

where Tvar is the period of the pulsation and rm is roughly the cylindrical radius at which the matching surface intersects the lower boundary of the computational box. Outflow variability produces the formation of knot-like structures: the introduction of radiation cooling (Tesileanu et al. this volume) will allow direct comparison with observational data.

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The simulations are performed with PLUTO1 [7]. A uniform resolution of 256 zones for every 100 AU is used whereas the simulations have been carried out up to a final time of 80 y. On the lower boundary we keep fixed all variables to their initial values, on the axis we apply axisymmetric boundary conditions and at the upper and right borders of the domain we prescribe outflow conditions.

3 Results In the left panel of Fig. 1, the initial setup (left) and the final configuration (middle) of the two-component jet are displayed. The model shows remarkable stability and reaches a well defined steady state in only a few years. In particular, the disk wind solution remains almost unmodified whereas the stellar component gets compressed around the axis. Moreover, a shock manifests during time evolution, located roughly along the diagonal line which crosses (10, 40) and (30, 100) (steady state plot). This shock is found to causally disconnect the acceleration regions from the jet propagation physics and the subsequent interaction with the outer medium. Note that there is no such “horizon” present in the initial setup. Furthermore, on the right plot of the left panel of Fig. 1, the same model is displayed when a monthly time variable

Fig. 1 Left panel: Logarithmic density contours (the code unit is 1012 g cm3 ) and magnetic fieldlines for the initial two-component jet model (left), final steady state (middle) and when a p monthly flow variability is applied (right). Right panel: The quantity 103 2 Tp(roughly related with emissivity) is plotted for the yearly variable model. Although max.103 2 T / D 53:9, the color bar uses a lower maximum value to enhance the displayed features

1 A versatile shock-capturing numerical code suitable for the solution of high-Mach number flows. Publicly available at http://plutocode.to.astro.it

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velocity is applied on the lower boundary (Eq. 2). Evidently, despite the strong gradients seen in the density and the wiggling of the magnetic fieldlines, the general structure is retained, proving the stability of the two-component jet model. The fact that the system remains very close to the initial configuration, demonstrates that the analytical solutions provide solid foundations for realistic twocomponent jet scenarios. Consequently, specific YSO systems can be addressed more accurately by constructing analytical outflow solutions with the desirable characteristics, before merging them into a two-component regime. On the right panel of Fig. 1 a quantity related to emissivity is plotted in larger scales when a yearly variability is applied. Near the base, the numerical solution remains close to the initial ADO and ASO models. However, farther away along the flow the fluctuations create knot-like structures, which may be related with jet variability. In fact, note that the model is associated with a condensations spacing 100 AU, similar to the knot spacing of HH30. Finally, although not presented in this article, an other important parameter is the one controlling the relative contribution of each component, `B , with which we can effectively and smoothly switch the model from a totally magneto-centrifugal wind to a pressure driven jet [5].

4 Conclusions To sum up (taking also into account the results of [4] and [5]), most of the technical part concerning two-component jets, e.g. 2.5D stability, steady states, parameter study, time variability etc., is now at hand, providing us with all the necessary ingredients to address YSO jets. Namely, with (a) the proper analytical solutions, i.e. desirable lever arm, mass loss rate etc., (b) the correct choice of the mixing parameters and (c) an enforced time variability that effectively produces knot structures, we are now ready to qualitatively study different and realistic scenarios, address observed jet properties and ultimately understand the various outflow phases of specific T Tauri stars. Acknowledgements The present work was supported in part by the European Community’s Marie Curie Actions - Human Resource and Mobility within the JETSET (Jet Simulations, Experiments and Theory) network under contract MRTN-CT-2004 005592 and in part by the HPC-EUROPA++ project (project number: 211437), with the support of the European Community - Research Infrastructure Action of the FP7 Coordination and support action Programme.

References 1. Edwards, S., Fischer, W., Hillenbrand, L., & Kwan, J., ApJ, 646, 319–341, (2006) 2. Ferreira, J., Dougados, C., & Cabrit, S., A&A, 453, 785–796, (2006) 3. Kwan, J., Edwards, S., & Fischer, W., ApJ, 657, 897–915, (2007)

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4. Matsakos, T., Tsinganos, K., Vlahakis, N., Massaglia, S., Mignone, A., & Trussoni, E., A&A, 477, 521–533, (2008) 5. Matsakos, T., Tsinganos, K., Vlahakis, N., Massaglia, S., Trussoni, E., Sauty, C. & Mignone, A., submitted to A&A 6. Matt, S., & Pudritz, R., ApJ, 678, 1109–1118, (2008) 7. Mignone, A., Bodo, G., Massaglia, S., Matsakos, T., Tesileanu, O., Zanni, C., & Ferrari, A., ApJS, 170, 228–242, (2007) 8. Sauty, C., Trussoni, E., & Tsinganos, K., A&A, 389, 1068–1085, (2002) 9. Vlahakis, N., & Tsinganos, K., MNRAS, 298, 777–789, (1998) 10. Vlahakis, N., Tsinganos, K., Sauty, C., & Trussoni, E., MNRAS, 318, 417–428, (2000)

Jets from Young Stellar Objects: From MHD Simulations to Synthetic Observations Ovidiu Tes¸ileanu, Andrea Mignone, and Silvano Massaglia

Abstract With the recent improvements in available observational data, simulating the radiative processes in YSO jets will provide a valuable tool for model discrimination. We have added a radiative cooling module and time-dependent ionization state computation for 29 ion species to our MHD code PLUTO. Also, post-processing routines for the realistic computation of emission lines are now available. From 2D simulations, synthetic surface brightness maps and positionvelocity diagrams for the line emissions can be directly computed, to be compared with observations.

O. Tes¸ileanu () Universit`a degli Studi di Torino, via P. Giuria 1, I-10125 Turin, Italy and Research Centre for Atomic Physics and Astrophysics, RO-077125, Bucharest, Romania e-mail: [email protected] A. Mignone Universit`a degli Studi di Torino, via P. Giuria 1, I-10125 Turin, Italy Osservatorio Astronomico di Torino, via Osservatorio 20, I-10025 Pino Torinese (TO), Italy S. Massaglia Universit`a degli Studi di Torino, via P. Giuria 1, I-10125 Turin, Italy

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1 Introduction Herbig-Haro (HH) objects are regions of shocked gas in collimated supersonic jets which emerge along the rotation axis of accretion disks around forming stars. The Hubble Space Telescope provided new high-resolution data by resolving the subarcsecond scale structure and cooling layers of dozens of HH objects. There are several radiative numerical codes employed in astrophysics, with var` [5], or ious implementations of cooling and ionization treatments (e.g., YGUAZU ASTROBEAR [3]). The MHD simulation code we use for our astrophysical applications – PLUTO – is a freely distributed application developed and maintained at the Turin University – Turin Astronomical Observatory [2]. In this work, the new cooling function MINEq (Multi-Ion Non-Equilibrium) embedded in PLUTO is used in order to generate, from numerical simulations, maps directly comparable with observations. The main advantage of the MINEq approach is the full ionization state computation during the MHD simulation, which allows for better predictions of emission line intensities. The procedures of converting the MHD simulation data to synthetic intensity maps, PV diagrams and synthetic spectra will be presented. Because of the various scales of the ionization/recombination, cooling and dynamical processes, resolutions of 1011 cm are needed for the postshock zone on grids with total dimensions of the order of 1016 cm, forcing the use of Adaptive Mesh Refinement (AMR) techniques.

2 The Cooling Module The newly developed cooling module for the PLUTO MHD code integrates a complex ionization network of 29 ion species: H I, H II, He I, He II, C I to V, N I to V, O I to V, Ne I to V, and S I to V, and the collisionally excited line emission for these ion species in the approximation of a 5-level atom. The total energy E is evolved according to the standard MHD equations: @E C r Œ.E C pt / v  .v B/ B D SE ; @t

(1)

where SE is a radiative loss term, and pt  p C jBj2 =2 denotes the total pressure (thermal C magnetic) of the fluid. For each ion, we solve the additional equation @. X ;i / C r . X ;i v/ D S ;i @t

(2)

coupled to the original system of conservation laws. In (2), the first index ( ) corresponds to the element, while the second index (i ) corresponds to the ionization stage. Specifically, X ;i  N ;i =N is the ion number fraction, N ;i is the number density of the i -th ion of element , and N is the element number density. We denote the whole set of ions for all possible and i with X  fX ;i g.

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The source term S ;i accounts for ionization and recombination. The total line emission from these species enters in the source term SE in (1) and should give a good approximation of radiative cooling for the above conditions [4]:   SE D  Nat Nel  .T; X/ C LFF C LIR ;

(3)

where .T; X/ is the radiative cooling function due to collisionally-excited line radiation, LFF denotes the free-free (bremsstrahlung) losses from HC and HeC , while LIR accounts for the energy lost during ionization/recombination processes. The number densities Nat and Nel are, respectively, the total atom and electron number densities, readily determined from the mass density and the known chemical composition of the plasma. For a detailed description of the cooling function, testing and sample applications to astrophysics, we refer to [6].

3 Synthetic Observations The current efforts concentrate on the computation of “synthetic observations” from the MHD simulations, that is emission maps and emission line ratios to be directly compared with the available observations. This Chapter will describe the method and algorithms of the passage from the MHD simulation (left and middle panels of Fig. 1) results to synthetic observations.

3.1 Emission Maps The input data from PLUTO is the result of a steady jet simulation with an internal perturbation evolving into a shock (Fig. 1, left) after 100 yrs of propagation

Fig. 1 Left panel: Logarithmic map of density of the whole computational domain of the AMR jet simulation with the perturbation; Middle panel: Area of the propagating shock, density logarithmic map with refinement levels superimposed; Right panel: Emissivity logarithmic map in ˚ units of erg cm3 s1 [S II]6731 A,

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along the jet. The simulation was performed employing Chombo AMR with 7 levels of refinement, at a very high equivalent resolution (6144  24576 cells) on a cylindrical grid of 10  40 jet radii (rj D 2:5 1015 cm). At this resolution, the ionization/recombination and cooling scales for a moderate shock as the one presented should be resolved. From the MHD data we will enlarge and analyze the area containing the propagating shock (middle panel in Fig. 1). The jet setup is the following: jet speed, vjet D 150 km s1 ; ambient medium density, na D 103 cm3 ; jet density, njet D 104 cm3 ; velocity perturbation producing the shock, v D 0:25 vj . The perturbation is created using a special technique explained in [1], in order to produce a single forward shock instead of the usual forward-reverse pair. After computing the 2D emissivity maps in the desired emission lines (e.g., right panel of Fig. 1), each cell is cylindrically integrated in 3D (the emissivity is multiplied with the cell volume), with a defined inclination with respect to the line of sight, and then projected on a plane perpendicular on the line of sight. The angular step around the z axis of the cylindrical integration can be customized. An user-defined Gaussian point source function (PSF) is applied to the resulting map to simulate the effect of the observing instrument. In the presented map in the left panel of Fig. 2 the variance of the Gaussian is  D 0:1 in simulation lengths units on both axis and the jet axis forms an angle of 80ı with the line of sight.

3.2 Position-velocity Diagrams For the Position-velocity (PV) diagram presented in the right panel of Fig. 2 the slit was defined on the axis of the jet, with a length 1:0 in simulation units starting at z D 23:125 and a width of 0:1 starting at r D 0:05. The slit comprises all the relevant shock and post-shock zones along the z axis. The ranges and resolution of the PV diagram are user-defined on the velocity axis and set up by the slit dimensions and emission map resolution on the position axis. In the presented examples, the PV diagram dimensions are 206 cells for the position

˚ line: Source surface brightness in units of erg cm2 s1 arcsec2 , with the Fig. 2 [S II]6731 A slit location superimposed (left panel); PV diagram for the defined slit, same units with velocity bins of 1 km s1 (right panel)

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Fig. 3 Emission spectra, logarithmic scale in units of erg cm2 s1 arcsec2 Hz1 : (from left to right) before the shock, immediately behind the shock, farther behind the shock

axis and 100 cells for velocities (radial velocities from 50 to 50 km=s with 1 km=s bins). The PV diagrams are convolved with a smoothing Gaussian function that has the  D 0:1 on the position axis and  D 5 km=s on the velocity axis.

3.3 Synthetic Spectra It is possible to create a synthetic spectrum, as it would be observed from a defined slit. For each of the points along the slit (at the simulation resolution) a spectrum is generated. The Doppler shifts due to the radial velocity are computed, as well as the thermal line broadening. ˚ at three positions along the slit are preIn Fig. 3 the spectra from 6200 to 6800 A ˚ sented. The emission lines computed and visible in this region are: H˛ 6562.8 A, ˚ [OI] 6300C6363 A, ˚ [SII] 6716C6731 A. ˚ The lines are very [NII] 6548C6583 A, weak ahead of the shock front, as expected, and sharply increase after. It is worthwhile to note that the highest emission from the [NII] doublet occurs after the [OI] and [SII] lines peaked – which is normal considering the slower ionization of N with respect to S and the fact that OI is a neutral specie.

4 Conclusions The newly-developed cooling function provides a powerful tool for investigating the stellar jets and gaseous nebulae. It is, in the current configuration, suited for the study of radiative shocks in stellar jets. The synthetic observations have finally become feasible with the latest developments in available computer power, and will be the tool for jet model discrimination with the increasing resolution of observations. Acknowledgements The present work has been supported by the European Union (contract MRTN-CT-2004-005592) within the Marie Curie RTN JETSET.

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References 1. Massaglia, S., Mignone, A., & Bodo, G. 2005, A&A, 442, 549 2. Mignone, A., Massaglia, S., Bodo, G., et al. 2007, ApJS, 170, 228 3. Poludnenko, A., Varni`ere, P., Cunningham, A., Frank, A., Mitran, S. 2005, Lecture Notes in Computational Science and Engineering 41, pp 331–340, Springer 4. Raga A.C., Mellema G., Lundqvist P. 1997, ApJS, 109, 517 5. Raga, A.C., Navarro-Gonz`alez, R., Villagr`an-Muniz, M. 2000, Rev.Mex.AA, 36, 67 6. Tes¸ileanu, O., Mignone, A., Massaglia, S. 2008, A&A, 488, 429

Molecular Cooling in Large Scale Simulations of Protostellar Jets Jamie O’Sullivan and Max Camenzind

Abstract We discuss the case for J-shocks in large scale protostellar outflows as sources of the observed infrared molecular emission and describe our model for simulating such jets using the PLUTO code with an additional molecular chemistry and cooling module. Results are presented from initial gas-chemistry jet simulations showing the effect of the large scale flow on the chemical properties, and we describe our ongoing work on improving the model via comparison with more detailed models.

J. O’Sullivan () Landessternwarte (ZAH), Heidelberg, Germany e-mail: [email protected] M. Camenzind Landessternwarte (ZAH), Heidelberg, Germany e-mail: [email protected]

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1 Introduction and Aim The origins of the molecular emissions associated with protostellar jets are the subject of ongoing investigation in the numerical modelling community. Several protostellar jets are seen to have extensive emission in the infrared wavelengths, particularly in H2 ro-vibrational lines, giving evidence to the presence of shockheated molecular material. A notable example is HH46/47, seen in Fig. 1, which shows a distinct “cocoon” of molecular emission. There have been a number of suggestions to explain the physical origin of the molecular outflow, as described in the contribution of Downes in these proceedings. In the case of the jet-driven outflows, where the molecular emission is supposed to come from gas shocked by the jet’s bow-shock, there is still the matter of determining whether the emission is from surviving or reformed molecules in the post-shock region of a J-type shock, or from the milder conditions of C-shocks, where molecules are less prone to being dissociated. In some sources, such as HH211, there is some evidence to indicate that the emission comes from molecular gas in C-type bow-shocks (see [10]). However, given the variety of shock speeds and the difficulty of ascertaining the strength of magnetic fields in most sources, the presence of conditions conducive to the production of C-shocks is usually not guaranteed. Therefore, can J-shocks in large scale jets account for the cocoon-like emissions from molecular hydrogen seen in some sources? In current modelling approaches, it is becoming more commonplace to use hydro- or magnetohydrodynamical codes in conjunction with some non-ideal or micro-physical processes to simulate astrophysical jets. Although multi-fluid phenomena such as the ambi-polar diffusion needed for C-shocks are out of reach for single fluid codes, we can attack the problem by using large scale simulations

24μm Hotspot

Outflow NE HH 47A

Outflow SW Bow shock Cavity

Star jet

24μm Hotspot

Compact jet

jet?

HH47C

Fig. 1 Enhanced Spitzer image from [11] showing the infrared emissions, which coincide with the molecular H2 emissions earlier detected by e.g. [4]. The prevalent SW cocoon-like emission (bottom-right) is the result of the jet impinging on the ambient gas of the molecular core containing the protostar

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with chemistry to investigate the feasibility of J-shocks as an explanation for the molecular emission, as well studying the broader effect of molecular cooling on the properties of the jets in general.

2 Method The engine of our code is the finite volume magnetohydrodynamical code PLUTO (see [9] for a detailed description). In addition to the hydrodynamic solver used to simulate the jet gas, we have implemented a chemical network to simulate the gas chemistry and calculate the amount of molecular cooling. Our present metal-free gas chemistry model accounts for the dominant gas-phase pathways for the production of molecular hydrogen (see for example [1]). The model solves the species concentrations for HI, e , HC , H , HC 2 and H2 , with 15 reactions between these species. The rate equations for the chemistry network have the general form @ni D Ci .T; nj /  Di .T; nj /ni @t

(1)

and we use a quasi-implicit BDF scheme of the form below to cope with the stiffness of the system, CitCt t C nti ntCt D (2) i 1 C DitCt t where the pseudo-implicit values for Ci (T,nj ) and Di (T,nj ) are obtained by using the updated species fractions as they are calculated. This scheme is highly robust and stable, albeit at the expense of some accuracy [2]. We have heretofore not included any treatment of dust, and there may be some justification for this in the temperature and density regime considered (see [7]). However, it is planned to include dust at a later stage in order to investigate this assertion. The cooling losses are calculated from the species abundances. Cooling due to atomic hydrogen (i.e. collisional ionisation, collisional excitation, and radiative recombination) is treated as in [3], and the ro-vibrational losses due to molecular hydrogen were implemented following [6]. These losses are incorporated with the hydrodynamics via a split-source step.

3 Simulation Results So far we have carried out intermediate scale simulations, modelling the jets as highly supersonic (Mach 12) inflows of gas with number density 104 cm3 at a temperature of 104 K into a medium with a number density and temperature both lower by a factor of 10. The injected jet beam was given an ionisation of 10% and the ambient medium was given a molecular hydrogen mass fraction of 10%.

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The simulations were carried out on a domain of 400  2000 points in axisymmetric cylindrical coordinates. The physical domain corresponds to a distance of 1600 AU, which is covered by the jet in some 170 years. Shown in Fig. 2 are the physical variables, which show the characteristic effects of the cooling on the morphology. The ionisation is seen to decrease along the beam of the jet as recombination occurs. Figure 3 shows the chemical species fractions. The abundances of H and HC 2 , though strongly dependent on the temperature, remain low rho (x 1e4 cm–3) 2.0

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throughout the domain, and thus cannot contribute appreciably to the H2 fraction. This serves as a reminder that large-scale aspects of the flow have a strong hand in the chemical evolution on the time-scale of the flow, and that one must be judicious in selecting dominant chemical processes on which to expend computing resources.

4 Ongoing Work In order to gain information on how to constrain or augment the set of reactions and cooling terms included, we use a simple test problem to compare the network with a more complex network. The stationary J-shock is an ideal problem for comparing the choice of chemical and cooling terms, as well as for carrying out parameter studies. The jump conditions for a hydrodynamic shock are evaluated, and the equations

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are solved in steady state for the post-shock flow, with the chemistry integrated at each step and the cooling effects coupled to the hydrodynamic flow via the pressure term in the momentum equation. The resulting profile for the evolution of the chemistry and flow variables for a theoretical stationary shock is readily compared with results from detailed models such as that presented in [5]. The initial comparison shows that the post-shock dissociation agrees quite well, but our cooling is slightly slower and the H2 formation rate lower. The inclusion of some previously omitted cooling terms, as well as dust-catalyzed H2 formation as described in [8] bring our simple network into surprisingly close range of the comparison model. These changes will be incorporated into the model for upcoming simulations. Acknowledgements The present work is supported by the European Community’s Marie Curie Actions-Human Resources Mobility with the JETSET (Jet Simulations Experiments and Theory) Network under Contract MRTN-CT-2004-005592. The work has been performed under the Project HPC-EUROPA (RII3-CT-2003-506079), with the support of the European Community - Research Infrastructure Action under the FP6 Structuring the European Research Area Programme.

References 1. Abel, T., Anninos, P., Zhang, Y., Norman, M., New Astronomy, 2, (1997), p181 2. Anninos, P., Zhang, Y., Abel, T.,Norman, M., New Astronomy, 2, (1997), 209 3. Cen, R., ApJ, 78, (1992) 4. Eisloeffel, J., Davis, C., Ray, T., Mundt, R. ApJ, 422, (1994) 5. Flower et al., MNRAS, 341, (2003) 6. Galli, D., Palla, F., 209, (1997) 7. Glover, S., ApJ, 331, (2003) 8. Hollenbach, D., McKee, C., ApJSS, 41:555–592 9. Mignone et al., ApJ Supplement Series, 170, (2007) 10. O’Connell et al., A&A, 431, 223–234, (2005) 11. Velusamy et al., ApJ, 668, 2, (2007)

Survival of Molecules in MHD Disk Winds Despina Panoglou, Sylvie Cabrit, Paolo J.V. Garcia, and Guillaume Pineau des Forˆets

Abstract Magnetohydrodynamical (MHD) models have been constructed and observations have been conducted thoroughly for atomic jets in the past. Observations of molecular jets implied the need of further modelling of jets that would include molecular chemistry. In this work we imposed a molecular chemical network on a self–similar steady state MHD disk wind, in order to study the origin and formation mechanisms of molecular jets. We calculated the radiation field as coming from the star and hot spots on its surface; this leads to an X-ray ionization rate, which, together with UV field induced photoreactions, raises the ionization fraction at the base of the flow. The main heating mechanism is the drag heating between charged and neutral particles, and after an initial increase of the temperature at low altitudes, it is balanced mainly by adiabatic and molecular cooling. Hence the temperature is maintained at low values, and molecules are indeed able to survive.

1 Introduction In the past, various attempts have been made towards the chemical modelling of winds. Glassgold et al. [9] computed the chemical evolution of the steady state spherical winds of Ruden et al. [14], with the temperature fixed from the original

D. Panoglou () Faculdade de Ciˆencias, Universidade do Porto, Portugal Universit´e Pierre et Marie Curie – Paris 6, France e-mail: [email protected] S. Cabrit LERMA, Observatoire de Paris-Meudon, France e-mail: [email protected] P. J. V. Garcia Faculdade de Engenharia, Universidade do Porto, Portugal e-mail: [email protected] Guillaume Pineau des Forˆets Institut d’Astrophysique Spatiale, Orsay, France e-mail: [email protected] K. Tsinganos et al. (eds.), Protostellar Jets in Context, Astrophysics and Space Science Proceedings, DOI 10.1007/978-3-642-00576-3 55, c Springer-Verlag Berlin Heidelberg 2009 

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solution. Safier [15] calculated the evolution in self–similar atomic disk winds that were not computed in consistency with the accretion disk. Garcia et al. [8] did the same with disk winds connected to the accretion disk [5]. Shang et al. [16] modelled the thermal evolution of X-wind gas that was irradiated by X-rays but included no chemistry. Molecular jets are usually attributed to shocks. The molecules that are observed can be considered (a) to appear in the jet through shocks with the gas entrained from the infalling envelope [13] or (b) they may be formed in shocks while the originally atomic gas propagates along the jet axis [11]. No modelling has taken place so far for the chemistry that dominates the evolution of molecular abundances inside jets. Towards this direction is our new model of molecular MHD disk winds, which includes the radiation field from the star, and where both the far UV and X-ray components are allowed to affect the gas, by heating, ionization, dissociation and excitation of the different species. Such thermal and chemical modelling of disk winds would allow us (a) exclude the presence of molecules in the jet itself if they cannot survive under difficult conditions for molecular maintainance or (b) add this hypothesis as an explanation for the observations of molecular jets, along with the formation in shocks.

2 Chemical Modelling of Disk Winds In this contribution we show the results of a chemical model of disk winds, so as to evaluate the chemical composition of the wind gas. For this, we imposed molecular chemistry from a model of shocks in star forming regions [7] to the dynamics of a warm self–similar disk wind [1]. The elemental abundances and the distribution of matter amongst the gas and the solid state is the same as used by Flower and Pineau des Forˆets [6]. The size distribution of dust is given by Mathis et al. [12] and its optical properties by Draine and Malhotra [4]. We take into account the sublimation of dust in the region close to the star, assuming that the sublimation temperature is 1,500 K. The dust temperature is calculated at each point by simple radiation equilibrium. The radiation field comes from the star, with an additional component of radiation flux from hot spots on the stellar surface, in order to account for the observed excess emission. The UV component of the radiation field is attenuated by the dust along the line of sight, while the X-rays are attenuated by all the nuclei that they find in their path. The results show that the temperature of the gas remains in moderate values ('2; 000 K, Fig. 1a). The gas temperature starts at the same value as the dust temperature, 100 K. For a small distance the dust temperature rises and then starts falling to the value of 70 K, as it follows the behavior of the attenuated UV flux. The evolution of the gas temperature is calculated by taking into account all the heating and cooling sources. All along the disk wind flow line, the ambipolar diffusion heating is the dominating heating mechanism. The ambipolar diffusion comes from the drift between charged and neutral particles, and the corresponding drift velocity

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(Fig. 1d) is reversely proportional to the ionization fraction (Fig. 1g). At altitude z ' 1 AU, the ionization fraction of the gas exhibits a low peak, the drift speed rises accordingly, and so does the gas temperature. From then on, the molecular cooling from excitation of H2 and other molecules almost balances the ambipolar diffusion and the temperature forms a plateau–like profile.

3 Molecular Constituent of the Wind Gas The abundances of molecules depend on a series of variables, which make complicated the study of the formation mechanisms in a chemical network composed of 150 species and 1,150 reactions. It must be mentioned that the majority of the reactions depend on the temperature and the drift speed, whose high values enhance

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the reaction rates. Exceptions from this feature are the rates of reactions induced by the UV field (Fig. 1f) and X-rays (Fig. 1e), which at the base of the wind increase because of the large negative derivatives of visual extinction and X-ray optical depth (Fig. 1c). But the drop in attenuation soon becomes less steep, and the decreasing radiation flux from the star (/1=R2 , where R is the spherical distance from source) dominates causing photoreaction rates to decrease. As a result there is only in part dissociation of H2 (Fig. 1h) and CO (Fig. 1i), since the flow time scale is too short to allow a total destruction: H2 dissociates mainly through neutral–neutral endothermic reactions such as O C H2 ! OH C H

(1)

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or through collisional dissociation by hydrogen atoms H2 C H ! H C H C H H2 is self–shielded against photoreactions [3], since its photodissociation is dominated by line absorption instead of continuum absorption. CO is also self–shielded and cross-shielded by H2 [10], but it is not efficiently protected, and the main dissociation mechanism for CO is photodissociation to atomic carbon and oxygen. By the end of the flow line 60% of the initial H2 and 50% of CO survive. Other molecules such as H2 O and OH appear to form out of the dissociation of H2 and CO, but the conversion from one form to another is not always a direct reaction. Photodissociation to OH is rapid in the H2 O chemistry, but it is mostly balanced by formation through (2). At the base of the flow H2 O starts with a very low fractional abundance (106 ), drops a little by the ion–neutral reaction HCOC C H2 O ! H3 OC C CO and then rises quickly through (2). The increase in the abundance of OH from 108 to 105 occurs mainly through the electron–ion recombination H3 OC C e ! OH C H2 in the beginning and through (1) along the main jet zone.

4 Conclusions We showed that it is possible that observations of molecular jets are not necessarily due to entrained gas from the envelope surrounding the protostar or because of shocks that create appropriate conditions for the formation of molecules. The presence of molecules can as well be due to molecular gas driven outwards from the

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accretion disk in bipolar flows by large scale magnetic fields. The complex chemistry induced from high volume densities, rather high ionization fraction, low drift speed and high attenuated radiation flux could potentially provide sufficient protection from dissociation of molecules in a disk wind. The next goal is the consideration of three thermal fluids instead of one, which would be more accurate in handling both reactions and heating/cooling rates. Giving up on the assumption of all particles being in uniform temperature may unveil various phenomena that were not taken into account so far but could be significant in the evolution of the fluid. Acknowledgements This work was supported by the European Community’s Marie Curie Actions Human Resource and Mobility within the JETSET (Jet Simulations, Experiments and Theory) network under contract MRTN-CT-2004 005592.

References 1. Casse, F., and J. Ferreira, 2000, A&A, 361, 1178. 2. Draine, B. T., 1978, ApJS, 36, 595. 3. Draine, B. T., and F. Bertoldi, 1996, ApJ, 468, 269. 4. Draine, B. T., and S. Malhotra, 1993, ApJ, 414, 632. 5. Ferreira, J., 1997, A&A, 319, 340. 6. Flower, D. R., and G. Pineau des Forˆets, 2003, MNRAS, 343, 390. 7. Flower, D. R., G. Pineau des Forˆets, and T. W. Hartquist, 1985, MNRAS, 216, 775. 8. Garcia, P. J. V., J. Ferreira, S. Cabrit, and L. Binette, 2001, A&A, 377, 589. 9. Glassgold, A. E., G. A. Mamon, and P. J. Huggins, 1991, ApJ, 373, 254. 10. Lee, H.-H., E. Herbst, G. Pineau des Forˆets, E. Roueff, and J. Le Bourlot, 1996, A&A, 311, 690. 11. Lim, A. J., A. C. Raga, J. M. C. Rawlings, and D. A. Williams, 2002, MNRAS, 335, 817. 12. Mathis, J. S., W. Rumpl, and K. H. Nordsieck, 1977, ApJ, 217, 425. 13. Micono, M., G. Bodo, S. Massaglia, P. Rossi, and A. Ferrari, 2000, A&A, 364, 318. 14. Ruden, S. P., A. E. Glassgold, and F. H. Shu, 1990, ApJ, 361, 546. 15. Safier, P. N., 1993, ApJ, 408, 115. 16. Shang, H., A. E. Glassgold, F. H. Shu, and S. Lizano, 2002, ApJ, 564, 853.

Sheared Magnetic Field and Kelvin Helmholtz Instability Matteo Bocchi, Hubert Baty, and Max Camenzind

Abstract In the context of plasma flows and jet stability, I studied the evolution of the Kelvin Helmholtz Instability (KHI) in the presence of a magnetic field reversal. The aim of the work is to compare this case with the known results from the uniform magnetic field case. Thus, I performed magnetohydrodynamic numerical simulations of a single shear layer and of a slab jet. I found significant differences in the growth rates and the behavior of the instability in the single shear layer simulations. Especially for high magnetic field strengths, the KHI presented higher growth rates and a more unstable nature. In the slab jet simulations, the magnetic field is amplified around the jet even for supersonic Mach numbers.

M. Bocchi () ZAH - Landessternwarte, Koenigstuhl 10 - 69117 Heidelberg, Germany e-mail: [email protected] H. Baty Observatoire Astronomique de Strasbourg e-mail: [email protected] M. Camenzind ZAH - Landessternwarte e-mail: [email protected]

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1 Introduction Astrophysical Jets show remarkable stability properties, as they are observed to propagate for long distances keeping a high degree of collimation. It is still up to debate how jets survive magnetohydrodynamic (MHD) instabilities (see for example the reviews by [3]). The work I present here focuses on a particular model of the interface between the jet and the ambient medium embedded in a magnetic field reversal. Some jet launching models like the X-wind [7] predict the presence of antiparallel magnetic field lines above the accretion disk. Numerical simulations of the star-disk interaction have shown the same scenario [2]. This physical problem has been studied already with a discontinuous magnetic field profile by [4], with emphasis on the numerical aspects. The aim of the work is to study the evolution of the Kelvin Helmholtz Instability (KHI) growing on the interface layer, and to compare the results with the well known literature on the case with uniform magnetic field. An upcoming paper will contain all the details.

2 Model and Numerical Methods In this section I describe the model and numerical methods used to obtain the results.

2.1 The Model Consider a cartesian coordinate system. Plasma flows initially only along the x direction following a simple hyperbolic tangent profile: vx .y/ D

y  V0 tanh : 2 a

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The magnetic field is reversing with the same hyperbolic tangent profile, for the sake of simplification: y  : (2) Bx .y/ D B0 tanh a Here a is the width of the shear layer. V0 and B0 are parameters that vary from case to case, so they will be specified later in Sect. 3. Both vy and By are initially set to zero. The layer is assumed to be in pressure equilibrium. The density profile is chosen accordingly to keep the temperature constant in the domain.

2.2 Numerical Methods I implemented the initial conditions described above in the Pluto code. Pluto [5] is a finite volume, shock capturing fluid dynamical code designed to integrate the

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(R)(M)HD (special Relativistic)(Magneto)HydroDynamic equations in conservative form; several solvers and integration options are available, as well as non-ideal modules, like cooling and resistivity (http://plutocode.to.astro.it). To prepare the MHD simulations, I used the linear code Ledaflow [6] to compute the growth rates of the instability.

3 Results This section provides the results obtained both from the numerical simulations of a single shear layer as described in Sect. 2, and from simulations of two layers to form the so called slab-jet. For the single layer simulations, I found significant differences from the uniform case in term of growth rates, saturation and disruption levels. The slab-jet simulations presented an interesting behavior and enhancement of the magnetic field at the interface, an effect previously believed to be present only in subsonic jets (see [8]).

3.1 Single Layer Simulations Single layer simulations are performed using a Mach number equal to 1. V0 is chosen accordingly. The magnetic field strength is controlled by the Alfv´en Mach number (MA ), that I varied between 2:5 and 100.

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My results from the Ledaflow code are summarised in Fig. 1. Two important differences are present from the uniform magnetic field case: 1. The wave vector k corresponding to the maximum growth rate does not depend on the magnetic field strength, but is constant for all the values of MA . In the uniform case, low values of MA shift the maximum towards smaller wave vectors. 2. The growth rates for MA smaller than 5 are sensibly higher than the uniform case.

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The previous points indicate that the character of the instability is different from the uniform case. This is confirmed by the saturation levels and by the ranges of unstable MA . For uniform magnetic fields we know from [1] that when 2 < MA < 4

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Fig. 1 Growth rates  as a function of the wave vector k for different values of the Alfv´en Mach number

Fig. 2 Density plot of a single layer simulation at 7:4 time units. For this realization MA D 7. The image is scaled with gray tones corresponding to the range Œ0:7; 1:16

the KHI is unstable, but stabilized by non linear processes. This effect is not present in my reversed field case, that is fully unstable for any value of MA > 2. Figure 2 represents the density structure at the time of saturation of the instability. One can clearly notice the plasmoid structures along the low density spiral structure typical of the KHI. These features proved to be magnetic islands driven by the KHI. The final state of the simulations is also a magnetic island positioned at the centre of the domain, an expected outcome of the initially reversed magnetic field. The disruption level, measured as the quantity of kinetic energy that is taken away from the initial flow, is shown in Fig. 3. For values of MA < 5 the reversed field is more desruptive than the uniform case. On the contrary, it is less desruptive for values of MA > 6.

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Fig. 3 Desruption levels of the KHI for different MA . Stars represent the uniform magnetic field case, while crosses represent the reversed field case

Fig. 4 Upper panel: Density plot at 80 time units. Lower Panel: Plot of the velocity parallel to the jet axis at 150 time units. The contours represent the Mach number. External dark line: M D 1, Inner line: M D 3

3.2 Slab-jet Simulation For Slab-jet simulations, the Mach number is set to 3, and the Alfv´en Mach number to 7. The simulations are performed in a spatial approach, so the plasma is flowing into the domain from the left side of the x axis, and is flowing out of the domain on the opposite side. I found several interesting results, summarised here: 1. Early in the instability evolution there is a transition from pinch to sinusoidal modes, and from short to long wavelengths (see Fig. 4, upper panel). 2. The jet is episodically disrupted by the instability, but the flow is later revived and keeps its coherence (see Fig. 4, lower panel). 3. The compression of the external medium due to the sinusoidal deformation of the jet causes a magnetic field enhancement propagated backwards towards the inflow boundary through Alfv´en waves (see light-color structures in Fig. 5). 4. The inner part of the jet and the plasma left by the jet disruption display a turbulent magnetic structure (see Fig. 5).

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Fig. 5 Gray scale plot of the magnetic field strength at 150 time units. Only the right half of the domain is shown

−

Acknowledgements The present work is supported by the European Communities Marie Curie Actions-Human Resources Mobility with the JETSET (Jet, Simulations, Experiments and Theory) Network under Contract MRTN-CT-2004-005592. Part of the simulations for this work were performed on the JUMP computer of the John von Neumann Institute for Computing, Forschungszentrum Juelich.

References 1. Baty, H., Keppens, R., Comte, P., Physics of Plasmas., 10, 4661, (2003) 2. Bessolaz, N., Zanni, C., Ferreira, J., Keppens, R., Bouvier, J., A&A., 478, 155, (2008) 3. Hardee, P.E., ApJ., 664, 26, (2007) 4. Keppens, R., T´oth, G., Westermann, R.H.J., Goedbloed, J.P., Jurnal of Plasma Physics., 61, 1, (1999) 5. Mignone, A., Bodo, G., Massaglia, S., et al., ApJS., 170, 228, (2007) 6. Nijboer, R., van der Holst, B., Poedts, S., Goedbloed, J.P., Comput. Phys. Commun., 101, 39, (1999) 7. Shu, F., Najita, J., Ostriker, E., et al., ApJ., 429, 781, (1994) 8. Viallet, M., Baty, H., A&A., 473, 1, (2007)

Jets from Class 0 Protostars: A Mid-IR Spitzer View Odysseas Dionatos

Abstract We present Spitzer - IRS staring and spectral mapping observations of the L1448-C and HH211-mm outflows. Atomic lines from the fundamental transitions of [FeII], [SiII] and [SiI] have been detected in both cases indicating for the first time the presence of an embedded atomic jet at low excitation conditions. Rotational H2 lines are also present in the IRS spectra, and can be attributed in the inner jet regions to warm molecular gas enveloping the atomic jet. In the analysis that follows, we probe for the excitation conditions of the H2 and atomic emission applying a number of emission line diagnostics and employing shock models.

1 L1448-C We have observed towards the driving source L1448-C and along its outflow up to the bowshock position using the low (SL) and high (SH and LH) resolution IRS modules. In the left panel of Fig. 1 we present the placement of the various IRS slits over an IRAC band 2 image and define four positions for which spectra were extracted. Additional NIR spectra from UKIRT/UIST were obtained in order to constrain the excitation conditions along the jet.

1.1 H2 Emission We have probed the H2 excitation conditions along the jet by means of excitation diagrams (Fig. 1, right panel). H2 fluxes were dereddened after calculating the optical extinction at each observed point by examining the silicate absorption feature. The gas temperature was derived from the slope of a least square fitted line

O. Dionatos () INAF - Osservatorio Astronomico di Roma, Italy e-mail: [email protected]

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Fig. 1 (left:) IRS slit distribution over an IRAC band 2 image of the L1448-C region; spectra were extracted for each of the four defined regions (CS - OF2). (right:) Excitation diagrams for the driving source (CN) and bow shock (OF2) positions. Spitzer (open circles) and UKIRT (open squares and arrows for upper limits) data points are optimally least square fitted (solid line)

(see Fig. 1. Derived temperatures range from 600 to 900 K, except the bowshock position where the near-IR rovibrational lines trace a second temperature component with T  2;700, probably arising from the unresolved cooling zone behind the bowshock. The H2 column density was derived from the intersection of the fitted line to the data points with the abscissa of the excitation diagram; combined with kinematical information given by sub-mm interferometric studies [3] mass flux rates of 107 Mˇ yr1 were derived for the molecular component. In the bow shock, near and mid-IR emission have been compared with a C-type bowshock model [8]. Shock velocities of 100 km s1 and pre-shock densities n.H2 /  105 cm3 do account rather well for the observed column densities.

1.2 Atomic Emission Strong emission from the fundamental transitions of [FeII], [SiII] and [SI] has been detected along the jet. The non-detection of higher excitation [FeII] lines falling in the IRS wavelength range in conjunction with the absence of the [NeII] line at 12.8 m suggest that the excitation conditions of the atomic jet are very low. We probe the excitation conditions by means of atomic/ionic line ratio diagnostics. In the left panel of Fig. 2 we plot the [SiII]34.8 m/[FeII]26 m ratio as a function of electron density for temperatures of 1,000 K and 2,000 K. The diagnostic was realized employing a statistical equilibrium model that takes into account the first 16 levels for [FeII] and a two level system for [SiII]. Hatched areas represent the observed ratios at the CN and OF2 positions respectively, implying values of ne < 400 cm3 The temperature of the atomic component can be constrained from the diagnostic diagram presented in the right panel of Fig. 2. In this, the upper limit ratio of [FeII]18 m/[FeII]26 m is plotted against the gas temperature, for different electron densities; derived upper limit temperatures are 2,500 K and 1,500 K for the OF2 and CN positions respectively.

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Fig. 2 Atomic line diagnostic diagrams. (left:) [SiII]34.8 m/[FeII]26 m ratio versus electron density and (right:) [FeII]18 m/[FeII]26 m ratio versus temperature. Horizontal hatched areas and lines mark the observed values

We estimate the mass flux for the two species using the methodology of [7], using as in the case of H2 kinematical information from sub-mm studies. The atomic component mass flux is of the order of 106 Mˇ yr 1 , an order of magnitude higher than the derived H2 values; comparing with the CO momentum flux of the corresponding molecular outflow it results that the atomic jet has the power to drive the large scale outflow, and thus may represent the primary jet component ejected by the source.

2 HH211-mm For the case of HH211 we have deployed the low resolution IRS modules in slitscan mode to cover an area of 5700  7300 and 15700  16800 for the Short Low (SL) and Long Low (LL) modules respectively centered on the Class 0 source HH211mm. A perpendicular slit step equal to the slit width of each module was used to perform subsequent observations and consequently, the pixel scale of the resulting spectral line maps differs substantially in the two cases, being 3:500 and 10:500 for the SL and LL modules respectively. As in the case of L1448 along the outflow we have detected the full series of H2 pure rotational lines (S(0)–S(7)) as well as atomic lines from the fundamental transitions of [FeII], [SiII] and [SI].

2.1 Morphology In the upper left panel of Fig. 3 we present the emission line map of the S(5) rotational H2 line. Contours shape a characteristic bipolar outflow pattern where H2

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Fig. 3 upper panel H2 S(5) left and [FeII] right spectral line maps of the HH211 outflow. Peaks of emission discussed in the text are labeled. Contour levels are form 2 1012 W cm2 sr1 with increment of 2.5 1012 W cm2 sr1 for the H2 and from 1012 W cm2 sr1 with an increment of 1013 W cm2 sr1 for the [FeII] maps. (lower panel:) The H2 S(5) line map overlaid on SiO J D 8–7 and near-IR emissions - plot is rotated by 26.6ı clockwise with respect to the H2 map

emission becomes significant for a projected angular distance >500 from the driving source HH211-mm. Further downwind, peaks of emission which can be attributed to shocked gas are named as B1–B3 for the southeast, and R1–R2 for the northwest lobes respectively. The limited spatial coverage of the SL maps does not cover though the southeast bow shock region. The lower panel of Fig. 3 shows the H2 S(5) map (green) overlaid on the SiO J D 8–7 emission of [6] on a near-IR image [4] (grayscale). Morphologically the H2 emissions in the near and mid-IR display very similar characteristics, while the mid-IR and SiO emissions are spatially coincident only in the southeast lobe B1 peak; in the northwest lobe, the extended pedestal as seen in the near and mid-IR is uncorrelated to the SiO jet and could be attributed to scattered light on the walls of a cavity. In the upper right panel of Fig. 3 we present the emission line map of the [FeII] line, which is substantially weaker in comparison to the H2 maps. The larger spatial extend of LL module maps does cover the bow-shock areas, where [FeII] emission peaks.

2.2 H2 Emission As in the case of L1448, we have employed excitation diagrams to derive the excitation conditions along the HH211 jet; however in the current analysis we used only

Jets from Class 0 Protostars: A Mid-IR Spitzer View Fig. 4 Kinetic temperature map produced by the H2 emission falling in the SL range. The white rectangle delineates the boarders of the mapped area. Shocked gas reaches temperatures up to 1,300 K, where quiescent areas display temperatures as low as 700 K

475 1500

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0

the S(2)–S(7) transitions that fall within the SL module. Initially we focused on the peaks of emission where all 6 H2 transitions are detected. We have derived the ortho to para ratio of 2.5–2.8 and visual extinction of 7–9 mag. by examining the alignment of the column densities S(5) and S(3) transitions with their neighboring data points. Considering an average AV D 8, the gas temperature was derived for all map points that 3 or more H2 lines were detected. In the temperature distribution of Fig. 4, left panel, peaks of temperature up to 1,300 K are spatially coincident with peaks of emission as seen in the upper left panel of Fig. 3, while in the more diffuse gas, the temperature is as low as 700 K.

2.3 Atomic Emission The atomic emission along the HH211 jet was treated as in the case of L1448. The derived electron density of
References 1. Asplund, M., Grevesse, N., & Sauval, A. J., 2005, ASPC, 336, 25 2. Dionatos, O., Nisini, B., Garcia Lopez, R., Giannini, T., Davis, C. J., Smith, M. D., Ray, T. P., De Luca, M., 2008, ApJ, in press 3. Girart, J. M.., & Acord, J. M. P., 2001, ApJ, 552, 63 4. Hirano, N., Liu, S.-Y., Shang, H., Ho, P. T. P., Huang, H.-C., Kuan, Y.-J., McCaughrean, M. J., & Zhang, Q., 2006, ApJL, 636, L141

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5. Kristensen, L. E., Ravkilde, T. L., Pineau Des Forˆets, G., Cabrit, S., Field, D., Gustafsson, M., Diana, S., & Lemaire, J.-L., 2008, A&A, 477, 203 6. Lee, C.-F., Ho, P. T. P., Palau, A., Hirano, N., Bourke, T. L., Shang, H., & Zhang, Q., 2007, ApJ, 670, 1188 7. Nisini, B., Bacciotti, F., Giannini, T., Massi, F., Eisl¨offel, J., Podio, L., & Ray, T. P., 2005, A&A, 441, 159 8. Smith, M. D., 1991, MNRAS, 252, 378

0.1500 Study of the Atomic and Molecular Jets in DG Tau Vanessa Agra-Amboage, Catherine Dougados, and Sylvie Cabrit

Abstract We present the results of a spectro-imaging study of the DG Tau microjet in [FeII] and H2 lines emission allowing for the first time to compare the atomic and the molecular jet properties. Our observations were taken using the Integral Field Spectrograph SINFONI at VLT, combined with adaptive optics, yielding a spectral resolution of 100 km/s and spatial resolution of 0:1500 D 20 AU. A nested flow morphology in both atomic and molecular components is observed. We found asymmetries in the atomic jet suggesting precession and a cavity in the molecular emission. We estimated the mass-loss rate for both components using different methods, we discuss in the text the implications of the values obtained.

V. Agra-Amboage () and C. Dougados Laboratoire d’Astrophysique de l’Observatoire de Grenoble, UMR5521 du CNRS, 38041 Grenoble Cedex, France e-mail: [email protected]; [email protected] S. Cabrit Observatoire de Paris, LERMA, UMR 8112 du CNRS, 61 Avenue de l’Observatoire, 75014 Paris, France e-mail: [email protected]

K. Tsinganos et al. (eds.), Protostellar Jets in Context, Astrophysics and Space Science Proceedings, DOI 10.1007/978-3-642-00576-3 58, c Springer-Verlag Berlin Heidelberg 2009 

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1 Introduction Jets from young stars are a ubiquitous part of the star formation process and their link with accretion into the star is also clear. They are launched almost certainly from the inner regions of the star-disk system. Although a number of models exist for their generation, the exact mechanism by which mass is ejected from accreting systems and collimated into jets remains an open problem. There are several possible scenarios for the origin of the outflow: the stellar surface, the inner edge of the accretion disk, an extended range of disk radii or re-connexion sites in the stellar magnetosphere. The main aim of this work is to find observational constraints on the jet launching region in order to distinguish between these scenarios. To attain this objective we analyzed IR spectro-imaging observations taken on October 15th 2005 at the Very Large Telescope (VLT) using the integral field spectrograph SINFONI and combined with an adaptive optics module. We observed in both H and K bands with spatial resolution of 0.1500 and spectral resolution of 100 km/s. These observations allow us to probe the atomic jet in the inner 150 AU from the star, but also the molecular H2 transitions tracing the warm molecular component, allowing to probe the structure of the flow in a broad range of excitation conditions. We will discuss here the jet morphology and excitation conditions (mass fluxes in particular).

2 Jet Morphology In the deconvolved images, Fig. 1, the main jet characteristics are clear. The atomic blue jet is more collimated in HVC than in IVC and it is more collimated than the molecular component. The atomic HVC is especially collimated until 100 from the star and then it becomes wider, probably due to the knot at 1.200 . At lower velocities the jet shows a clear v-shaped morphology, opening gradually until the field limit. On the contrary, the red counter-jet is wider and shows a bow-shock structure in the two velocity components. On the other hand, molecular emission is confined in a small region (about 0.500 in the main blue emission). The red emission is faintly detected but it shows a bow shock structure. On the contrary, the blue emission is clearly detected and it shows a cavity at 0.300 from the star. This cavity was already mentioned in [3]. We measured the collimation degree on the blue side of the deconvolved images measuring the FWHM at different distances from the star. We obtain a half opening angle of 5ı for the atomic HVC and 13ı for the IVC. The molecular emission is much more wider, with an opening angle of 32ı . From the minimum FWHM value measured we derive a limit on the ejection radius of 0.1500 (17 AU) for H2 and even smaller (<7 AU) in the case of the atomic jet. In [2] similar widths were found in optical lines observed with STIS on the HST.

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Fig. 1 Continuum-subtracted maps, raw data (top panels) and deconvolved (bottom panels). H2 emission map is integrated over one velocity inerval, centered at 0 km/s. HV atomic emission is centered in 175 km/s and in C225 km/s for the blue and red jet respectively. Atomic IV is centered in 50 and C75 km/s

In addition, in the atomic HVC we detect a hint of wiggling. If it is real, it has a period of 1.200 that, assuming a typical jet velocity of 300 km/s at Taurus distance (140 pc), means a period of 3 years. A detailed model will be used to probe the jet precession possibility and accurately measure its precession period.

3 Atomic [FeII] Component: Excitation Conditions and Mass-loss Rates A very useful technique to constrain the physical conditions of the gas is to use line ratios between forbidden lines. In particular [FeII]1.53 m/1.64 m ratio is very sensitive to the electron density. Figure 2 gives the line ratio map integrated over the whole line, that is, it gives the density distribution in the jet. The red counter-jet is quite uniform with a density of 104 cm3 , whereas in the blue jet density decreases with the distance to the star. Far from the star we find values of 104 cm3 increasing to values of 105 cm3 . The blue jet shows in addition asymmetries about the jet axis, especially close to the star, exhibiting higher densities to the East and reaching the high density limit, with densities higher than 5  105 cm3 .

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Fig. 2 Electron density distribution obtained from the [FeII]1.53 m/1.64 m ratio. Contours are in cm3 . The high density limit is reached for electron densities higher than 5  105 cm3 . Close to the star, at the East we do not plot contours higher than 1  105 cm3 because electron density in this region is higher than 5  105 cm3

Another parameter important to constrain the models is the mass-loss rate. We used three different methods to estimate it (described in detail in [1]). The first method is based on estimates of the jet density, nH (obtained from ne and xe ) and the jet cross-section rj , and assumes a filling factor of 1. The second method uses the observed luminosities of the optically thin forbidden emission line, [FeII]1.64 m, and it assumes that Te , ne and xe are uniform in each pixel. The last method used uses also line luminosity, but assuming that this emission is due to single shocks inside the jet. We used models computed by [6] to obtain the shock flux mass. To obtain the jet mass flux we correct MP s of Vj=Vs D 4 from [9] and we assume one shock every 50 AU as [5] show in the case of HH30. Figure 3 shows the mass-loss rate as a function of the distance to the star for the three methods. The stars mark the value found by [4] using the same method based on jet density and cross-section as us. The difference with our estimation comes mainly from the different rj used, they found a two times higher value than us. In fact this method depends strongly on the estimation of this parameter. Lavalley [8] found also higher values using this method, and even higher than [4], for the optical ˚ with higher observed jet cross-section. For the other two methods, line [OI]6,300 A [8] found for all distances very similar mass fluxes as us. On the other hand, the

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Fig. 3 Mass-loss rate as a function of the distance to the star for the atomic blue jet for the three different methods (HVC on the left and IVC on the right). We also plot the [FeII]1.64 m luminosity for each kinematical component. Stars mark the values obtained by [4]

good agreement between the method based on the jet density and the cross-section and that based on the emission on a uniform slab (using the luminosity estimations) for the HVC suggests that the Fe is not depleted into grains and that all the Fe is in gaseous form. That is not the case for the IVC, where we found a difference between the two volume methods of one order of magnitude. Moreover, for the two methods using Fe luminosities, we found that IVC is lower than HVC following the lower luminosity in this component. However, the method based on the cross-section is higher in the IVC because of the higher cross-section. Following equations described by [10] we estimated the accretion rate from the Br luminosity and we obtained MP D 107 Mˇ /yr. Hence the ejection-accretion ratio has values between 0.06 and 0.2.

4 Molecular Cavity (H2 ) In order to characterize the molecular emission we estimate the mass-loss ratio of H2 . We used two methods, both using the luminosity in H2 2.12 m, but one supposing a volume emission and the other supposing shocks. From a H2 luminosity integrated over the blue lobe of 1:6  105 Lˇ we estimate the total H2 mass in 3  108 Mˇ (details of calculations in [1]). Assuming a speed of 15 km/s [3] over an emitting region of 0.500 , we estimate a mass-flux of 1:5  109 Mˇ =yr. If we suppose that this emission comes from a slow molecular MHD wind heated by ambipolar diffusion, and following [11] predictions, we derive a launching radius lower than 10 AU with a small magnetic lever arm ( < 3). In this case the ejectionaccretion ratio is 0.02. However, if we assume that this emission is generated by a shocked slow wide angle wind, we have two possibilities: J-shocks and C-shocks. J-shocks can not reproduce the observed H2 luminosity and 2-1S(1)/1-0S(1) ratio equal to 0.06 taken from [3]. Using C-shock models, computed by [7], we derive a shock mass-flux between 4108 and 4109 Mˇ =yr depending on ‘b’ (b  VA ). Hence, we derive a MP ej =MP acc of 0.015 for the first method and between 0.04 and 0.4 for the second.

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5 Conclusions We summarize the main conclusions of this work:  We found a nested flow morphology in both atomic and molecular components.

Moreover H2 emission shows a cavity encompassing atomic flow.  We observed asymmetries in the morphology of the jet and in the electron den-

sity. They are probable precession signatures, a detailed modelling is underway.  We derived mass fluxes. In the atomic HVC, we found that values obtained

with the model based on the jet density and cross-section are compatible with those obtained from the Fe luminosity, assuming a uniform slab and suggesting no depletion into grains. For the molecular component, two possible scenarios were presented, a slow molecular MHD wind heated by ambipolar diffusion or a chocked slow wide angle wind. In this last case, only C-Shocks are compatible with the observations. A detailed study including kinematics and a thorough discussion of H2 excitation processes will be presented in [1]. Acknowledgements Vanessa Agra-Amboage and Sylvie Cabrit wish to acknowledge financial and travel support through the Marie Curie Research Training Network JETSET (Jet Simulations, Experiments and Theory) under contract MRTN-CT-2004-005592.

References 1. V. Agra-Amboage, C. Dougados, S. Cabrit, P. Garcia, (DG Tau) in prep, (2008) 2. F. Bacciotti, R. Mundt, T.P. Ray, J. Eisl¨offel, J. Solf, M. Camezind, ApJ, 537, L49, (2000) 3. T.L. Beck, P.J. McGregor, M. Takami, T.S. Pyo, ApJ, 676, 472, (2008) 4. D. Coffey, F. Bacciotti, L. Podio, ArXiv e-prints, (2008) 5. P. Hartigan, A. Frank, P. Varni´ere, E.G. Blackman, ApJ, 661, 910, (2007) 6. P. Hartigan, J. Raymond, R. Pierson, ApJ, 614, L69, (2004) 7. L.E. Kristensen et al., in prep, (2008) 8. C. Lavalley, Etude de la Morphologie et de la Cin´ematique de l’Emission des Raies interdites autour des Etoiles T Tauri. Ph.D. thesis, AA (Universit´e Joseph Fourier (Grenoble 1)), (2000) 9. C. Lavaley-Fouquet, S. Cabrit, C. Dougados, A&A, 356, L41, (2000) 10. J. Muzerolle, L. Hartmann, N. Calvet, AJ, 116, 2965, (1998) 11. D. Panaglous et al., in prep, (2008)

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Velocity Resolved IR Diagnostics of Class I Jets Rebecca Garc´ıa L´opez, Brunella Nisini, Teresa Giannini, Jochen Eisl¨offel, Francesca Bacciotti, and Linda Podio

Abstract We present VLT-ISAAC spectra at medium resolution of the HH34 and HH46-47 protostellar jets covering diagnostic lines of the ionised ([FeII] 1.644, 1.600 m) and molecular (H2 2.122 m) gas. These data have been used to measure the electron density (ne ) and mass ejection flux (MP jet ) as a function of velocity. In both jets, the electron density is higher in the Low Velocity Component (LVC), although the mass flux is mainly contributed from the High Velocity Component (HVC) that is associated to the large scale jet. In the HH34 H2 Position-Velocity diagram (PVD) the HVC and LVC appear separated suggesting that they may originate from physically distinct regions.

R. Garc´ıa L´opez (), B. Nisini, and T. Giannini INAF-Osservatorio Astronomico di Roma, Via Frascati, 33, 00040 Monte Porzio Catone, Italy e-mail: [email protected] J. Eisl¨offel Th¨uringer Landessternwarte Tautenburg, Sternwarte 5, D-07778 Tautenburg, Germany F. Bacciotti INAF-Osservatorio Astrofisico di Arcetri, Largo E. Fermi 5, 50125 Florence, Italy L. Podio Dublin Institute for Advanced Studies, Fitzwilliam Place 31, Dublin 2, Ireland K. Tsinganos et al. (eds.), Protostellar Jets in Context, Astrophysics and Space Science Proceedings, DOI 10.1007/978-3-642-00576-3 59, c Springer-Verlag Berlin Heidelberg 2009 

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HH 46–47

Fell (1.64 μm)

Fell (1.64 μm)

HH 34 H2 1-0S(1) 4

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rA

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Offset (arcsec)

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rA

A6

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–300 –200 –100 0

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–200

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–200 –100 0 100 200 –200–100 0 100 200 300 400 Velocity (km/s) Velocity (km/s)

Fig. 1 Continuum-subtracted PV diagrams for the [FeII] 1.644 m and H2 2.122 m lines of the HH34 and HH46-47 jets in the region nearest to the source. A P.A. of 15ı and 57ı were adopted for HH34 and HH46-47, respectively. Contours show values of 5, 15, 45, 135, 260  for both lines in the case of HH34, while for the HH46-47 jet [FeII] contours indicate values of 3,9,. . . ,243  and 4.5,13.5,. . . ,364.5  for the [FeII] and H2 lines

1 Introduction Protostellar jets play an important role in the star formation process. They are believed to remove angular momentum from the star-disc system, disrupt infalling material, etc. Little is known about the jet origin, however. Several optical studies at the base of jets from classical TTauri stars (CTTSs) have been performed in an attempt to get some information about the launching mechanism. Very few studies have been done, however, regarding the internal regions of jets from Class I jets near the source. These objects are highly extincted and thus not detected at optical wavelengths close to the central source. In this context, infrared (IR) observations turn out to be a very useful tool, firstly, because of extinction and secondly, because Class I objects exhibit a rich IR spectrum. In fact, several atoms and molecules have bright transitions falling in the near-IR (NIR) section of the spectrum (e.g. [FeII], and H2 ) that can be used to diagnose important physical properties of the jets. Here, we present velocity-resolved NIR spectra of the HH34 and HH46-47 jets (two classical Class I jets) obtained with the spectrograph ISAAC at VLT. We have detected important diagnostic lines of the ionised ([FeII] 1.644, 1.600 m) and molecular (H2 2.122 m) gas, that allow us to study the kinematics of the atomic and molecular components, and to derive ne and MP jet as a function of the jet radial velocity. The kinematic features near the source and the derived physical parameters are briefly compared with those found for CTTS jets.

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2 Kinematics of the HH34 and HH46-47 Jets near the Exciting Sources The position-velocity diagrams (PVDs) of the [FeII] 1.644 m and H2 2.122 m lines for the HH34 and HH46-47 jets are shown in Fig. 1. On the PVD y-axis, the distance from the source is reported, while the x-axis represents the velocity with respect to the local standard of rest (LSR) and corrected for a cloud velocity of 8 km s1 and 20 km s1 for HH34 and HH46-47. The [FeII] component in both jets presents a similar behaviour. Close to the central source, the [FeII] lines broaden and emission at lower velocities, down to 0 km s1 , appears within 200 from the central sources. Inside 100 emission is also detected at positive velocities, that in the case of the HH34 jet reconnects with the spatially resolved red-shifted knot rA (newly detected here). For both jets, we identify the high blue-shifted velocity corresponding to the large scale jet as the HVC and the emission component at lower velocity as the LVC, in analogy to the HVC and LVC observed in the forbidden emission line (FEL) regions of T Tauri stars. In CTTS jets the LVC is, however, confined within a distance of d 6 200 AU from the source, while in HH34 and HH46-47 it can reach distances up to 1,000 AU. We remark that in T Tauri FELs one usually observes two separated peaks spatially located at different offsets from the central source, with the HVC being commonly displaced further downstream (e.g. [7]). Here, the HVC has a peak on-source while the LVC is seen as a weaker shoulder of the brightest component. This behaviour could be due to the lower resolution of our observations as a consequence of the larger distance of HH34 and HH46-47 with respect to other studied T Tauri stars. At variance with [FeII], the H2 PVDs for both objects show a different behaviour. In fact, the H2 PVD for the HH34 jet shows spatially- and kinematically-separated LVC and HVC, and only the LVC is visible down to the central source. This component is close to 0 km s1 LSR velocity and the blue-shifted and red-shifted jets differ by less than 10 km s1 . The HVC appears at a distance of 200 from the central source, at the position of the knots A6 and rA. At intermediate velocities between these two components, no emission is seen in the H2 PVD. This suggests that the two components correspond to physically distinct regions. It could be that the LVC of the H2 lines in HH34 is excited by oblique shocks occurring at the wall of a cavity created by the interaction of a wide angle wind with the ambient medium. Another possibility is that the H2 gas originates from the external layers of a disc-wind. The H2 PVD for the HH46-47 jet shows, on the other hand, a single velocity component near the source.

3 Diagnostics of Physical Parameters The electron density in the atomic jet component can be derived from the ratio of the [FeII] 1.600/1.644 m (e.g., [9]). Taking advantage of the velocity resolved profiles

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in both the 1.644 and 1.600 m lines, we have measured the electron density in the different velocity components. With this aim, we have extracted the spectra of the two lines at the positions of the inner knots along the flows and measured the 1.600 m/1.644 m intensity ratio in each pixel along the spectral profile. Figure 2 shows the normalised profiles of the two lines and their ratio as a function of velocity for the knots A6 and S of HH34 and HH46-47. The spatial intervals used to extract the spectra of the knots are given in Table 1. The plotted ratio gives a qualitative indication on how ne varies in the different velocity components, that is, a higher ratio indicates a higher electron density. Therefore, Fig. 2 shows that the electron density increases with lower velocities in both the jets in the regions closest to the central sources. The derived values of the electron density in the two components are given in Table 1. Our result is in agreement with what is found in the FEL regions of T Tauri stars [5, 6]. In such studies, the LVC was not spatially resolved and the different values for ne found between the HVC and LVC were interpreted assuming that the LVC originates from a dense compact region close to the disc surface, while the HVC is associated with a more extended high-velocity jet displaced further out. Such an interpretation was supported by the spatial offset from the central source often found between the two components of the forbidden lines detected in spatially-unresolved spectra of T Tauri stars. Spatially-resolved measurements of the electron density in the HV and LV components have been performed in DG Tau [2, 8]. In this case, at variance with HH34, the electron density near the source has been found to increase with velocity. On the other hand, in the DG Tau micro-jet the total density in the LV component is higher than in the HV component because of the much lower ionisation fraction. Therefore, the density structure in HH34 and in the T Tauri stars might be similar, but they differ in the excitation conditions.

Fig. 2 Normalised line profiles (lower panels) of the [FeII] lines 1.644 m (solid line) and 1.600 m (dotted line), and their ratio in each velocity channel (upper panel) for different extracted knots along the HH34 and HH46-47 jets. Higher ratios indicate higher electron densities

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Table 1 MP jet and ne values for the most internal knots of HH34 and HH46-47 Jet Knot rt AV ne (HVC) ne (LVC) MP jet (HVC) HH34 HH46-47

A6 S

(00 ) (2.6, C0.9) (2.8, C3.4)

(mag) 7.1 6.6

(103 cm3 ) 10.5 22.5 5.4 7.6

(Mˇ yr1 ) 5.2 108 2.6 108

MP jet (LVC) (Mˇ yr1 ) 6.6 109 4.9 109

The MP jet value was obtained here from the luminosity of the [FeII] 1.644 m line, adopting the relationship MP D mH  .nH V /  vt = lt , where  D 1:24 is the average atomic weight; mH and nH are the proton mass and the total density; V is the emitting region; and vt and lt are the velocity and length of the knot, projected perpendicular to the line of sight. The product nH V can be expressed as  1 C L.li ne/ hAi fi FFee ŒFHe where Ai and fi are the radiative rates and fractional C

population of the upper level of the considered transition and FFee is the ionisation

 fraction of the iron having a total abundance with respect to hydrogen of FHe . We have assumed that all iron is ionised, and has a solar abundance of 2:8  105 ([1], i.e., no dust depletion). To compute the fractional population, we have used the ne values derived separately for the LV and HV components, while we have assumed a single temperature for both components, equal to 7000 K for HH34 [10] and 10,000 K for the HH46-47 jet. Tangential velocity values have been derived from the radial velocities assuming an inclination angle i D 34ı for HH46-47 and i D 22:7ı for HH34 [4, 3]. The luminosity of the line has been computed by integrating the extinction corrected flux of the knot in the same range of velocity used to calculate the electronic density. In Table 1, the derived values of MP jet for both jets are reported. Our results indicate that MP jet (LVC) is lower than MP jet (HVC) for both jets, despite the fact that ne in higher in the LVC. In jets from CTTSs, [6] have shown that MP jet (LVC) should be lower than in the HVC: from the analysis of not spatially-resolved optical spectra, they concluded that the LVC in jets from CTTSs could be responsible for carrying the majority of mass and momentum only if the emitting region were smaller than 1 AU. Acknowledgements The present work was partly supported by the European Community´s Marie Curie Actions - Human Resource and Mobility within the JETSET (Jet Simulations, Experiments and Theory) network under contract MRTN-CT-2004005592.

References 1. Asplund, M., Grevesse, N., Sauval, A.J. In: T.G. Barnes III, F.N. Bash (eds.) Cosmic Abundances as Records of Stellar Evolution and Nucleosynthesis, Astronomical Society of the Pacific Conference Series, vol. 336, pp. 25–+, (2005) 2. Bacciotti, F., Mundt, R., Ray, T.P., Eisl¨offel, J., Solf, J., Camezind, M. ApJ, 537, L49–L52, (2000) 3. Eisl¨offel, J., Mundt, R. A&A, 263, 292–300, (1992)

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4. Eisl¨offel, J., Mundt, R. A&A, 284, 530–544, (1994) 5. Hamann, F., Simon, M., Carr, J.S., Prato, L. ApJ, 436, 292–300, (1994) 6. Hartigan, P., Edwards, S., Ghandour, L. ApJ, 452, 736–+, (1995) 7. Hirth, G.A., Mundt, R., Solf, J. A&AS, 126, 437–469, (1997) 8. Lavalley-Fouquet, C., Cabrit, S., Dougados, C. A&A, 356, L41–L44, (2000) 9. Nisini, B., Caratti o Garatti, A., Giannini, T., Lorenzetti, D. A&A, 393, 1035–1051, (2002) 10. Podio, L., Bacciotti, F., Nisini, B., Eisl¨offel, J., Massi, F., Giannini, T., Ray, T.P. A&A, 456, 189–204, (2006)

Laboratory Astrophysics: Episodic Jet Ejections Alberto Marocchino, Jeremy P. Chittenden, Andrea Ciardi, Francisco A. Suzuki-Vidal, and Chantal Stehle

Abstract Recent experiments performed at Imperial College on the pulsed-power MAGPIE facility have successfully shown the formation of magnetically driven radiatively cooled plasma jets formed from radial wire arrays, which are relevant to the study of launching mechanisms in astrophysical jets. The experiments have been now extended to study the episodic ejection (25 ns) and the interaction of jets and magnetic bubbles with an ambient gas. The dynamics of the interaction is investigated through three-dimensional resistive magneto-hydrodynamic simulations using the code GORGON. Comparison with experiments is offered to validate the results. The ablation process as well as current reconnection is described and analyzed. The complex three-dimensional structure and the confinement/collimation effect offered by the magnetic field are investigated. The scenario is modified introducing a background gas (Ar,  6:7  103 kg/m3 ), collimation effects are investigated for the new set-up.

A. Marocchino () Imperial College, Prince Consort Road, London, SW7 2BW, U.K e-mail: [email protected] J.P. Chittenden Imperial College, Prince Consort Road, London, SW7 2BW, U.K e-mail: [email protected]

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1 Introduction: The Role of Laboratory Astrophysics In this work we aim to clarify and possibly explain the YSO-jet launching mechanism and high collimation angle via laboratory experiments. The ability to generate supersonic plasma jets has been successfully demonstrated on laser and Z-pinch facilities [1]. Plasma jets that are generated with pulsed power machines are unique in several aspects. The use of pulsed power machines and ablating wires allow formation of an hydrodynamic continuous jet, rather than the plasma bullets typically generated in laser experiments. Moreover, jets generated with pulsed power machines are in a relevant cooling regime and include a topological relevant magnetic field, a scenario much more difficult to achieve with lasers. The key aspect of pulsed power is the capability to reach particular physical regimes that are relevant to astrophysics. This contribution concentrates on magneto-hydrodynamics (MHD) simulations of laboratory astrophysics experiments performed on the MAGPIE pulsed power generator at Imperial College London. Laboratory experiments have been accurately reproduced with the numerical code GORGON [2, 3]. In the first section multiple mass ejections are investigated as well as the high collimation of the ejections. In the second section the effect of an added background gas is studied.

2 Episodic Jet Emission 2.1 Radial Foil Z-pinch and YSO Multi Bubble Formation The HH 34 and HH 111 are among the best studied of all HH outflows. Since their discovery, detailed spectroscopic imaging and proper motion studies of these objects have been carried out at optical, infrared and radio wavelengths. In the past few years, several authors have studied numerically the effect of the magnetic field on the dynamical evolution of HH objects [6,7]. Since HH jets are characterized by a knotty structure, the magnetocentrifugal [8] mechanism is investigated as a mechanism that repeats itself several times until the young stellar object gets exhausted. Since astrophysical measurements are still too poor to be directly compared both with mathematical models and simulations, laboratory astrophysics might offer a valid help in understanding the evolution and role played by the magnetic field. In order to reproduce in the laboratory the multiple mass ejection the radial foil set-up for pulsed power machines has been used. The foil array consists of 6 m thick, Al foil. The radial foil generates multiple ejections (4 ejections for a single pulse 480 ns), and consequently generates a structure that resembles the typical HH-jet scenario.

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2.2 Radial Foil General Dynamics Initially the current flows from the cathode to the anode via the top part of the Al foil, the foil thus thermally ablates from the top. In particular the foil ablation proceeds as energy is deposited in the foil through Ohmic heating and to a much less degree through thermal conduction. Very early in the current pulse (probably around 20 ns) the foil develops a two-component structure consisting of a dense, cold, and high-resistivity foil core with on top a relatively hot (1–2 eV), low resistive plasma. Because of the marked differences in the resistivity, currents preferentially flow in the hot low resistive, plasma which is accelerated in the axial direction by the Lorentz’s force (Fig. 1 a). The Lorentz’s force is now (200 ns) strong enough to push upward the foil still present in the cathode proximity. The ablated plasma is snowploughed on axis, and there compressed (cartoon b, Fig. 1). The plasma compression on axis is characterized by the production of X-rays. Some of these X-rays would eventually hit the central electrode thus make it ablate. Immediately after the X-ray emission, the central pinch becomes unstable and starts breaking up. The first pinch, and the first cathode ablation, occurs at about 240 ns. While the first jet is now launched, the cathode ablates mass is accelerated on axis by the Lorentz’s force (cartoon d - Fig. 1). The newly formed pinch will undergo a similar evolution, the scenario presented in the cartoons b-c-d is then repeated for the second and subsequent magnetic bubbles. a

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Fig. 1 These four cartoons illustrate the overall evolution for a foil array. Frame (a) refers the to early ablation process. Frame (b) shows the pinch formation and Lorentz’s force direction. Frame (c) shows the cathode ablation and current path swapping. In the last frame, (d), the formation of the second bubble is shown. The current path is also drawn highlighting magnetic field trapping

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2.3 Knots Collimation At approximately 320 ns, all three jets have already been launched, and are moving upward along the z-direction. The magnetic field is fundamental to explain the jet formation and subsequent launching. Moreover the magnetic field plays a fundamental role in the early stage of jet collimation. Around the launched jet the magnetic field stays approximately 100 T for 40 ns, offering a collimation effect. The magnetic field explain the collimation in the early stage, but not far away from the cathode, where it still stays highly collimated. The magnetic field is about 250 T around the just broken pinch, but it decays to a value of 8–10 T just after 80 ns. If synthetic XUV images are produced, the launched jets seems somehow connected, thus the plasma connecting them is highly twisted in between. The same phenomenon happen in astrophysics, X-ray diagnostic shows collimated and connected knots, however that might be a pure projection effect, while the knots are just connected via a highly twisted low dense plasma.

3 Episodic Jet Emission in an Ambient Medium 3.1 Astrophysical Motivation The issue of whether YSO jets with small but extended structure are confined due to some combination of magnetic-fields and surrounding medium is a question that is not yet resolved. Both magnetic-field and background medium play a role, but it is not yet clear in which proportion and in which stages. The magnetic field seems to play a fundamental role in the early stage, while at later stages the surrounding media seems to keep the knots collimated on axis. Prompted by these unresolved issues, laboratory experiments are performed in order to investigate and try to explain these unclear phenomena. The experiments are performed on the MAGPIE facility. The set-up is a modification of the radial foil geometry. A nozzle injects Ar gas at ambient temperature just above the Al foil.

3.2 X-ray Emission Regions Analysis of the X-ray emissions allows us to identify the brighter X-ray source region and to understand emission effects when the enhanced mass on axis, initially, and the jet, later on, interact with the ambient medium. The regions with the highest temperature are clearly the central pinch, the magnetic bubble and the magnetic cavity. However the presence of the background gas enhances the temperature at the shock interface (about 4–5 eV, 1–2 eV higher than the background plasma) and where the top part of the head of the jet starts penetrating the background gas.

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The X-ray emissions originate from localized regions inside the cocoon, these regions correspond to the central pinch, the jet-heads and the top part of the first jet when interacting with the ambient gas. While the first two regions are identical also for a case with no gas, the third region develops only in the case of background gas. The interaction generates a shock, the temperature on the shock interface increases and consequently the region becomes a bright X-ray region.

3.3 Bubble Confinement Generally the bubble morphology does not change with the introduction of a background gas. However a few aspects are different and need to be discussed. The background gas is not offering, as it might be be though, direct confinement. The background Ar gas is dragged away from the initial shock wave expanding in it, consequently the bubble forms and expands in an Al plasma. The background Al density for a foil with an ambient gas is about 15–20% denser than the case with no gas ( bg D 0:07 kg/m3 ). The density increase is due to the fact the the Al expanding plasma is confined within the expanding shock wave. The little mass density difference influences both the expansion velocity that reduces from about 50 km/s (case of no gas) to about 30 km/s, and the bubble edge density is 3 times denser than in the case with no gas (0.85 kg/m3 , case with gas). In the case with gas the central pinch results in being much more stable than the case without, Fig. 2. In particular the background gas is enveloping the head of the jet breaking the radial expansion, making the shape of the jet more elongated and overall denser. The jet is in fact about twice as long as the case with no background gas and just half wide. The overall density is about twice as dense as the case with no Ar gas. The jet collimation is offered by the gas thermal pressure, the gas envelops the jet and keeps it more collimated; at this stage the magnetic field does not offer any collimation effects. The magnetic field does not diffuse into the neutral Ar background gas.

Fig. 2 The image offer a direct density comparison for the case of a standard radial foil (left image) and the same experiment with background gas (right image). The image helps in highlighting the shock that generates at the interface, the higher jet collimation, the different jet densities, and the higher pinch stability in the case with background gas

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References 1. S. V. Lebedev et al. – Magnetic tower outflows from a radial wire array Z-pinch – Monthly Notices of the Royal Astronomical Society, 361:97, (2005) 2. A. Ciardi et al. – The evolution of magnetic tower jets in the laboratory – Physics of Plasmas, 14:056501, (2007) 3. J. P. Chittenden at al. – X-ray generation mechanisms in three-dimensional simulations of wire array Z-pinches – Plasma Physics and Controlled Fusion, 46:B457, (2004) 4. J. Bally et al. – X-rays from the vicinity of the protostar L1551 IRS 5: reflection or fast shocks?– The Astrophysical Journal, 584:843–852, 2003 5. F. Favata et al. – The discovery of an expanding X-ray source in the HH154 protostellar jet– Astronomy & Astrophysics, 2008 6. S. O’Sullivan – Numerical Simulations of Steady and Pulsed non-Adiabatic magnetized jets from young stars– Astronomy and Astrophysics, 363:355–372, (2000) 7. J. M. Stone and P. Hardee – Magnetohydrodynamic Models of Axisymmetric Protostellar Jets – The Astrophysical Journal, 540:192, (2000) 8. D. Lynden-Bell – On why discs generate magnetic towers and collimate jets – Monthly Notices of the Royal Astronomical Society, 341:1360–1372, (2003)

Parameter Study in Disk Jet Systems A Focus on Magnetization Petros Tzeferacos, Attilio Ferrari, Andrea Mignone, Silvano Massaglia, Gianluigi Bodo, and Claudio Zanni

Abstract In the present work we discuss how the strength of magnetic fields determines the characteristics of solutions in models where the collimated outflow and the accretion disk are treated consistently. We perform a complete analysis of the range of magnetic field by non-relativistic 2.5 dimension numerical simulations using the PLUTO code. The main results are that magnetic fields around equipartition with plasma pressure allow steady super-fast-magnetosonic collimated jet solutions; magnetic fields below equipartition correspond to intermittent collimated outflows, while magnetic fields above equipartition lead to sub-alfvenic winds. These results

P. Tzeferacos (), A. Ferrari, A. Mignone, and S. Massaglia Dipartimento di Fisica Generale, Universit´a degli Studi di Torino, Via Giuria 1, 10125 Torino, Italy e-mail: [email protected] A. Ferrari Department of Astronomy and Astrophysics, University of Chicago, USA A. Mignone and G. Bodo Osservatorio Astronomico di Torino, Viale Osservatorio 20, 10025 Pino Torinese, Italy C. Zanni Laboratoire d’Astrophysique de Grenoble, 414 rue de la Piscine, BP 53, 38041 Grenoble, France

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allow to conclude that the configuration proposed by Blandford and Payne to interpret supersonic jets is viable both for equipartition and weaker magnetic fields.

1 Introduction Many objects studied in the context of stellar and extragalactic astronomy fall into the category of accretion ejection systems, such as young stellar objects, microquasars, X-ray binaries, gamma-ray bursts, extended radio galaxies and active galactic nuclei. All these outflow scenarios are associated with inflows onto a central compact object with a strong gravitational pull. Typically a large scale magnetic field is assumed to account for the acceleration, since a thermal or radiation pressure drive alone results inadequate. From the mid seventies Lovelace [13] and Blandford [1] showed that a force-free poloidal field anchored in a Keplerian disk can extract energy and angular momentum creating a Poynting flux jet. This idea was developed in the MHD regime by Blandford and Payne [2], who showed how Poynting jets can transfer their energy and momentum to outflowing matter via a magneto-centrifugal mechanism capable of reaching super-fast-magnetosonic speeds. In these studies the disk was treated as a fixed boundary and no self-consistent treatment of the inflow/outflow dynamics was attempted. More recent works, for example Ferreira and Pelletier [8], Ferreira [9] and Casse and Ferreira [4, 5] have studied analytic “cold” and “warm” steady state outflow solutions and linked them to thin accretion disks . Viscous and resistive transport coefficients have been taken into account allowing for anisotropic magnetic diffusivity between poloidal and toroidal fields; they conclude that super-fast-magnetosonic outflows can be obtained with plasma ˇ values around unity, for a limited range of Prandtl numbers and larger toroidal diffusivity. Numerical simulations allow instead to investigate time-dependent solutions, often still limited to the study of the outflow treating the disk as a boundary condition [19, 16, 11], or referred to the entire disk-jet system but for very short timescales [18,10]. Casse and Keppens [6] have followed the evolution of an accretion-ejection system for longer timescales but without using an energy equation replaced by a simple polytropic equation of state; more recently they updated their work including non-adiabatic effects [7]. In this work we address the problem of the stability of the inflow/outflow dynamics on the basis of compressible MHD numerical simulations, analyzing the effects of the most relevant physical parameters with respect to the possibility of reaching steady state solutions over long time scales of integration. In a previous paper, Zanni et al. [20] have discussed the importance of resistive effects. Here we concentrate on the strength and configuration of the initial magnetic field associated with the accretion disk. Simulations have been performed in a 2.5D resistive framework, utilizing the numerical code PLUTO [14].

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2 Model Description In order to model the accretion ejection system we solve numerically the equations of resistive magnetohydrodynamics, neglecting all transport coefficients but for magnetic resistivity m . The latter is parametrized following the Shakura and Sunyaev [17] paradigm. We allow for anisotropy in the magnetic diffusion through a tensor representation of m and we specify the toroidal diffusion to be three times stronger than the poloidal. Even though we solve the energy equation, we presume that the heating caused by Ohmic dissipation is completely radiated. This is done to avoid a rise in the disk’s entropy, being thus closer to the “cold” analytic solutions of Ferreira [9]. The strength of the magnetic field, in respect to thermal pressure, is defined by the magnetization parameter D B 2 =P to define the strength of the magnetic field. Note that this parameter is used in the analytical formalism [8, 9, 4]. We study four different cases of magnetization, below and above the limits set in the previously mentioned studies. From weak to strong field configurations, cases 1 to 4 correspond to values of  equal to 0.2, 0.6, 2.0 and 6.0 respectively. Only the second case is within the analytical limits. The initial inclination of the magnetic field lines is not studied here and is set to a value, the same for all cases, that satisfies the Blandford and Payne criterion [2]. The initial configuration adopted considers a thin disk rotating at slightly subKeplerian speed with an embedded magnetic field with purely poloidal field lines exiting at some angle from the disk equatorial midplane. The field is initialized through it’s flux function to ensure solenoidality. The profiles of all other primitive variables are derived starting from an equatorial radial self-similarity assumption [2] and imposing a force equilibrium in both radial and vertical directions, as dictated by the following equations:

u2 @˚g @P D 

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In pressure equilibrium with the disk’s surface, we overimpose a static hot corona with a radial stratification Our domain is a rectangular region spanning radially from 0 to 40r0 and vertically from 0 to 120r0 . The temporal integration reaches 63 revolutions of the disk’s inner radius. Uniform resolution of Œ512  1536 is used throughout the domain. Axisymmetry is assumed on the rotation axis whereas equatorial symmetry is supposed for the disk’s midplane. At the upper and right boundaries special care has been shown in order to retain the magnetic field’s solenoidality and to avoid artificial collimation. Since the origin is inside the computational domain, an appropriate

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internal boundary Œ1:0ri n  0:5ri n  is imposed to cope with the singularity of the gravitational potential and the radial self-similar model there. No artificial outflow is allowed to escape that zone.

3 Results-Conclusions In all cases a current carrying outflow is formed from the early stages of the simulations (Fig. 1) The magnetic tension and an outward slip of the footpoint within the disk (as magnetic diffusion increases with VA ) quickly collimate the outflow in high  configurations. The initial inclination of the magnetic field lines amplifies the braking and the angular momentum is quickly redistributed, resulting in a slower rotation of the disk. For a weak magnetic field the collimaton is somewhat inefficient (Fig. 2). This decrease in the field’s curvature has a direct impact both in the ejection efficiency of the system and it’s ability to accelerate the outflowing plasma. Monitoring the outflow/inflow rates at a control volume set in the ejection region we see that there is a clear plateu after the initial transient phase for low  cases. The rate is not sustained for magnetically dominant configurations (Fig. 3). The ejection efficiency found for case 2 is close to the analytical results of Ferreira [9], for a set of parameters close to ours. Only low magnetization configurations are able to cross the Alfv´en and superfast magnetosonic critical surfaces (Fig. 4). The terminal velocities evaluated along field lines for cases 1 and 2 range between 4–5 times the Keplerian velocity of the respective footpoint. High  cases are four to five times slower, with case 4 barely crossing the slow surface at the innermost part. This is attributed both to the field’s topology and the elevated ejection rates. The Poynting flux transformation to kinetic energy for cases 1 and 2 is almost complete (magnetic to kinetic energy

Fig. 1 The poloidal current distribution for the four cases studied, after 6 rotations of the inner radius. A butterfly topology [9] is found in sub-equipartition cases. The ejection region is identified in a range between 1 and 10 ri n . Notice also that the disk’s structure cannot be supported by the thermal pressure in cases 3 and 4 due to the strong magnetic compression. This confirms [8]

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Fig. 2 Sample poloidal magnetic field lines for the four cases studied after 63 rotations. While for low magnetization the expected open topology is achieved, strong fields collimate very prominently. Notice also the formation of a slow and dense outflow in the outer part of the disk for cases 1 and 2, as the pressure gradient wins over magnetic compression

Fig. 3 Mass ejection to accretion rates as a function of time, taking into consideration the counter jet. The notation used is: Case 1 dotted, case 2 solid, case 3 dashed and case 4 dot-dashed line. Cases 3 and 4 decay into smaller values though as the ejection is not sustained. These two cases do not reach a steady configuration in contrast to the clear plateu reached for low magnetization configurations

Fig. 4 Snapshots, after 35 rotations, of the poloidal velocity over the fast magnetosonic. The line indicates the Alfv´enic surface and the poloidal velocity vectors are superimposed. The lenght of the arrows is normalized to the maximum poloidal speed reached in case 1. The position of the Alfv´enic surface is closer to the disk as  decreases

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flux ratios 102 –101 at the upper end of the domain). Magnetically dominant configurations, on the other hand, seem defficient in accelerating the outflowing plasma with ratios of 1–3. We have also shown that for weakly magnetized disks the magnetic pressure drive wins over the centrifugal one in smaller altitudes, while the situation reverses for high  configurations where we only have co-rotation; Nonetheless, the centrifugal acceleration is not efficient in above equipartition cases since the field’s inclination and the slow rotation do not favour the mechanism. The clear plateau in the ejection efficiency and the preservation of the MHD invariants (not shown here) advocate in favour of the quasi stationarity of the solutions for magnetization below equipartition. It is also clear that no such claim can be made for the strongly magnetized cases. This statement applies only to the innermost part of the magnetized disk, since the revolutions of the outer part are not enough to reach a quasi stationary state. Acknowledgements PT would like to thank N. Vlahakis and K. Tsinganos for many fruitful discussions and comments. We would also like to thank J. Ferreira, S. Cabrit and A. K¨onigl for all the valuable suggestions. AF thanks R. Rosner and D. Lamb for hospitality and support at the University of Chicago. This work was supported in part by the U.S. DoE under grant No. B523820 to the Center of Astrophysical Thermonuclear Flashes at the University of Chicago and the EU Marie Curie Research Training Network JETSET under contract MRTN-CT-2004-005592.

References 1. Blandford, R.D. 1976, MNRAS, 176, 465 2. Blandford, R.D., and Payne, D.G. 1982, MNRAS, 199, 883 3. Balasara, D.S., Spicer, S.D. 1999, J. Comput. Phys, 149, 270 4. Casse, F. and Ferreira, J. 2000a, A&A, 353, 1115 5. Casse, F. and Ferreira, J. 2000b, A&A, 361, 1178 6. Casse, F. and Keppens, R. 2002, ApJ, 581, 988 7. Casse, F. and Keppens, R. 2004, ApJ, 601, 90 8. Ferreira, J. and Pelletier, G. 1995, A&A, 295, 807 9. Ferreira, J. 1997, A&A, 319, 340 10. Kato, S.X., Kudoh, T. and Shibata, K. 2002, ApJ, 565, 1035 11. Krasnopolsky, R., Li, Z.Y. and Blandford, R.D., 1999, ApJ, 526, 63 12. Londrillo, P. and Del Zanna, L. 2004, J. Comput. Phys, 195, 17 13. Lovelace, R.V.E. 1976, Nature, 262, 649 14. Mignone, A., Bodo, G., Massaglia, S., Matsakos, T., Tesileanu, O., Zanni, C., and Ferrari, A. 2007, ApJS, 170, 228 15. Ogilvie, G.I., and Livio, M. 2001, ApJ, 553, 158 16. Ouyed, R., and Pudritz, R.E. 1997, ApJ, 482, 712 and 484, 794 17. Shakura, N.I. and Sunyaev, R.A. 1973, A&A, 24, 337 18. Uchida, Y., and Shibata, K. 1985, PASJ, 37, 515 19. Ustyugova, G.V., Koldoba, A.V., Romanova, M.M., Chechetkin, V.M., and Lovelace, R.V.E. 1999, ApJ, 516, 221 20. Zanni, C., Ferrari, A., Rosner, R., and Massaglia, S. 2007, A&A, 469, 811

Part VIII

Posters

Early Stage Development of the Jetset Database Periklis Rammos, Emma T. Whelan, Jos´e Gracia, Stephane Dudzinski, and Philippe Grange

We are developing an online Database. Similar databases exist, like the The Youngest Protostars, by Dirk Froebrich, from which we can take some ideas. This database features separate tables for each object, few data for each observation and a link to the corresponding publication. Our current goals are to have user-entered data, a query-able database and separate tables for each object (source and outflow). The fields of the database are defined by Work Package managers. The data are stored on the server, as a text file for now. The user interface is an html website, and the communication between the two is done with php scripts, activated by the user (Fig. 1). For now only two functions are available: Entering new data (Fig. 2), and viewing data (Fig. 3).

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P. Rammos (), E.T. Whelan, J. Gracia, S. Dudzinski, and P. Grange Dublin Institute for Advanced Studies, e-mail: [email protected]

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Fig. 2 Screen shot of the window for entering new data

Fig. 3 Screen shot of the window for viewing data

For the future, we plan to replace the text file with a proper MySQL database, which will allow us to use queries. Also we intend to specify levels of access, such as who can view and/or edit what data. Another interesting feature would be including images or spectra that the user can upload. Finally, we intend to explore alternative implementations such as SAADA or Ruby on Rails. Acknowledgements The present work was supported by the European Community’s Marie Curie Actions - Human Resource and Mobility within the JETSET (Jet Simulations, Experiments and Theory) network under contract MRTN-CT-2004 005592.

Shaping Planetary Nebulae by Jets Muhammad Akashi

Abstract We conduct 2D axisymmetrical hydrodynamical simulations to investigate the interaction of a collimated fast wind (CFW; wide jets) with a spherical AGB wind. The code includes radiative cooling. We find that the shape of the planetary nebula (PN) is sensitive to the exact mass loss history of the AGB wind, and the opening angle of the CFW. Some typical PN morphologies are obtained, but many other observed morphologies seem to require more ingredients than what we assume in our present simulations, e.g., equatorial AGB wind, and ionization and fast wind during the PN phase. The hot bipolar bubble formed by the jets is an X-ray source.

1 Numerical Simulations Our simulations were performed using Virginia Hydrodynamics-I (VH-1), a high resolution multidimensional astrophysical hydrodynamics code developed by John Blondin and co-workers (Stevens et al., 1992; [2, 1]). We have added radiative cooling to the code at all temperatures T > Tmin , where the gas temperature is bound from below at Tmin D 300–1000 K. Radiative cooling is carefully treated near contact discontinuities, where large temperature gradients can lead to unphysical results. We simulate axisymmetrical morphologies. This allows us to use axisymmetrical grid, and to simulate one quarter of the meridional plane. There are 208 grid points in the azimuthal ( ) direction of this one quarter and 208 grid points in the radial direction. The radial size of the grid points increases with radius. In these simulations the grid extends from 1015 cm to 4  1017 cm. Before the jet is launched (t D 0), the grid is filled with slow wind having a speed v1 D 10 km s1 and mass loss rate MP 1  105 Mˇ yr1 , with small variations

M. Akashi () Department of Physics, Technion – Israel Institute of Technology e-mail: [email protected]

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Fig. 1 Left: Density plot at t D 2000 yr (after jet launching starts) for a case where the slow wind mass loss rate has increased by a factor of five to MP 1 D 5  105 Mˇ yr1 950 yr before jet starts. The vertical axis is the symmetry axis and the horizontal axis is in the equatorial plane. Slow spherical wind: at t D 0 v1 D 10 km s1 , MP 1 D 5  105 Mˇ yr1 for r < 3  1016 cm, and MP 1 D 105 Mˇ yr1 for r > 3  1016 cm. The jet has a half opening angle ˛ D 20ı , v2 D 600 km s1 , and MP 2 D 2  108 Mˇ yr1 . Arrows indicate flow direction: v > 200 km s1 (long arrow), 20 < v 200 km s1 (medium arrow), and v 20 km s1 (short arrow). This general case can account for PNs such as NGC 6886 or NGC 3918. Right: Density plot at t D 530 yr for another case. Slow spherical wind: v1 D 10 km s1 , MP 1 D 5  106 Mˇ yr1 . The jet has a half opening angle ˛ D 40ı , v2 D 600 km s1 , and MP 2 D 2:5  106 Mˇ yr1 . This general case can account for some morphological features in He 2-104, such as the instabilities in the outer lobes

between the runs. We launch a collimated fast wind from the first 20 zones attached to the inner boundary of the grid. The jet is uniformly ejected within an angle (half opening angle) ˛ (0    ˛). For numerical reasons a weak slow wind is injected in the sector ˛ <   90ı ( D 0 along the symmetry axis – vertical in the Fig. 1).

2 Results and Summary We try to explain PN shaping and formation processes by simulating jets interacting with the AGB wind. The parameter space for these types of flows is huge. We run more than 200 simulations, in which we tried many values of the CFW (jets) mass loss rate, velocity, and opening angle. Another parameter that was used is the mass loss rate history of the AGB wind. For example, in a model presented here (Fig. 2), we assume that some period t before the beginning of the jet-launching phase the slow wind mass loss rate increased by some factor k; e.g., in the run presented in Fig. 2 t D 950 yr and

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Fig. 2 Left: Density plot at t D 1300 yr for a case with a slow wind of v1 D 10 km s1 , and mass loss rate of MP 1 D 3  105 Mˇ yr1 for r < 3  1016 cm, and MP 1 D 5  106 Mˇ yr1 for r > 3  1016 cm at t D 0. The jet has a half opening angle ˛ D 30ı , v2 D 600 km s1 , and MP 2 D 107 Mˇ yr1 . The initial temperature was 1000 K, and the temperature is limited from below at 1000 K. Right: Density plot at t D 1560 yr for a case with a slow wind of v1 D 10 km s1 , and mass loss rate of MP 1 D 3  105 Mˇ yr1 for r < 3  1016 cm, and MP 1 D 5  106 Mˇ yr1 for r > 3  1016 cm at t D 0. The jet has a half opening angle ˛ D 30ı , v2 D 600 km s1 , and MP 2 D 107 Mˇ yr1 . The initial temperature was 10,000 K, and the temperature is limited from below at 1,000 K. The general shape is similar to MZ 3. The arrows are just as in Fig. 1. Note that left and right are the same beside the initial temperature

k D 6. We found that the mass loss history of the AGB wind plays a significant role in the shaping process. In some runs the AGB wind initial temperature was set to 10;000 K instead of 1;000 K, with a noticeable influences on the final shape of the PN (compare the two panels in Fig. 2). Taking the AGB wind to be at 10;000 K is relevant to system where the accreting companion is a WD. If the accretion rate is high enough the WD sustains a constant nuclear burning, and it is very hot and luminous, as super-soft X-ray sources (e.g., [3]). Such an accreting WD might maintain the entire AGB wind ionized. We were able to reproduce some physical properties of PNs, but not all desirable properties. It is very difficult, and might be impossible, to get the entire range of properties with our limited numerical code. We must add more ingredients to the code. Examples are pulsed jets, precessing jets, dense equatorial outflow and ionization fronts at later times. We presented results of four numerical runs that migh match real PNs. Our main results are as follows: 1. We show that different PNs morphologies could be strongly dependent on the AGB wind mass loss rate history.

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2. We have an evident that the final morphology is very sensitive also to the assumed initial temperature. For Ti D 1; 000 K and Ti D 10; 000 K we get two different evolutions of the planetary nebula (Fig. 2). 3. The dense finger along the symmetry axis (in Fig. 2, top) result from instabilities. In reality, it would be more extended and not so narrow. It is forced to the symmetry axis by the numerical code. 4. In one case we obtain low density finger (in Fig. 2) protruding to the upper right form a torus. However, it is also a result of an instability, and we expect that in real systems there will be several such fingers. In particular, if there is a departure from axisymmetry due to the orbital motion, we would expect the fingers to be similar but not identical in the two sides of the equatorial plane. 5. In all cases that we run we saw that there is a strong dependence on the half opening angle of the jet. This will be the focus of a future paper.

References 1. Blondin J.M., 1994, The VH- 1 Users Guide. Univ. Virginia., VA, USA 2. Blondin J.M., Kallman T.R., Fryxell B.A., Taam R.E., 1990, ApJ, 356, 591 3. Starrfield S., Timmes F.X., Hix W.R., Sion E.M., Sparks W. M., Dwyer S.J., 2004, ApJ, 612, L53

New Herbig-Haro Objects in the Gulf of Mexico Tina Armond, Bo Reipurth, and Luiz Paulo R. Vaz

Abstract We present the result of a survey for Herbig-Haro objects in the “Gulf of Mexico”, a part of the dark cloud L935. Images taken through H˛ and [S II] filters at SUBARU Suprime Cam covering an area of 340  270 reveal 35 new Herbig-Haro objects in this active star forming region, little studied so far.

T. Armond () Centro de Astrof´ısica da Universidade do Porto, Rua das Estrelas 4150-762 Porto, Portugal e-mail: [email protected] B. Reipurth Institute for Astronomy, University of Hawaii, 640 N. Aohoku Place, Hilo, HI 96720, USA e-mail: [email protected] L.P. R. Vaz Depto. de F´ısica, ICEx, UFMG, CP 702, 30123-970 Belo Horizonte, MG, Brazil e-mail: [email protected]

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1 Introduction The “Gulf of Mexico” is part of the dark cloud L935, a dust lane between the North America (NGC7000) and the Pelican Nebulae (IC5070). All of them are part of a single HII region, W80, at a distance of about 550 pc at the nearest part. Signs of active star formation have already been found in the Gulf of Mexico, around the T Tauri stars LkH˛ 185–189 [4]. Herbig-Haro objects, H˛ emission-line stars and infrared excess stars were found in a previous optical and near-infrared survey covering an 70  140 area [1,2]. This survey was made with the University of Hawaii 2.2 m telescope on Mauna Kea in 2002. New deeper and wider images obtained with the SUBARU Suprime Cam cover the whole gulf (340  270 ). The images taken through H˛ and [S II] filters show more Herbig-Haro objects outside the region previously surveyed and increase the amount of details of the known objects.

2 Herbig-Haro Objects In our first observations in 2002, we have identified 28 Herbig-Haro objects in the region. Most of them appear only in the [S II] images; the few that are detected in H˛ are stronger in [S II], showing that these are low excitation shocked jet material, and not photo-ionized nebulae. Some of the flows were also detected in near-infrared H2 images. Broadband I images were also checked to prevent reflection nebulae to be identified as a Herbig-Haro object. The new SUBARU [S II] and H˛ images reveal more details of the known Herbig-Haro objects. As the images cover a larger region, they show seven new HH objects and also a couple of nebulous stars. This is a sign that star formation might be occurring in the whole L935 dark cloud, possibly triggered by the expansion of the shock wave from the exciting source(s) through the molecular shell [3]. In our previous survey, 35 H˛ emission-line stars were identified, most of them near the optical cluster of LkH˛ 188. Previously obtained near-infrared images and recently released SPITZER images reveal a chain of embedded sources from the optical cluster to a larger structure, already seen through the MSX infrared telescope, as well as some embedded outflows. It is possible to relate many of the Herbig-Haro objects to probable sources, visible in the optical and/or the near-infrared. But only a more detailed study can certainly define this. We conclude that star formation is much more active in that area than previously suspected. Spectroscopic observations together with the infrared images from SPITZER can clarify the nature of the sources driving the flows. The optical cluster of H˛ emission stars is a fraction of the total number of stars being formed there. The presence of Herbig-Haro flows, the reddened sources shown in JHK and the recent SPITZER images show that this is a large site of widespread star formation, still partially hidden behind the dark cloud, and deserves a profound study.

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Acknowledgements Thanks for the financial support from CNPq/Brazil (Processes 200430/ 2001-7 and 201958/2007-4), FCT/Portugal (Project PTDC/CTE-AST/65971/2006) and the Local Organizing Committee of this Conference.

References 1. Armond, T., Reipurth, B., Vaz, L. P.: In: Bol. da Soc. Astronˆomica Brasileira, 23, 92 (2003) 2. Armond, T., Reipurth, B., Bally, J., Aspin, C.: submitted to A&A (2009) 3. Bally, J., Scoville, N.Z.: ApJ, 239, 121 (1980) 4. Herbig, G.H.: ApJ, 128, 259 (1958)

Launching Jets from MRI-driven Accretion Discs Steffen Brinkmann and Max Camenzind

Abstract Jets are launched from a wide variety of compact objects ranging from young stellar objects to stellar black holes and AGN. The dynamics of accreting systems are mainly determined by the gravity of the central object, the rotation of the accreting matter and the magnetic fields. It is widely accepted that the emergence of turbulence and the transport of angular momentum is dominated by magnetic turbulence rooted in the magnetorotational instability (MRI), which is a powerful source of angular momentum transport [1, 2]. In order to study the dynamics of accretion-ejection systems, we perform global direct numerical simulations with varying magnetic configurations.

1 Numerics and Setup The MHD code used for the simulations is PLUTO [5], a versatile collection of Godunov-type solvers for relativistic and non-relativistic HD and MHD. We apply an HLLD-algorithm with linear reconstruction and a Runge-Kutta timestepping

S. Brinkmann () and M. Camenzind Landessternwarte (ZAH), Universit¨at Heidelberg, K¨onigstuhl 12, 69117 Heidelberg, Germany e-mail: [email protected] K. Tsinganos et al. (eds.), Protostellar Jets in Context, Astrophysics and Space Science Proceedings, DOI 10.1007/978-3-642-00576-3 65, c Springer-Verlag Berlin Heidelberg 2009 

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scheme of 2nd order. The simulations were carried out on a 2.5 dimensional (i.e. axisymmetric), spherical, static grid, resolving the radial and poloidal directions with 512  256 grid cells. The simulated region extends from 4 to 90 gravitational radii. A pseudo-Newtonian potential ˚ D GM=.r  2Rg / [6] emulates the general relativistic effects of a non-rotating Black Hole and yields an analytical solution for the torus which is hydrodynamically stable [4]. We place a weak, poloidal magnetic field entirely inside the torus, setting the mean plasma-ˇ D 103 . We tested 4 different setups of the magnetic field, a dipole, a quadrupole, a sextupole and an octupole. The loops of the multipole setups were vertically stacked with alternating helicity.

2 Results Accretion starts immediately after the onset of the simulation and reaches saturation at a mean value that depends on the initial setup of the magnetic field. Higher orders of field geometry accrete faster and more efficient, because angular momentum is transported radially by the r--component of the stress-tensor, [3] and for higher order of the magnetic field, more magnetic energy is situated in the radial component, which causes enhanced transport of angular momentum. Furthermore, the wavelength of the initial magnetic perturbations is closer to the fastest growing mode of MRI for the higher order setups. As soon as accretion sets in, fast jet-like outflows form in both hemispheres (see Fig. 1). They reach velocities close to the speed of light. The jet-like outflows rotate not only in the same sense of the disc rotation, but remarkably also in the other direction. We found two patterns of retrograde rotation (see Fig. 1): (1) The outflow builds up layers which rotate in opposite directions. (2) The outflow of both hemispheres rotate in opposite directions.

Fig. 1 Toroidal velocity. Colours show the phi -component of the velocity in units of c for different setups at different times. From left to right: dipole at 1000 lct , dipole at 1600 lct , sextupole at 400 lct and sextupole at 1000 lct . A color version of this plot can be found in the electronic version of this book and in Appendix A (Fig. A.33)

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Both types of counter-rotation were observed in different setups at several times, as the unstable dynamical behaviour of the outflow also reflects in the rotation pattern. We assume that the azimuthal acceleration is in some way caused by the magnetic field, which dominates the dynamics of the inner part of the accretion flow. The details of this process however need further investigation both analytically and by simulations.

References 1. Balbus, S. A. & Hawley, J. F. 1991, ApJ, 376, 214 2. Balbus, S. A. & Hawley, J. F. 1998, Reviews of Modern Physics, 70, 1 3. Balbus, S. A. & Hawley, J. F. 2003, Lecture Notes in Physics, 614, 329348 4. Igumenshchev, I. V., Chen, X., & Abramowicz, M. A. 1996, MNRAS, 278, 236 5. Mignone, A., Bodo, G., Massaglia, S., et al. 2007, ApJS, 170, 228 6. Paczynsky, B. & Wiita, P. J. 1980, A&A, 88, 23

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Properties of Jet Emitting Discs C´eline Combet and Jonathan Ferreira

1 Context When a star forms, both disc accretion onto the central object and bipolar ejections of material are observed. It is now widely accepted that these two phenomena are coupled. Therefore, the presence of jets must affect the underlying disc structure. Here, we calculate the latter taking into account the presence of the jets. The results for these jet emitting discs (JED) should apply to the inner regions of accretion discs (r . 1 AU) where jets are believed to originate from [5]. They are to be compared to the structure of the so-called standard accretion disc (SAD) model, where no jet is present. This work is published in [2], where more details and references can be found.

C. Combet () Dept. of Physics and Astronomy, University of Leicester, LE17RH, Leicester, United Kingdom e-mail: [email protected] J. Ferreira Laboratoire d’Astrophysique de Grenonble, BP53, 38041 Grenoble C´edex 9, France e-mail: [email protected]

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2 Including the Jets in the Disc Structure Calculation The analytical treatment used in this work is very similar to that of the early studies of steady-state SADs (e.g., [7]). The extra ingredient is the use the Magnetised Accretion-Ejection Structures model properties (MAES, [4]) that allow us to include the effects of the jet torque on the disc. Indeed, it has been shown that for such steady-state accretion-ejection structures to exist, the jet torque should be much larger than the viscous torque, thus departing from the well-studied SAD. The heating term of the disc can be written in a generic form as QC D f  .GM? MP a =8 r 3 / where f represents the fraction of gravitational potential energy that contributes to heat the gas, M? the mass of the central object and MP a the mass accretion rate. For standard discs, f D fSAD D 1. For a JED, however, most of the gravitational energy escapes with the jet and only a small fraction contributes to the viscous heating. In the MAES model, it has been shown [1] that f D fJED  " 1 where " is the disc aspect ratio. For the cooling, we assume that the disc is optically thick and radiates like a black body. Equating heating and cooling terms allows us to analytically determine the radial variations of all the disc physical quantities.

3 Summary of Results and Perspective The main findings are that a JED is thinner, lighter and cooler that a SAD displaying the same accretion rate [2]. In particular, the surface density of a JED is about two orders of magnitude smaller than that of its equivalent SAD, and its radial velocity much larger. The picture is then the following: a inner light thin disc launching the jet is supplemented by a standard structure at larger radii. The implications of this structural change of the inner disc are manifold:  We have shown that the presence of a JED, despite its small radial extension,

should affect the spectral energy distribution (SED) of the disc. A more sophisticated SED calculation is nevertheless required to quantify this result.  It has been recently showed by [6] that a steep density gradient could halt the type I migration of planetary embryos. The transition from an inner JED to an outer SAD would naturally produce such conditions and stop planetary embryos at the JED outer radius. Numerical simulations are now needed to explore this point further.  The JED reduced thickness and density suggest that it is more prone to ionisation (X-rays from the central object and cosmic rays). This questions the existence of dead zones (neutral regions decoupled from the magnetic field, important for planet formation) in these discs. Preliminary work suggests that the location of the X-ray source (on the disc midplane or slightly above) drastically affects the ionisation rate in JEDs. This is still work in progress [3].

Properties of Jet Emitting Discs

References 1. Casse, F. & Ferreira, J., A&A, 361, 1178 (2000) 2. Combet, C. & Ferreira, J., A&A, 479, 481 (2008) 3. Combet, C. & Ferreira, J., in preparation (2009) 4. Ferreira, J., 1997, A&A, 319, 340 (1997) 5. Ferreira, J., Dougados, C., & Cabrit, S., A&A, 453, 785 (2006) 6. Masset, F. S., Morbidelli, A., Crida, A., Ferreira, J., ApJ, 642, 478 (2006) 7. Pringle, J., Ann. Rev. Astron. Astroph., 19, 137 (1981)

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The H2 Velocity Field of Inner Knots in HH 212 Serge Correia, Hans Zinnecker, Stephen Ridgway, and Mark McCaughrean

Abstract High resolution R  50 000 long-slit spectroscopy of the inner knots of the highly symmetrical protostellar outflow HH 212 was obtained in the 1-0 S(1) line of H2 at 2.12 m with a spatial resolution of 0.45 arcsec using Phoenix at Gemini South. At the resulting velocity resolution of 6 km s1 multiple slit orientated observations of the northern first knot NK1 clearly show double-peaked line profiles consistent with either a radiative bow shock or dual forward and reverse (or jet) shocks. In contrast, the velocity distribution of the southern first knot SK1 remains single-peaked suggesting a significantly smaller jet velocity and possibly a different density variation in the jet pulses in the southern flow compared to the northern flow. Comparison with an analytical semi-empirical model of bow shock emission allows us to constrain parameters such as the bow inclination to the line of sight, the bow shock and jet velocities for each flow. Although a few features are not reproduced by this model, it confirms the presence of several dynamical and kinematical asymmetries between opposite sides of the HH 212 bipolar jet. The position-velocity diagrams of both knots exhibit complex dynamics broadly consistent with emission from a bow shock and/or jet shock which does not exclude jet rotation, although a clear signature of jet rotation in HH 212 is missing. Alternative interpretations for the variation of radial velocity across these knots, such as a variation in the jet orientation, as well as for the velocity asymmetries between the flows are also considered (more details in S. Correia, H. Zinnecker, S. Ridgway, and M. McCaughrean, A&A, in print).

S. Correia () and H. Zinnecker AIP, An der Sternwarte 16, 14482 Potsdam, Germany e-mail: [email protected]; [email protected] S. Ridgway NOAO, PO Box 26732, Tucson, AZ 8526, USA e-mail: [email protected] M. McCaughrean Astrophysics Group, School of Physics, University of Exeter, Exeter EX4 4QL, UK e-mail: [email protected]

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Knot Velocity Structures The position-velocity (PV) diagrams (not represented here) are qualitatively consistent with a bow shock model but the following features are not reproduced : - One-sided tail (blue-shifted for NK1, red-shifted for SK1) which likely arise from a dual forward and reverse shock structure, i.e. a combination of a jet shock (or reverse shock or Mach disk) and a bow shock. - Increasing velocity towards the rear of the bow. The absence of “spur” structure in the PV diagrams may be indicative of a C-type shock and/or an increase of entrained material in the bow wings.

Transverse Velocity Gradients and Jet Rotation The data could be consistent with jet rotation in HH 212 with upper limits of 1 km s1 but there is no clear signature. The sense of rotation would be in agreement with that of the protostellar core in both lobes.

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Asymmetries in the Jet : Velocity & Orientation Our model yields bow shocks velocities with respect to the LSR of 110 km s1 for NK1 and 55 km s1 for SK1 (bow shock velocities with respect to the preshocked material are respectively 55 km s1 and 31 km s1 ). There is therefore a factor two difference in velocity between the two lobes. The reason for such intrinsic velocity asymmetry is unclear, but its correlation with jet width/collimation which seems to apply to HH 212 and to other jets could give some insights. Our model is also consistent with a somewhat different jet orientation between the lobes with inclinations to the line of sight of 83ı and 87ı for NK1 and SK1, respectively.

Magnetic Fields in Low-Mass Star Forming Regions: Alignment to Jets/Outflows? Rachel L. Curran and Antonio Chrysostomou

Abstract We present results of an analysis of the alignment of the large scale magnetic field direction to the jet/outflow direction for the largest sample (to date) of low-mass star-forming regions observed using submillimetre polarimetry. Cumulative distribution functions of the difference between the mean position angles of the field vectors and the jet/outflow axis reveal no correlation. However, visual inspection of the maps reveal alignment in 4 out of the 8 low-mass regions.

1 Introduction: Magnetic Fields and Jets/Outflows Current theories for the launching and collimation of jets/outflows from protostars rely on the magnetic field being 60ı to the disk plane for launching the jet, and then toroidal in order to collimate the jet [1]. Observations of the direction of the

R.L. Curran () Osservatorio Astronomico di Palermo, Piazza del Parlamento, 1, 90134 Palermo, Italy e-mail: [email protected] A. Chrysostomou Joint Astronomy Centre, 660 N. A’ohoku Place, University Park, Hilo, Hawaii 96720, U.S.A. e-mail: [email protected]

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field, both through the disk and along the jet, will help to place constraints on such models. Whilst it is not yet possible to observe the magnetic field on such small scales, it is interesting to see if the larger scale magnetic fields show any relation to the jet/outflow directions. The magnetic field directions were derived from observations using the Submillimetre Common User Bolometer Array (SCUBA; [5]), in conjunction with the polarimeter [4], on the James Clerk Maxwell Telescope (JCMT) in Hawaii 1 (see [3] for a discription of the data).

2 Analysis and Conclusions Previous studies of T-Tauri stars [6] found no relation between the magnetic field and the sample of T-Tauri stars as a whole (both with and without jets), although there was a relation between the magnetic field direction and the T-Tauri stars with jets. Such a relation between the magnetic field on these scales and the jet/outflow direction may indicate the large scale magnetic field influences the outflow direction. The weighted mean of the B-vector position angles from our JCMT data were calculated and compared to the position angle of the jet/outflow axis. The magnetic field vectors are not true (i.e. undirectional) vectors. They have a 180ı ambiguity and so range 0ı –180ı . Figure 1 shows the cumulative distribution function for our sample and the function expected for a randomly orientated sample. A

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1 The James Clerk Maxwell Telescope is operated by The Joint Astronomy Centre on behalf of the Science and Technology Facilities Council of the United Kingdom, the Netherlands Organisation for Scientific Research, and the National Research Council of Canada.

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Kolmogorov-Smirnov test reveals that there is a 98.8% chance that the low-mass sample is randomly orientated, compared to 84.9% for the high-mass sample [2]. This method of analysis may not be the best way of summarising the magnetic field direction within such regions, as we observe several sources in which there are abrupt changes in direction of the magnetic field. For such regions, modal averages (in reasonable sized bins) may be more representative. Also, visual inspection of the polarimetry maps is a good way of analysing the alignment. Acknowledgements R.L.C. acknowledges funding from SFI under grant number 04/BRG/P02741, and from the Marie Curie Fellowship Contract No. MTKD-CT-2005-029768 of the project “Young stellar objects, their surroundings and jets: advanced observational and MHD studies”.

References 1. Blandford, R.D., & Payne, D.G., 1982, MNRAS, 199, 883 2. Curran, R.L. & Chrysostomou, A. 2007, MNRAS, 382, 699 3. Curran, R.L. & Chrysostomou, A. 2008, in prep. 4. Greaves, J.S., et al. 2003, MNRAS, 340, 353 5. Holland, W.S., et al. 1999, MNRAS, 303, 659 6. M´enard, F. & Duchˆene, G., 2004, A&A, 425, 973

Interacting Knots in Jets: Simulations vs. Observations Fabio De Colle and Alessio Caratti o Garatti

Abstract Recent observations of stellar jets show time variability in the emission and velocity of the knots along the jets, on time scales of a few years [1, 2]. We compare results obtained by numerical simulations of protostellar jets with observations. The simulations include a detailed treatment of the evolution of ionic/atomic species, and the variability is assumed to be due to the interaction between knots with different shock speeds and densities. Method & Initial conditions To study the knot interaction, we run a set of two-dimensional axisymmetric simulations using the mescal code [3], changing density, shock velocities and mass of the knots. The codes include rate equations for the calculation of the ionization and recombination of a set of 17 atomic/ionic species [4]. We produce synthetic [SII] and

F. De Colle () Dublin Institute for Advanced Studies, 31 Fitzwilliam Place, Dublin 2, Ireland e-mail: [email protected] A. Caratti o Garatti Th¨uringer Landessternwarte Tautenburg, Sternwarte 5 07778 Tautenburg, Germany e-mail: [email protected]

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H˛ emission maps, properly convolved to take into account projection and instrumental effects. From the resulting synthetic maps, we measure the time variability of the knots, in particular, their flux and velocity variations. Finally, our results are compared with observations. We run 5 models of axisymmetric jet. Model m0 : the jet is injected with a density njet D 103 cm3 , a radius Rjet D 3  1015 cm, and moves into an environment with density nenv D 100 cm3 . The initial temperature is T D 103 K everywhere. The injection velocity varies sinusoidally with time, with period D 15 yr, an average velocity v0 D 200 km s1 and a velocity variation amplitude v different from zero for 6 < t = < 6:5 (v D 50 km s1 ) and 7 < t = < 7:5 (v D 180 km s1 ). The computational domain is (Lz , Lr ) D (1018 ,1.25 1017 ) cm, resolved by 2048  256 grid points. The other models differ from model m0 in the following. Model m1 : njet D 103 cm3 and nenv D 102 cm3 . Model m2 : njet D 500 cm3 . Model m3 : v D 230 km s1 for 7 < t = < 7:5. Model m4 : the mass of the impacting knot is 5 times the mass of the target. Results & Discussion Numerical simulations of stellar knot interaction reproduce at least qualitatively the observed variations in luminosity (Fig. 1) and velocity:  Both observations and simulations show a sudden increase and a subsequent slow

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 The luminosity variations are of the same order (a factor of 2–10) in observations

and simulations (m0 , m2 , and m3 ), with similar time scales (10–30 years, in m0 , m2 , and m3 ).  The simulations reproduce the observed accelerations and decelerations. The velocity variation time scales in the simulations are 2–3 times longer than the observed, while the variation amplitude is similar.  The models produce different effects on the duration and intensity of the variability. The model that best reproduces the observation is m3 . Acknowledgements The present work was supported in part by the European Community’s Marie Curie Actions Human Resource and Mobility within the JETSET (Jet Simulations, Experiments and Theory) network under constract MRTN-CT-2004-005592. The authors acknowledge the Irish Centre for High-End Computing for the provision of computational facilities and support.

References 1. Bonito R., Fridlund C. V. M., Favata F. et al. 2008, A&A 484, 389 2. Caratti o Garatti A., Eisl¨offel J., Froebrich D. et al., in preparation 3. De Colle F., PhD Thesis (2005) UNAM Mexico 4. Raga A. C., De Colle F., Kajdic P. et al. 2007, A&A 465, 879

Wide Field JCMT HARP-B CO(3-2) Mapping of the Serpens Cloud Core Odysseas Dionatos, Brunella Nisini, Teresa Giannini, Claudio Codella, John Richer, and Mario Tafalla

1 Abstract We present 12 CO(3-2) JCMT–HARP-B map covering 46000  23000 on the Serpens cloud core (a detail is presented Fig.1, left panel). This homogenous dataset presents an ideal environment to probe the mass accretion and ejection properties through the study of outflow properties. CO emission traces the swept-up ambient gas that has been entrained by the protostellar jet/wind; backtracking the outflow lobes, we can associate the outflow activity to individual protostellar cores.

2 Outflow activity, Momentum flux vs bolometric luminosity On this dataset, we have calculated basic outflow properties such as the column density, the total mass, momentum, kinetic energy and momentum flux following the method descibed by [2], assuming two values for the optical depth ( D 0; 1) and an arbitrary 45ı inclination angle. We have employed Spitzer photometric data [7] along with continuum sub-mm [3], mm[5, 6] and cm[4] observations to reconstruct the SED of the sources driving the outflows, and derive their bolometric luminosity.

O. Dionatos (), B. Nisini, and T. Giannini INAF - Osservatorio Astronomico di Roma, Italy e-mail: [email protected]; [email protected]; [email protected] C. Codella Istituto di Radioastronomia, Sezione di Firenze, Italy e-mail: [email protected] J. Richer Cavendish Laboratory, Cambridge, UK e-mail: [email protected] M. Tafalla Observatorio Astronmico Nacional, Madrid, Spain e-mail: [email protected]

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Fig. 1 (left:) High velocity CO J D 3-2 outflow map around SMM1 (detail of the area observed) superimposed on an IRAC 8 m Spitzer image. Solid (blue) contours delineate blueshifted gas (integrated over 30 km s1 < VLSR < 0km s1 , while dashed (red) contours delineate redshifted gas (integrated over 18 km1 < VLSR < 1 ). Filled and dashed contours start at 0.8 K km1 with an 1 K km1 increment. (right:) Mass flux versus bolometric luminosity diagram; details for the various points presented can be found in the text. A color version of this plot can be found in the electronic version of this book and in Appendix A (Fig. A.35).

These result to be up to 10 smaller than previous estimations based on IRAS data, a fact that can be attributed to the superior spatial resolution of the Spitzer in comparison to IRAS. In Fig.1 we plot the momentum flux derived for our sample of outflows as a function of the source bolometric luminosity. On the same plot the correlation found from [1] on Class I sources (dashed line) is also reported. In our data sample, stars represent sources that are doubtfully correlated with outflows, while arrows signify lower limit measurements either in mass flux (vertical) or in bolometric luminosity (horizontal) in the absence of a complete sampling. Sources designated red are from the sample of [1] (filled squares), for which we have calculated new values of Lbol with Spitzer photometry (filled triangles), in order to check if the correlation derived in [1] could be affected by a less accurate bolometric luminosity of their sample. A shift to lower luminosities is observed also for these sources but does not seem to be sufficient to invalidate the correlation. Most of the Serpens sources stand well above the correlation of [1], which confirms their classification as Class 0 YSO

References 1. Bontemps, S., Andre, P., Terebey, S., & Cabrit, S. 1996, A&A, 311, 858 2. Choi, M., Evans, N. J., II, & Jaffe, D. T. 1993, Revista Mexicana de Astronomia y Astrofisica, 27, 91

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3. Davis, C. J., Matthews, H. E., Ray, T. P., Dent, W. R. F., & Richer, J. S. 1999, MNRAS, 309, 141 4. Eiroa, C., Torrelles, J. M., Curiel, S., & Djupvik, A. A. 2005, AJ, 130, 643 5. Enoch, M. L., Glenn, J., Evans, N. J., II, Sargent, A. I., Young, K. E., & Huard, T. L. 2007, ApJ, 666, 982 6. Williams, J. P., & Myers, P. C. 1999, ApJ, 518, L37 7. Winston, E., et al. 2007, ApJ, 669, 493

Numerical Simulations of Herbig Haro Objects: A Low Excitation HH Object Alejandro Esquivel, Alejandro C. Raga, and Fabio De Colle

Abstract The spectra of some low excitation HH objects could be explained by a variable-velocity jet impinging on a dense medium. The leading shock (i.e. “the head”) of the jet would move slowly into the medium, thus producing a low excitation spectra even though the velocity at the base of the jet would correspond to a high excitation one. We present a high-resolution numerical simulation of such a jet using a newly developed parallel two-dimensional adaptive-mesh hydrodynamical code.

A. Esquivel () and A.C. Raga Instituto de Ciencias Nucleares, Universidad Nacional Aut´onoma de M´exico, A. Postal 70-543, M´exico D.F. 04510, M´exico e-mail: [email protected]; [email protected] F. De Colle Dublin Institute for Advanced Studies (DIAS), 31 Fitzwilliam Place, Dublin 2, Ireland e-mail: fdc@cp,dias.ie

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1 Introduction Raga & Cant´o (2003) [2], hereafter RC03, proposed an analytical model to explain the case of low-excitation emission in the leading head of an HH object considering a variable jet impinging on a dense medium. The idea is that the “head”of the jet which impacts a high density medium will be slowed and cooled-down considerably. But the knots, that are the result of a variable velocity, travel more freely inside the cavity swiped by the head, catch up with it, and deposit their momenta keeping the motion of the head. The resulting spectra will be of low excitation after the head has cooled down (quickly due to the interaction with a high density medium), and will show flashes of high excitation only when the knots impact the leading shock from behind. At a given time, it is more probable to observe the jet at its low excitation phase. RC03 also provided analytical predictions of the position of the head as a function of time.

Fig. 1 Left panel: temperature stratification for t D 1300 yr. The abscissa corresponds to the outflow axis, and the ordinate to the cylindrical radius (which has been reflected for visual purposes). Right Panel: position of the leading shock as a function of time. The analytical prediction in RC03 is the solid line, the “C” signs are the position measured from the numerical simulation.

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2 The Model We used the newly developed parallel hydrodynamical code with adaptive mesh refinement (AMR) WALICXE [1] to model a HH object impinging on a dense medium. We used a block-based binary mesh with 7 levels of refinement, the maximum resolution is 2:4  1013 cm, in a computational domain of .4  1/  1017 cm (radial and axial, respectively), which is sufficient to (barely) resolve the cooling distances behind the leading shock. Both the jet and the (static) medium are initially at 1000 K, with hydrogen densities of 10 and 100 cm3 , respectively. The jet velocity varies as v.t / D v0 Œ1 C 0:5 sin .2= /, with v0 D 200 km s1 , and D 50 yr. For a more detailed description of the model we refer the reader to [1]. We show in Figure 1 some of the results of the simulation.

3 Conclusions/Summary We presented a numerical model of a HH object interacting with a high density medium using a newly developed parallel AMR hydrodynamical code. We compare the results with analytical predictions and obtain a reasonable agreement. The head of the jet,however, does not slow-sown as much as predicted in the analytical model. This can be justified because the analytical model considered a fixed jet cross-section. But the jet’s head, becomes narrower, reducing the cross section, and thus the drag with the medium.

References 1. Esquivel, A, & Raga, A. C., de Colle, F. 2008, ApJ, in prep. 2. Raga, A. C., & Cant´o, J. 2003, A & A, 412, 745

Soft X-rays from DG Tau: A Physical Jet Model ¨ Hans Moritz Gunther, Sean P. Matt, and Zhi-Yun Li

DG Tau is a classical T Tauri star (CTTS) showing an unusual X-ray spectrum, best described by two thermal components with different absorption columns. The soft X-rays are less absorbed than the hard, presumably coronal, component [4]. This rules out stellar accretion as the origin of the soft photons, which is a standard model for CTTS and successfully explains the emission in e.g. TW Hya [5]. Instead, the observations of DG Tau require an emission region above the circum-stellar absorption layer. A good candidate is the jet of DG Tau, which is resolved in X-rays out to a few arcseonds using Chandra [3]. Additionally there is a 30 AU offset between the hard, coronal and the softer X-ray emission of the central source [7].

H.M. G¨unther () Hamburger Sternwarte, Gojenbersgweg 112, 21029 Hamburg, Germany e-mail: [email protected] S.P. Matt Department of Astronomy, University of Virginia, P.O. Box 400325, Charlottesville, VA 22904, USA e-mail: [email protected] Z.-Y. Li Department of Astronomy, University of Virginia, P.O. Box 400325, Charlottesville, VA 22904, USA e-mail: [email protected]

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The jet has also been observed in the optical with HST/STIS. It consists of components with different velocities, where the faster components reach up to 600 km/s and they are more collimated than the slower components. This can be traced in H˛ and forbidden sulfur and oxygen lines [1]. The outermost wind is a molecular outflow [2]. We suggest that the soft, unresolved X-rays originate from shocks in a narrow, fast inner wind component bracketted by slower outflows as observed in the optical (Fig. 1, left). The geometry is cylindrical with the shock at the cylinder base. The outflowing matter is heated to X-ray emitting temperatures in the shock front and cools radiatively within the post-shock cooling length dcool , where faster velocities vshock lead to higher temperatures and larger dcool [6]. Several scenarios can lead to the formation of the shock: Stationary collimation shocks, wind shocks or internal working surfaces caused by unsteady launching velocities are possibilities. Using all available X-ray data from Chandra and XMM-Newton we fit a two temperature model. We explore the parameter space of the soft component, keeping the values for the hard emission fixed. Unfortunately the temperature is not very well constraint, because there is –as always– an ambiguity between soft emission and extra absorption (Fig. 1).

Fig. 1 left: Sketch of our model. The innermost and fastest component of the outflow passes through a shock front. The cooling zone has a cylindrical geometry. right: Estimated dimensions of the shock for a density 105 /cm3 . The contours encircle the 68%, 90% and 99% confidence regions.

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We divide the observed volume emission measure by the density taken from optical observations and by dcool to obtain the shock area with radius R (Fig. 1, right). As a result we find that our model successfully describes the observed spectra. In all cases the dimensions of the shock are only a few AU, below the resolution limit of the optical observations. Thus it is possible that the X-ray shock cannot be seen in the optical data. From vshock and the radius of the cylinder base the mass flux can be estimated: 1010 Msun /yr are sufficient to explain the emission; this is at least three orders of magnitude below the total mass loss in the jet. The high extinction towards the central source allows the spatially distributed emission to be detected in DG Tau, but it is possible that emission from a jet base also contributes in other CTTS to the observed spectra. A grating spectrum of the soft emission could help to narrow down the errors on our results significantly.

References 1. Bacciotti, F., Mundt, R., Ray, T. P., et al. 2000, ApJ, 537, L49 2. Beck, T. L., McGregor, P. J., Takami, M., et al. 2008, ApJ, 676, 473 3. G¨udel, M., Skinner, S. L., Audard, M., et al. 2008, A&A, 478, 797 4. G¨udel, M., Telleschi, A., Audard, M., et al. 2007, A&A, 468, 515 5. G¨unther, H. M., Schmitt, J. H. M. M., Robrade, J., et al. 2007,A&A, 466, 1111 6. Raga, A. C., Noriega-Crespo, A., & Velazquez, P. F. 2002, ApJ, 576, L149 7. Schneider, C. & Schmitt, J. H. M. M., 2008, A&A, 488, L13

Multifluid Simulations of the Kelvin-Helmholtz Instability in a Weakly Ionised Plasma Aoife C. Jones, Mohsen Shadmehri, and Turlough P. Downes

1 Initial Set-up These simulations demonstrate the Kelvin-Helmholtz instability in the presence of multifluid effects, by examining the behaviour of a weakly ionised three-fluid plasma, consisting of electrons, ions and neutrals. The simulations are carried out for three cases, that of ideal magnetohydrodynamics (MHD), multifluid MHD dominated by ambipolar resistivity, and multifluid MHD dominated by Hall resistivity. The code used is that described in O’Sullivan and Downes [1] and the initial parameters are the same as those used by Palotti et al. [2]. The initial set-up involves two fluids flowing parallel with velocities CV0 and V0 in the y-direction. The magnetic field is initially uniform, and aligned with the interface, Fig. 1a. The system is subjected to a small perturbation in the transverse velocity, as given by 2 vx D ıv sin.ky y/ exp  x 2 , where ıv V0 . The system is then allowed to

A.C. Jones () School of Mathematical Sciences, DCU, Ireland e-mail: [email protected]

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evolve with time. The Kelvin-Helmholtz instability is seen to form along the interface of the two plasmas. This instability is observed in the development of the Kelvin’s cat’s eye vortex, Fig. 1b.

2 Results In order to examine the growth of the Kelvin-Helmholtz instability, we plot the behaviour of two parameters with time for each of our three cases: the transverse R plasma kinetic energy, 12 v2x dxdy, Fig. 2a, and the total magnetic energy of the R system, 12 B 2 dxdy, Fig. 2b. For comparison, the respective plots from Palotti et al. [2] have been included, Fig. 2c and d. These results relate to a non-ideal MHD plasma using homogeneous resistivity and a series of values of magnetic Reynolds numbers.

3 Analysis and Conclusions Our results are preliminary, however they show good agreement with those by Palotti et al. [2]. We can also see a distinct difference in the evolution of KH instability between the ideal and non-ideal MHD plasmas. The growth times for the multifluid cases are clearly longer than in the ideal case, as seen from the slower increase in the kinetic energy. Even more interesting is the difference in the evolution

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of the magnetic field. In the ideal case, the total magnetic strength increases to a maximum, and then slowly dissipates. In the presence of ambipolar diffusion, this maximum reached is considerably lower. This is presumably due to magnetic reconnection. The process of magnetic reconnection is made possible only by the inclusion of multifluid effects, or by the use of numerical viscosity. In the case of the Hall effect, the magnetic field is seen to build steadily for the duration of the simulation. This would suggest the presence of a magnetic dynamo, and should prove interesting to investigate further. Future work will also include a comparison with analytical studies, such as Shadmehri and Downes [3].

References 1. O’Sullivan, S., Downes, T.P.: An explicit scheme for multifluid magnetohydrodynamics. Mon. Not. R. Astron. Soc. 366, 1329–1336 (2006) 2. Palotti, M.L., Heitsch, F., Zweibel, E.G., Huang, Y.-M.: Evolution of Unmagnetized and Magnetized Shear Layers. Astrophys. J. 678, 234–244 (2008) 3. Shadmehri, M., Downes, T.P.: The role of Kelvin-Helmholtz instability in dusty and partially ionized outflows. Mon. Not. R. Astron. Soc. 387, 1318–1322 (2008)

Large-scale 3D Simulations of Protostellar Jets: Long-term Stability and Jet Rotation Kai Cai, Jan Staff, Brian P. Niebergal, Ralph E. Pudritz, and Rachid Ouyed

K. Cai () Department of Physics & Astronomy, McMaster University, 1280 Main St. W., Hamilton, ON, L8S 4M1, Canada and Department of Physics and Astronomy, University of Kentucky, Lexington, KY 40506, USA e-mail: [email protected] J. Staff Department of Physics & Astronomy, Louisiana State University, 202 Nicholson Hall, Tower Dr., Baton Rouge, LA 70803, USA e-mail: [email protected] B.P. Niebergal and R. Ouyed Department of Physics & Astronomy, University of Calgary, 2500 University Drive NW, Calgary, Alberta, Canada, T2N 1N4 R.E. Pudritz Department of Physics & Astronomy, McMaster University, 1280 Main St. W., Hamilton, ON, L8S 4M1, Canada

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1 Introduction Jets from newborn stars are believed to be launched magneto-centrifugally from the sources [2, 7]. During the past few years, high resolution spectra of protostellar jets obtained by the Hubble Space Telescope (HST), especially those near the jet base (see, e.g., [9]), have made it possible for a direct comparison with jet simulation results. The observations of radial velocity shifts in several T Tauri jets ([1, 3]) that are suggestive of jet rotation, also reinforce the notion that the jets carry the bulk of the angular momentum away from accretion disks, favoring the disk wind mechanism ([7]). Here we present our latest three-dimensional time-dependent MHD simulations of jets launched from the surface of Keplerian disks. For simulation setups, see Staff et al. in this volume). The setup is similar to that used by [5] (hereafter OCP), but we start with field lines more open (negative ’s).

2 Results Staff et al. presents the case with initial threading magnetic field D 0:01 (‘potential’ configuration used in [6]), here we report on the simulation with D 0:25, a case resembles BP82 magnetic field configuration ([8]). Figure 1a(left) shows the density structure in a slice cut along the jet axis, after 1600 rotations of the inner disk. As shown by Fig.1 in Staff et al. and Fig. 1b (right), both jets exhibit strong magnetic “backbone” along the center, which act to reduce the average jet Alfv´enic Mach number, making the stabilizing process described in OCP plausible.

Fig. 1 Left: Snapshot of log. / of the D 0:25 simulation at t D 1900. The numbers on the colorbar is in units of logΠ=g=cm3 , while the axes of the plot is in units of ri . Right: The poloidal component of the magnetic field Bp at the same time. The numbers shown on the colorbar is in code units. A color version of this plot can be found in the electronic version of this book and in Appendix A (Fig. A.36).

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3 Summary and Conclusions We have extended OCP simulations, with different threading magnetic field configurations, to a much larger spatial scale, making a direct comparison with HST observations possible. The jets are able to maintain a long-term stability through a self-limiting process, as described in OCP. In addition, the jet tends to expand sideways, probably due to the pressure by the generated toroidal field. This may also help to stabilize the jet ([4]). The simulated jets exhibit complex density structure (possibly due to various instabilities) and a number of internal time-dependent shocks, yielding good prospects for producing forbidden line emission (see Staff et al.) as observed. In the future we hope to use the jet rotation profiles to constrain the initial magnetic field configuration threading the underlying accretion disk. Acknowledgements This work was made possible by the facilities of the Shared Hierarchical Academic Research Computing Network (SHARCNET:www.sharcnet.ca). K. C. acknowledges support from a CITA National Fellowship.

References 1. Bacciotti, F., Ray, T. P., Mundt, R., Eisl¨offel, J., & Solf, J. 2002, ApJ 576, 222 2. Blandford, R. D. & Payne, D. G. 1982, MNRAS 199, 883 3. Coffey, D., Bacciotti, F., Ray, T. P., Eisl¨offel, J., & Woitas, J. 2007, ApJ 663, 350 4. Moll, R., Spruit, H. C., & Obergaulinger, M., 2008, A&A, 492, 621 5. Ouyed, R., Clarke, D. A., & Pudritz, R. E. 2003, ApJ, 582, 292 6. Ouyed, R., & Pudritz, R. E., 1997, ApJ, 482, 712 7. Pudritz, R. E., Ouyed, R., Fendt, C. & Brandenburg, A. 2007, Protostars and Planets V, p.277 8. Pudritz, R. E., Rogers, C. S., & Ouyed, R. 2006, MNRAS 365, 1131 9. Ray, T. P., Dougados, C., Bacciotti, F., Eisl¨offel, J., & Chrysostomou, A. 2007, Protostars and Planets V, 231

Extragalactic Jets with Helical Magnetic Fields Rony Keppens and Zakaria Meliani

Abstract Extragalactic jets harbor dynamically important, organized magnetic fields. We explore with grid-adaptive, high resolution numerical simulations the morphology of AGN jets pervaded by helical field and flow topologies. We concentrate on the long term evolution of kinetic energy dominated jets, penetrating denser clouds. The jets have near-equipartition magnetic fields, and radially varying Lorentz factor profiles maximally reaching   22. The helicity of the beam magnetic field is effectively transported down the beam, with compression zones in between diagonal internal cross-shocks showing stronger toroidal field. The high speed jets have localized, strong toroidal field within the backflow vortices and a more poloidal field layer, compressed between jet beam and backflows. This layer stabilizes the jet beam. We infer emission intensity, suggesting a clear trend were highly structured beams are found for toroidal fields, while inner beam cross-shocks and thin hotspots are detectable for poloidal topologies. Significant jet deceleration only occurs beyond distances exceeding O.100Rj /, as the axial flow can

R. Keppens () and Z. Meliani Centre for Plasma-Astrophysics, K.U.Leuven, Belgium e-mail: [email protected]; [email protected]

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reaccelerate downstream to the internal cross shocks. This reacceleration is magnetically aided by field compression across the internal shocks that pinch the flow.

1 Equations, Setup, and Jet Propagation Characteristics The governing relativistic MHD equations express particle number conservation, energy-momentum conservation, and Maxwell’s equations for perfectly conducting plasmas. These can be written in conservation form @t U C r F.U/ D 0, where U D .D; S; E; B/T . Among the conserved variables used, B is the magnetic field in the fixed Lorentzian lab frame. To solve this set numerically, we use AMRVAC [1,2]. This code features an any-D implementation of AMR, and up to 3 AMR approaches: patch-based, block-based (octree), and hybrid block-based [1]. As initial condition, we set up a magnetic configuration in the entire domain, and inject a high Lorentz factor jet. The magnetic configuration is helical internal to the jet, while the jet B' B and rotation v' D p' . The pressure profile p.R/ velocity has vZ D ˛ p .R=a/ ensures radial force balance along Z D 0. We studied in [3] up to 8 models with varying field topology. We find in all cases that vortical patterns form at the contact, and complex backflows surround the jet beam. Recurrent cross-shocks are driven into the beam, leading to interacting diagonal cross-shock patterns. The magnetic field gets compressed by the reverse shock, enhancing collimation. The jet cocoon with backflow has a strong azimuthal field, and the helically magnetized jet beam gets surrounded by toroidal field. The propagation speed of the bow shock systematically exceeds the value expected from estimates using beam-averages, in accord with the centrallypeaked  .R/ injected at inlet. Only the jet beam, backflow, and a compressed region between beam and backflow contain significant magnetic pressure. The jet beam develops a near vertical field ‘sheet’, while the vortical patterns contain the most tightly wound fields. Vortices appear as more isolated protrusions into the jet cavity bounded by the bow shock. The rotational flows in these vortices are supersonic. We simulated the jet propagation until a (small) deceleration of the relativistic jet is evident, as seen in the decrease in Lorentz factor along the axis.

2 Power Maps for Varying Field We used the simulations to infer ‘synchrotron’ intensity maps from  2 v2 B 2 sin2 . /. The angle-dependent factor, together with the magnetic field variation with radius, explains the trend seen from more outer beam sensitivity (toroidal field cases), to inner beam sensitivity at cross-shock fronts (poloidal field cases). Two cases are shown in Fig. 1. In summary, we performed relativistic grid-adaptive MHD simulations of AGN jets. An analysis of the shok-dominated propagation dynamics showed that the magnetic helicity changes at the internal cross-shocks repeatedly

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Fig. 1 For a toroidal magnetic field topology (left), to a poloidal field model (right): arbitrary scaled synchrotron ‘power’. The plots show a zoom of the entire simulation domain

reaccelerate the jet. The synchrotron intensity maps indicate that one may distinguish the magnetic topology from their visual appearance.

References 1. B. van der Holst, R. Keppens, J. Comp. Phys. 226, 925–946 (2007) 2. B. van der Holst, R. Keppens, Z. Meliani, Comp. Phys. Commun. 179, 617–627 (2008) 3. R. Keppens, Z. Meliani, B. van der Holst, F. Casse, A&A 486, 663–678 (2008)

Jets from Collapsing Stars Volodymyr Kryvdyk

Abstract The formations of relativistic jets by stellar collapse are considered. These jets will form in the polar caps of magnetosphere, when a stellar magnetic field increases during collapse and the charged particles will accelerate. These jets will generate the non-thermal radiation. The analysis of particle dynamics and its emission in the stellar magnetosphere under collapse show that the collapsing stars can by powerful sources of the jets and the non-thermal radiation. The radiation flux grows with decreasing stellar radius and can be observed in the form of radiation burst with duration equal to the stellar collapse time. This radiation can be observed by means of modern astronomical instruments.

1 Introduction The several models are proposed for relativistic jets from compact objects. [1] and [2] assumed that relativistic electron accelerated and jets produced at large distances from holes in magnetized accretion disc by means electromagnetic fields surrounding rotating black holes. In [4] magnetic cannonball model the relativistic jets are generated in magnetosphere of compact objects by their collapse to black hole. [8] considered the formation of relativistic jets during the evolution of rotating helium stars, in which iron-core collapse does not produce a successful out going shock but instead forms a black hole. In this model simulated deposition of energy in the polar regions results in strong relativistic outflow jets. In paper [9] has studied of the possible production of supernovae and a variety of high-energy transients by black hole formation in massive stars endowed with rotation: the “collapsar model”. The energy of the jet and the explosion it produces depend upon the efficiency of MHD processes in extracting accretion energy from the disk. [10] performed

V. Kryvdyk () Dept. Astronomy, Faculty of Physics, Taras Shevchenko Kyiv National University, av. Glushkova 2/1, Kyiv 03022, Ukraine e-mail: [email protected]

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2.5-dimensional general relativistic magnetohydrodynamic (MHD) simulations of the gravitational collapse of a magnetized rotating massive star. The simulation results show the formation of a disklike structure and the generation of a jetlike outflow inside the shock wave launched at the core bounce. Magnetohydrodynamic simulations of a rotating massive star collapsing to a black hole are investigated in paper [3]. In this paper is perform two-dimensional, axisymmetric, magnetohydrodynamic simulations of the collapse of a rotating star in light of the collapsar model of gamma-ray bursts, and the formation of an accretion disk around a black hole and the jet production near the black hole is investigated. We consider the formation of relativistic jets and the radiation by collapse of stars having heterogeneous magnetospheres. The stellar magnetosphere compress during the collapse and its magnetic field increases considerably. A cyclic electric field is produced and the charged particles will accelerate, and the relativistic jets will formed in polar caps of stellar magnetosphere.

2 Jets in Magnetosphere of Collapsing Star The particle spectrum in the magnetosphere and its evolution during collapse was considered in paper [7]. The transformation of the stellar magnetosphere during collapse is show on Fig. 1. We can see that the initial stellar magnetosphere transform during collapse and the polar jets formed in magnetosphere.

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References 1. Blandford, R.D., Znajek, R.L.: Electromagnetic extraction of energy from Kerr black holes. Mon. Not. Roy. Astron. Soc. 179, 433–440 (1977). 2. Blandford, R. D., Payne, D. G.: Hydromagnetic flows from accretion discs and the production of radio jets. Mon. Not. Roy. Astron. Soc. 199, 883–894 (1982). 3. Fujimoto, S., Kotake, K., Yamada, S., Hashimoto, M., and Sato, K.: Magnetohydrodynamic Simulations of a Rotating Massive Star Collapsing to a Black Hole. Astrophys. J. 644, 1040–1056 (2006). 4. Hanami, H.: Magnetic Cannonball Model for Gamma-Ray Bursts. Astrophys. J. 491, 687–696 (1997) 5. Kryvdyk, V.: Electromagnetic radiation from collapsing stars. I. Power-series distribution of particles in magnetospheres. Mon. Not. Roy. Astron. Soc. 309, 593–598 (1999). 6. Kryvdyk, V. : High- energy emission from presupernova. Adv. Space Res. 33, 484–486 (2004). 7. Kryvdyk, V.: Formation of the relativistic jets by collapse star to black hole. Adv. Space Res. 42, 533–537 (2008). 8. MacFadyen, A.I., Woosley, S.E.: Collapsars: Gamma-Ray Bursts and Explosions in “Failed Supernovae”. Astrophys. J. 524, 262–289 (1999). 9. MacFadyen, A. I., Woosley, S.E., Heger, A.: Supernovae, Jets, and Collapsars. Astrophys. J. 550, 410–425 (2001). 10. Mizuno Y., Yamada S., Koide S, and Shibata K.: General Relativistic Magnetohydrodynamic Simulations of Collapsars. Astrophys. J. 606, 395–412 (2004).

Outflows in High-Mass Star Forming Regions Ana L´opez-Sepulcre, Claudio Codella, Riccardo Cesaroni, Maite T. Beltr´an, Nuria Marcellino, and Luca Moscadelli

Abstract A sample of 11 high-mass star forming regions (SFRs) in the earliest evolutionary phases has been mapped in the 13 CO(2-1) and C18 O(2-1) lines using the IRAM 30-m radio telescope. The goal is to detect molecular outflows and determine their parameters. Molecular outflows have been detected in six of our sources. Their mean mass and mass loss rate are 100 Mˇ and 5  103 Mˇ /yr respectively, suggesting that the powering Young Stellar Objects (YSOs) are massive and probably still in the accretion phase.

A. L´opez-Sepulcre (), R. Cesaroni, and L. Moscadelli INAF - Osservatorio Astrofisico di Arcetri, Largo E. Fermi, 5, 50125 Firenze (Italy) e-mail: [email protected] C. Codella INAF - Istituto di Radioastronomia, Sezione di Firenze, Firenze (Italy) M.T. Beltr´an Dep. d’Astronomia i Meteorologia, Facultat de F´ısica, UB Barcelona (Spain) N. Marcelino DAMIR, CSIC, Madrid (Spain)

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1 The Structure and Kinematics of Massive SFRs The formation of stars more massive than 8 Mˇ is still matter of debate. Two main scenarios have been proposed to explain the formation mechanism of high-mass stars: accretion through a scaled-up version of the low-mass star formation picture, with more massive disks and/or higher accretion rates, and coalescence of lowermass stars in dense clusters. A recent study of G24.78C0.08 and G31.41C0.31 [1] has revealed circumstellar toroids rotating perpendicular to the outflow axis, clearly supporting the accretion scenario. This encouraging result calls for a statistical study to check whether this is the general mechanism to form massive stars. In this paper we present single-dish 13 CO(2-1) and C18 O(2-1) observations at 1.3 mm carried out towards a sample of 11 high-mass SFRs with IRAM 30-m antenna at Pico Veleta (Granada, Spain) in September 2006. All the targeted sources have high luminosities (L > 1000 Lˇ ), present compact sub-mm continuum and H2 O, OH and CH3 OH maser emission, and most of them do not show any detectable free-free emission. These properties ensure that our sample is composed of embedded massive young stellar objects. For each source, 40  40 On-The-Fly maps were obtained using the multi-beam HERA receiver. The HPBW for both lines is 1100 , whereas the spectral resolution is 0.1 km/s. The comparison between the 13 CO(2-1) and C18 O(2-1) lines (Fig. 1) allowed us to observe line wings in six out of the 11 sources and to reveal self-absorbed profiles suggesting infall. The median mass and mass loss rate of the detected molecular outflows are, respectively, 100 Mˇ and 5  103 Mˇ /yr, indicating that the powering sources are massive YSOs. In several cases (e.g. Fig. 1 Right), the C18 O shows velocity gradients perpendicular to the outflow axis, suggesting the existence of circumcluster rotating envelopes. The mere existence of an outflow indicates that star formation is very active and supports follow-up interferometric observations in high density molecular tracers. G35.20–0.74 20 60 CO(2–1)

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Outflows in High-Mass Star Forming Regions

Reference 1. Beltr´an, M.T., Cesaroni, R., Neri, R. et al. 2005, A&A, 435, 901

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Astrophysical Jet Experiment Berenice Loupias, Claire Michaut, Chris D. Gregory, Emeric Falize, Jonathan Waugh, Dono Seiichi, S. Pikuz, Yasuhiro Kuramitsu, Alessandra Ravasio, Serge Bouquet, Wigen Nazarov, Youichi Sakawa, Nigel Woolsey, and Michel Koenig

Abstract We present an experimental characterization of jet propagation in an ambient medium. An intense laser (LULI2000) was used to generate the plasma jet using foam filled cone target. We observed, with several diagnostics, a perturbation in the interaction region between the jet and the ambient medium. The effect of the ambient medium on the jet velocity is also presented.

1 Context The experiments, when properly diagnosed, allow astrophysical theoretical models and computer simulation codes to be tested. This way requires the existence of scaling laws which ensure the complete similarity between the astrophysical object and the experiment. Scaling laws has been already demonstrated for radiative hydrodynamic phenomena (optically thin and thick medium) [1]. The collimated supersonic jets observed around young stars (YSO) can be studied in experiments [2, 3, 4, 6, 5, 7, 4]. The jets associated with YSO are often seen to have a chain of

C.D. Gregory, M. Koenig, B. Loupias, and A. Ravasio LULI, Ecole Polytechnique, France S. Bouquet, E. Falize, and C. Michaut () LUTH, Observatoire de Paris, CNRS, Universite Paris-Diderot, France S. Bouquet and E. Falize D´epartement de Physique Th´eorique et Appliqu´ee, CEA-DIF, France J. Waugh and N. Woolsey Department of Physics, University of York, UK Y. Kuramitsu, Y. Sakawa, and D. Seiichi Graduate School of Engineering, Osaka University, Japan S. Pikuz Multicharged Ions Spectra Data Center of VNIIFRTI, Mendeleevo, Russia W. Nazarov University of St. Andrews, School of Chemistry, UK

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high visible emission knots and to terminate with a bow shock. Astrophysical jets are very complex for an experiment to probe the whole of their dynamics. In this experiment, we selected and examined a part of the dynamics: jet interaction with the interstellar medium (ISM).The bow shock structure observed in astronomical data appear with a very perturbated and fragmented shape [8].

2 Experimental Setup and Results For this experiment we used a long pulse laser (LULI2000, 1 kJ in 1.5 ns) to generate the plasma jet and a short pulse beam (PICO2000, 100 J in 1 ps) to produce protons for radiography. Targets were brominate dopped foam filled cone (entrance hole diameter 500 m, exit hole diameter 100 m, density 50 or 100 mg.cm3 ) which have been used in previous experiment where the jet parameters were measured for its propagation in vacuum. A solid target as a pusher is placed over the entrance hole to drive a strong shock through the cone and results in the expulsion of a plasma jet (more details are reported in [2]). The plasma jet generation at the rear side of the target is a very important point, to limit any jet or ambient medium interaction with the laser beam. The jet velocity has been measured between 80–145 km.s1 from the transverse self optical pyrometer (SOP) which decreased when ambiant gaz density increased. In the Fig. 1, the probe delay is 30 ns from the arrival of the drive laser beam. The effects of ambiant medium are a high density feature propagating in front of the jet, as the flow acts as a piston driving a shock through the ambient gas. The shock front is not uniform, and small scale perturbations are seen around the leading edge.

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Due to the higher density in this shock front, the optical probe is not able to provide any quantitative measurement of the electronic density. With proton radiography, we observed the same perturbations and same dimensions than in the visible interferogram and the whole structure can be probed. We have performed plasma jet propagation experiment through an ambient medium by varying jet and/or environment densities. For all shots, we were able to measure electronic density, propagation and radial velocity, jet shape due to a large panel of diagnostics. The propagation of the supersonic jet through the ambient medium at the target rear side suggests that the observed highly perturbated region is the bow shock launching by the plasma jet.

References 1. 2. 3. 4. 5. 6. 7. 8.

Falize E, Bouquet S, Michaut C: J. Phys.: Conf. Ser. 112, 042016 (2008) Loupias B, Gregory CD, Falize E et al: to be published in Astrophys. Space Sci. (2008) Loupias B, Falize E, Gregory CD et al: J. Phys.: Conf. Ser. 112, 042022 (2008) Loupias B, Koenig M, Falize E et al: Phys. Rev. Lett. 99, 265001 (2007) Gregory CD, Howe J, Loupias B et al: Astrophys. J. 676, 420 (2008) Blue BE, Weber SV, Glendinning SG et al: Phys. Rev. Lett. 94, 095005 (2005) Lebedev SV, Ciardi A, Ampleford DJ et al: Plasma Phys. Control. Fusion 47, 465 (2005) Hartigan P: Astrophys. J. 339, 987 (1989)

The Angular Momentum of Dense Clumps in Elephant Trunks Veronica Lora, Alejandro C. Raga, and Alejandro Esquivel

Abstract The radiation from newly born massive stars photoionize and erode the parental molecular cloud, producing structures such as the so-called elephant trunks. At the head of an elephant trunk, the interaction of theshock (driven by the photo evaporation process) with previously existing density perturbations leads to the formation of dense clumps. Some of these clumps have enough mass to be autogravitating, and therefore can eventually form new stars. We carry out a 3D simulation of this process, and from the results we compute the angular momenta of these collapsing clumps. We show that the angular momenta of the clumps have preferential directions, which in principle indicate the directions in which jets will eventually be ejected from the star+accretion disk systems that will be formed.

1 Model and Setup We used a version of the hydrodynamic adaptive grid code YGUAZU-A that includes a two temperature equation of state (for neutral and ionized material) and plane-parallel radiative transfer of ionizing photons at the Lyman limit. We used a computational box of .3  1:5  1:5/  1018 cm and we had a resolution of 5:8  1015 cm. We let the model run from t D 0 to 600 kyr then we focused on the spatially connected structures with density above 3  1018 gcm3 and compute the angular momentum respect to the center of mass of each clump. We then computed the angle between the x-y components of the angular momentum and found that most

V. Lora () Instituto de Astronom´ıa, Universidad Nacional Aut´onoma de M´exico, Ap. 70-264, 04510 D.F., M´exico e-mail: [email protected] A.C. Raga and A. Esquivel Instituto de Ciencias Nucleares, Universidad Nacional Aut´onoma de M´exico, Ap. 70-543, 04510 D.F., M´exico e-mail: [email protected]; [email protected]

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Fig. 1 Clumps at all times and the angle made with the x-direction of the angular momentum

of the clumps make a 80–90 angle with the direction of the incomming radiation (x-direction). See Fig. 1. We then consider only the most massive clumps (M > 0.5 Mˇ ) and follow the mass evolution. We found that these dense clumps remain above the Jeans limit enough time to eventualy collapse and form stars.

2 Conclusions Photoevaporation fragments neutral clouds in many dense clumps with different masses. We computed the angular momenta of these clumps and gave more attention to the most massive/lasting ones.In general the x-y angle associated to each clump tend to be high (80–90), implying that the angular momentum is perpendicular to the x-axis (direction of the radiation) thus the jets of the protostar are expected to have the same orientation, which is observed.

References 1. Bally, J. & Reipurth, B., 2003, AJ, 126(2), 893. 2. Esquivel, A. & Raga, A.C., 2007, MNRAS, 377, 383. 3. Kajdiˇc, P. & Raga, A.C., 2007, AJ, 670, 1173.

A Precessing Jet in the NGC 2264 G Outflow Carolyn McCoey, Paula S. Teixeira, Michel Fich, and Charles J. Lada

Abstract We (Teixeira et al., 2008) present IRAC imaging of the NGC 2264 G protostellar outflow region. A jet in the red (eastern) outflow lobe is clearly detected in all four IRAC bands, and is shown to continuously extend over the entire length of the red outflow lobe as seen in CO. The easternmost part of the jet exhibits multiple changes of direction, which we find can be largely explained by a slowly precessing jet. The changes in the jet direction may be sufficient to account for a significant fraction of the broadening of the outflow lobe, as observed in CO emission. NGC 2264 was observed with IRAC onboard the Spitzer Space Telescope as part of the Spitzer Guaranteed Time Observation program 37 [2]. Figure 1 shows a colour composite image of the red lobe (east) and part of the blue lobe (west) of the NGC 2264 G outflow region, where we detect emission from a jet in all four IRAC bands (brightest in band 2, 4.5 m). The jet exhibits several changes of direction, at .˛1 ; ı1 /.J2000/ D .06h 41m 17:0s , C09ı 560 2000 ), .˛2 ; ı2 /.J2000/ D .06h 41m 25:5s , C09ı 550 4500 ) and .˛3 ; ı3 /.J2000/ D .06h 41m 28:5s ; C09ı 550 5200 /. We used a simple nonrelativistic jet model from [5] in order to model the observed jet morphology. Figure 2 shows (overplotted on a divided 4.5 m/3.6 m image of the jet) the best fitting precession model, which is an anti-clockwise rotating jet with a precession angle of 8ı , an angle of 185ı in the plane of the sky, and an inclination angle (of the blue lobe) to the line of sight of 82ı . NGC 2264 G was mapped in CO(2-1) by [6, 4] and their low- and high-velocity contours are also plotted in Fig. 2. Assuming the high velocity CO traces a molecular component of the jet, we may infer that the radial velocity of the jet is 40 km s1 . For a typical jet velocity of 200 km s1 this corresponds to an inclination angle to the line of sight i > 78ı (i.e., within 22ı from the plane of the sky). Lada and Fich [6] concluded from the highly bipolar and collimated nature of the outflow that the

C. McCoey () and M. Fich University of Waterloo and University of Western Ontario P.S. Teixeira and C.J. Lada, Harvard-Smithsonian Center for Astrophysics now ESO-Garching, Cambridge, MA 02138, USA e-mail: [email protected]

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Fig. 1 IRAC image of the NGC 2264 G outflow region. The jet extends eastwards over 1.1 pc, assuming a distance of 800 pc [9] from VLA 2. The vertical lines mark the locations where the jet changes direction, and the horizontal arrows he location of the red and blue CO lobes [7]. A color version of this plot can be found in the electronic version of this book and in Appendix A (Fig. A.37)

Fig. 2 A divided 4.5 m/3.6 m jet image with CO contours overlaid. The Psolid and dotted contours correspond to high (<39 km s1 ) [<27 km s1 ] and low (<8 km s1 ) [<3 km s1 ] velocity red- [blue-]shifted CO, respectively [4]. The solid line is the precession model fit

opening angle of the flow measured from the jet axis, be 8ı and that the outflow should lie within 20ı of the sky. Both of these constraints are consistent with our precession model. Furthermore, taking again a typical protostellar jet velocity of 200 km s1 we estimate a precession period of 7300–8500 yr, which compares well with the upper limit of 104 yr found by [6] for the dynamical time-scale of the outflow. Finally, we note that the motion of a precessing jet may significantly broaden an outflow. Our precession model has an opening angle of 16ı , compared with 25ı for the low velocity CO gas in NGC 2264 G, [6]: therefore jet precession accounts for roughly 64% of the observed extent of the outflow lobes. Acknowledgements PST acknowledges support from the scholarship SFRH/BD/13984/2003 awarded by the Fundac¸a˜ o para a Ciˆencia e Tecnologia (Portugal). MF&CM are supported by a Natural Sciences and Engineering Research Council of Canada (NSERC) Discovery Grant. This work is based [in part] on observations made with the Spitzer Space Telescope, which is operated by the Jet Propulsion Laboratory, California Institute of Technology under a contract with NASA.

References 1. Davis, C. J., & Eisl¨offel, J., 1995, A&A, 300, 851 2. Fazio, G. G., et al., 2004, ApJS, 154, 10 3. Fich, M., & Lada, C. J., 1997, ApJL, 484, L63

A Precessing Jet in the NGC 2264 G Outflow 4. Fich, M., & Lada, C. J., 1998, ApJS, 117, 147 5. Gower, A. C., Gregory, P. C., Unruh, W. G., & Hutchings, J. B.,1982, ApJ, 262, 478 6. Lada, C. J. & Fich, M., 1996, ApJ, 459, 638 7. Margulis, M., & Lada, C. J., 1986, ApJL, 309, L87 8. Teixeira, P. S., Mc Coey, C., Fich, M., & Lada, C. J., 2008, MNRAS, 384, 71 9. Walker, M. F., 1956, ApJS, 2, 365 10. Ward-Thompson, D., Eiroa, C., & Casali, M. M., 1995, MNRAS, 273, L25

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Line Diagnostics of Large Scale Jets from Classical T Tauri Stars: The Case of DG Tau Fiona McGroarty, Linda Podio, Francesca Bacciotti, and Tom Ray

Abstract Large-scale optical outflows have been observed from a number of Classical T Tauri stars. Recently the “BE” method has been applied to these outflows allowing us to determine the density, ionisation fraction and temperature of gas in the outflow. Some results of the study of the HH 158 jet from DG Tau using this method are briefly described here.

1 Introduction The “BE” method [1, 2] allows electron density (ne ), ionisation fraction (xe ), electron temperature (Te ) and hydrogen density (nH D ne =xe ) of gas in an outflow to be obtained directly from observational data. This is based on the fact that collisionally excited shocks emit optical forbidden emission lines (FELs) and uses the ratios of FELs in [SII], [OI] and [NII] to determine these physical parameters. This technique has been used successfully on a number of outflows to date [3]. Here we apply it to large-scale outflows from Classical T Tauri stars (CTTSs) of 0.5 pc [4]. Preliminary results of our study of the HH 158 jet from the CTTS DG Tau are presented here.

F. McGroarty () Department of Physics and Astronomy, Queen’s University Belfast, Belfast BT7 1NN, Northern Ireland e-mail: [email protected] L. Podio Dublin Institute for Advanced Studies, 31 Fitzwilliam Place, Dublin 2, Ireland e-mail: [email protected] F. Bacciotti INAF/Osservatorio Astrofisico de Arcetri, Florence, Italy e-mail: [email protected] T. Ray Dublin Institute for Advanced Studies, 31 Fitzwilliam Place, Dublin 2, Ireland e-mail: [email protected] K. Tsinganos et al. (eds.), Protostellar Jets in Context, Astrophysics and Space Science Proceedings, DOI 10.1007/978-3-642-00576-3 81, c Springer-Verlag Berlin Heidelberg 2009 

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2 Results and Analysis – DG Tau Jet (HH 158) These observations were taken in February and September 2007 using the Intermediate Dispersion Spectrograph on the Isaac Newton Telescope on La Palma. One pixel projects to 000.4 on the sky. Figure 1 shows how ne , xe , Te and nH vary along the jet. For the two inner jet sections (at 200.6 and 100 ) only ne can be determined. The general trends seen in this data with increasing distance from the source are: ne decreases, Te increases across each jet section but there is a decrease from section 1 to section 2 and nH shows an exponential-like decrease. xe values are generally quite high (0.5) with the exception of the inner-most measurable value which is lower by a factor of 2. These trends and the values of the physical parameters are in good agreement with previous studies using this method. For further discussion of these results and the data for the outflows from CW Tau and RW Aur see [5].

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Line Diagnostics of Large Scale Jets from Classical T Tauri Stars: The Case of DG Tau

References 1. Bacciotti, F., Chiuderi, C., & Oliva, E. 1995, A&A, 296, 185 2. Bacciotti, F., & Eisl¨offel, J. 1999, A&A, 342, 717 3. Podio, L., Bacciotti, F., Nisini, B., Eisl¨offel, J., et al. 2006, A&A, 456, 18 4. McGroarty, F., & Ray, T. P. 2004b, A&A, 420, 975 5. McGroarty, F., Podio, L., Bacciotti, F., & Ray, T. P. in preparation

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Relativistic Two-Component Hydrodynamic Jets Zakaria Meliani and Rony Keppens

Abstract Astrophysical jets from various sources seem to be stratified, with a fast inner jet and a slower outer jet. As it is likely that the launching mechanism for each component is different, their interface will develop differential rotation, while the outer jet radius represents a second interface where disruptions may occur. We explore the stability of stratified, rotating, relativistic two-component jets, in turn embedded in static interstellar medium. In a grid-adaptive relativistic hydrodynamic simulation with the AMRVAC (Adaptive Mesh Refinement version of the Versatile Advection code), the non-linear azimuthal stability of two-component relativistic jets is investigated. We simulate until multiple inner jet rotations have been completed. We find evidence for the development of an extended shear flow layer between the two jet components, resulting from the growth of a body mode in the inner jet, Kelvin-Helmholtz surface modes at their original interface, and their nonlinear interaction. Both wave modes are excited by acoustic waves which are reflected between the symmetry axis and the interface of the two jet components. Their interaction induces the growth of near stationary, counterrotating vortices at

Z. Meliani () and R. Keppens Centre for Plasma-Astrophysics, K.U. Leuven, Belgium e-mail: [email protected]; [email protected]

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the outer edge of the shear flow layer. The presence of a heavy external jet allows their further development to be slowed down, and maintains of a collimated flow. At the outer jet boundary, small-scale Rayleigh-Taylor instabilities develop, without disrupting the jet configuration.We demonstrate that the cross-section of two-component relativistic jets, with a heavy, cold outer jet, is non-linearly stable.

1 Equations, and Setup We set up a two-component structured jet with a typical total kinetic luminosity flux set to Ljet;Kin D 1046 ergs=s. The external radius of the two-component jet is taken to be Rout  0:1 pc. We set the inner jet radius Rin D Rout =3, and we ensure that this inner component carries only a small fraction fin D 0:1% of the total kinetic luminosity flux, while the external jet carries the remaining fout D 99:9%. For the initial condition, we adopt a uniform outflow poloidal Lorentz factor z;out  3, for the outer, slow jet within Rin < R < Rout . The inner jet has a fast outflow set to poloidal Lorentz factor z;in  16:6. Two-component jet models also indicate that the spin of the inner beam is higher than the spin of the external jet. Therefore, the initial rotation adopts two different profiles, one for the inner, and one for the external jet. This is inevitable when these two jet components are launched from different regions and with different mechanisms. We choose for the external jet a radially decreasing rotation profile, with a decrease faster than Keplerian, as this component is launched from an accretion disk and the external streamlines in the jet expand faster than the inner streamlines. Moreover, we assume that the angular momentum extracted along each streamline varies little, since we adopt a toroidal velocity varying as 1=R. For p the initial inner jet rotation profile, we take a toroidal velocity increasing with R from the axis. This toroidal velocity vanishes on axis and the equivalent classical centrifugal force given by this profile is constant. We consider a discontinuity in the toroidal velocity at the boundary between the two jets when fixing V';in D 0:05 and V';out D 0:005. Finally, the pressure is deduced by supposing initial transverse equilibrium between the pressure gradient and the centrifugal force, by making the assumption of a constant (but different for inner versus outer jet) local ‘effective polytropic index’. The governing equation of state is taken as a Synge-type relation, also used in [2].

2 Nonlinear Evolution and Stability We find that the interaction between the two rotating jet components induces the development of surface and body normal modes. These azimuthal instabilities decelerate the inner jet and decollimate the external jet. A cross-section is shown in Fig. 1, and the observed deformation should influence the 3D structure of the jet. More details are found in [1].

Relativistic Two-Component Hydrodynamic Jets

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References 1. Z. Meliani, R. Keppens, A&A 475, 785–789 (2007) 2. Z. Meliani, C. Sauty, K. Tsinganos, N. Vlahakis, A&A 425, 773–781 (2004)

The Physical Properties of the RW Aur Bipolar Jet from HST/STIS High-Resolution Spectra Stanislav Melnikov, Jochen Eisl¨offel, Francesca Bacciotti, Jens Woitas, and Tom Ray

1 Introduction A well-collimated bipolar jet emanating from the classic T Tauri star RW Aur A (HH 229) has been traced out to 1500 in the redshifted [2] and over 10000 in the blueshifted lobe [3]. Long-slit spectra of the RW Aur jet were taken on 10. Dec. 2000 with the STIS spectrograph on board the HST. The spectrograph slit of width 000 :1 and length 1000 was moved across the jet in seven positions from NW to SE, keeping it parallel to the jet axis. As a result, seven spectra covering the jet in both dimensions were obtained. The observations taken with the G750M grating cover the three prominent forbidden emission doublets – [O I]  6300,6363, [N II]  6548,6583, and [S II]  6716,6731. The current S. Melnikov (), J. Eisl¨offel, and J. Woitas Th¨uringer Landessternwarte Tautenburg (TLS), Sternwarte 5, D-07778 Tautenburg, Germany F. Bacciotti I.N.A.F. – Osservatorio Astrofisico di Arcetri, Largo E. Fermi 5, I-50125 Firenze, Italy T. Ray Dublin Institute for Advanced Studies, 31 Fitzwilliam Place, Dublin 2, Ireland e-mail: [email protected] K. Tsinganos et al. (eds.), Protostellar Jets in Context, Astrophysics and Space Science Proceedings, DOI 10.1007/978-3-642-00576-3 83, c Springer-Verlag Berlin Heidelberg 2009 

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study presents a continuation of the analysis of the HST/STIS spectroscopy of the RW Aur jet started in Woitas et al. [5]. For extraction of the physical quantities we followed the well-established method described in Bacciotti and Eisl¨offel [1].

2 Diagnostic Results from Spectroscopy: Basic Physical Parameters With this method, we were able to extract the parameters up to 300 :9 in redshifted and up to 200 :1 in blueshifted jet. The electron density ne at the basis of the jet is about nH D 4:6  105 cm3 and decreases gradually with distance from the source. The electron temperature Te in the redshifted jet decreases fast from 2  104 K to 104 K and then varies around the latter value. Contrary to this, in the blueshifted section the Te starts from 104 K and varies to higher values. The ionisation fraction xe varies between 0.04 and 0.4. In both lobes xe increases within the first few arcseconds and then starts to decrease. The variation of ne ; xe ; Te along the jets appears to be correlated with position of knots in both lobes. The hydrogen density xH in the brighter redshifted lobe seems to be higher (log.nH / D 4:8) than in the fainter blueshifted section, with a mean value of log.nH / D 4:4. With these values, the RW Aur jet seems to be one of the densest T Tauri jets studied so far.

3 Radial Velocity Asymmetry and Mass Flux Rate Our study confirms a strong radial velocity (RV) asymmetry in the opposite lobes: the averaged RV is 190 km s1 in the blueshifted lobe, and only C100 km s1 in the opposite direction. The velocity in the red jet has no prominent dependence with distance from the source. There is only a bump of RV at the position of the bright knots at 100 :8 and some decrease after 300 . In contrast, the RV in the blue jet does not exhibit any correlation with knots and shows a gradual decrease along the lobe. The derived physical quantities, together with the kinematical parameters of the jet, allow us to estimate the mass flux rate along the jet flows. Used method is based on measured forbidden line ratios [4]. The calculated mean mass flux MP j is similar in the investigated sections of both lobes: 2:5  109 Mˇ yr1 for the red- and 2:3  109 Mˇ yr1 for the blueshifted jet. The MP j =MP acc ratio along the jet lobes lies within the range predicted by MHD models (0.01–0.1) or less than its lower limit.

4 Comparison with Other T Tau Jets The RW Aur jet seems to be one of the densest objects studied so far, but other properties are quite similar to other T Tauri jets. This suggests that the jet density (i.e. the amount of ejected material) may be linked to the intensity of the accretion

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processes, taking into account that RW Aur is a very strong accretor. The behaviour of the RW Aur jet in regions close to the source is similar to that of HH 30 and DG Tau jets investigated with similar resolution. Acknowledgements We thank Linda Podio for her code for the jet mass flux calculation. J.E. and S.M. acknowledge support from the Deutsches Zentrum f¨ur Luft- und Raumfahrt. This study was supported in part by the European Community’s Marie Curie Research and Training Network JETSET.

References 1. Bacciotti, F., Eisl¨offel, J.: Ionization and density along the beams of Herbig-Haro jets. A&A 342, 717–735 (1999) 2. Dougados, C., Cabrit, S., Lavalley, C., M´enard, F.: T Tauri stars microjets resolved by adaptive optics. A&A 357, L61–L64 (2000) 3. Mundt, R., Eisl¨offel, J.: T Tauri Stars Associated with Herbig-Haro Objects and Jets. AJ 116, 860–867 (1998) 4. Nisini, B., Bacciotti, F., Giannini, T., Massi, F., Eisl¨offel, J., Podio, L., Ray, T.P.: A combined optical/infrared spectral diagnostic analysis of the HH1 jet. A&A 441, 159–170 (2005) 5. Woitas, J., Ray, T.P., Bacciotti, F., Davis, C.J., Eisl¨offel, J.: ApJ 580, 336–342 (2002)

Stability of Magnetized Spine-Sheath Relativistic Jets Yosuke Mizuno, Philip E. Hardee, and Ken-Ichi Nishikawa

Abstract In order to investigate the stability of magnetized spine-sheath relativistic jets, we have performed 3D RMHD simulations of weakly and strongly magnetized relativistic jets embedded in a weakly and strongly magnetized stationary or mildly relativistic (0.5c) sheath flow using the RAISHIN code. In the numerical simulations a jet with Lorentz factor,  D 2:5 is precessed to break the initial equilibrium configuration. The prediction from a normal mode analysis of the linearized RMHD equations of increased stability of a weakly-magnetized system with mildly relativistic sheath flow to Kelvin-Helmholtz instabilities and the stabilization of a strongly-magnetized system with mildly relativistic sheath flow is confirmed by the numerical simulations.

1 RMHD Spine-Sheath Simulations Relativistic jets have been observed in galaxies and quasars (AGNs) [10,1], in black hole binary star systems (microquasars) [8], and are thought responsible for the gamma-ray bursts (GRBs) [7]. Recent GRMHD simulations of jet formation [6, 3, 5] indicate that highly collimated high speed jets driven by the magnetic fields threading the ergosphere may themselves reside with a broader wind or sheath outflow driven by the magnetic fields anchored in the accretion disk.

Y. Mizuno () Center for Space Plasma and Aeronomic Research, The University of Alabama in Huntsville, 320 Sparkman Drive, NSSTC 2104, Huntsville, AL 35805, USA e-mail: [email protected] P.E. Hardee Department of Physics and Astronomy, Univesity of Alabama, Tuscaloosa, AL 35487, USA K.-I. Nishikawa CSPAR, The University of Alabama in Huntsville, Huntsville, AL 35805, USA

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Circumstantial evidence such as the requirement for large Lorentz factors suggested by the TeV BL Lacs when contrasted with much slower observed motions has been used to suggest the presence of a spine-sheath morphology [2]. We have investigated the stability properties of highly relativistic jet flows allowing for the effects of strong magnetic fields and relativistic flow in a sheath around a highly relativistic jet by 3D RMHD numerical simulations. The numerical simulations are performed with the 3D GRMHD code “RAISHIN” using Cartesian coordinates in special relativity [9]. We consider the following initial conditions for the simulations: a “preexisting” jet is established across the computational domain. The jet speed is vj D 0:9165c and  D 2:5. The initial magnetic field is assumed to be uniform and parallel to the jet flow. A precessional perturbation is applied at the inflow by imposing a small transverse velocity. In order to investigate the effect of an external wind, we have performed simulations with no external wind (ve D 0:0c) and a mildly relativistic external wind (ve D 0:5c). We have performed two sets of simulations. In the weakly magnetized (RHD) simulations, the Alfv´en speed is much smaller than sound speed. In the strongly magnetized (RMHD) simulations, the Alfv´en speed is about twice the sound speed. The weakly magnetized simulations show growing velocity oscillation from the helical Kelvin-Helmholtz (KH) instability. These simulations are fully evolved to the non-linear phase. An external wind reduces the growth of KH instability and delays the onset of the non-linear phase. The strongly magnetized simulation without an external wind also shows growing velocity oscillation, albeit more slowly. The strongly magnetized simulation with a wind shows damped velocity oscillation. Comparison of the different cases shows that stronger magnetic field reduces the growth rate of the helical KH instability, that a wind reduces the growth rate and in combination with strong magnetization leads to damping of the KH instability and to jet stabilization. These numerical results agree with theoretical predictions [4]. Acknowledgements Y.M., P.H., and K.I.N acknowledge partial support by NSF AST-0506719, AST-0506666, NASA NNG05GK73G, NNX07AJ88G, and NNX08AG83G. The simulations have been performed on Columbia Supercomputer at NAS Division in NASA Ames Research Center.

Stability of Magnetized Spine-Sheath Relativistic Jets

References 1. A. Ferrari, ARAA 36 539 (1998) 2. G. Ghisellini, F. Tavecchio, & M. Chiaberge, A&A 432 401 (2005) 3. J.F. Hawley, & J.H. Krolik, ApJ 641 103 (2006) 4. P.E. Hardee, ApJ 664 26 (2007) 5. P. Hardee, Y. Mizuno, & K.-I. Nishikawa, Ap & SS 311 281 (2007) 6. J.C. McKinney, MNRAS 368 1561 (2006) 7. P. M´esz´aros, Rep Prog Phys 69 2259 (2006) 8. I.F. Mirabel, & L.F. Rodr´ıguez, ARAA 37 409 (1999) 9. Y. Mizuno, P. Hardee, & K.-I. Nishikawa, ApJ 662 835 (2007) 10. C.M. Urry, & P. Padovani, PASP 107 803 (1995)

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Chemical Models of Hot Molecules at Shocks in Outflows Hideko Nomura and Tom J. Millar

It is observationally known that outflows are associated with many young stars, and they are thought to originate from accretion disks around the stars. Meanwhile, molecular line observations have shown that some molecules such as CH3 OH and SiO are very abundant at shocks and/or clumps in outflows [1]. We have constructed a chemical model at shock fronts in outflows associated with young stars by calculating gas-phase chemical reactions initiated by icy mantle evaporation from dust grains [3]. Most of the rate coefficients of the gas-phase reactions are taken from the UMIST database RATE06 [4]. A parameterized two-fluid, steady 1D MHD shock model [2] is adopted. Our results show that different physical conditions lead to a variety of molecular abundance ratios in outflows. For example, the molecular abundances at a shock front depend on the shock velocity. At high velocity molecular hydrogen is partly dissociated, and some molecules which are easily destroyed by atomic hydrogen, such as CH3 OH and H2 CO, becomes less abundant. Accordingly, the abundances

H. Nomura () Department of Astronomy, Kyoto University, Japan e-mail: [email protected] T.J. Millar ARC, School of Mathematics and Physics, Queen’s University Belfast, UK

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of their daughter species, such as CH3 OCH3 , also drop (Fig. 1, left). If the dust temperature is low, gas-phase molecules accrete on dust grains in a dense post-shock region, depending on their binding energies onto the grains. Meanwhile, the dust temperature and gas velocisty in a disk at the upper stream also affect the molecular abundances in the outflow. Some nitrogen-bearing daughter species, such as HCN and CH3 CN, becomes abundant if they move into the outflow after they spend enough time in a hot inner region of the disk (Fig. 1, right). Our results suggest that observations of molecular abundance ratios will trace the physical structure and the chemical processes, especially gas-grain interactions, at shock fronts in outflows.

References 1. Bachiller, R., P´erez Guti´errez, M., Kumar, M.S.M., & Tafalla, M.: A&A, 372, 899 (2001) 2. Jim´enez-Serra, I., Caselli, P., Mart´ın-Pintado, J. & Hartquist, T.W.: A&A, 482, 549 (2008) 3. Nomura, H. & Millar, T.J.: A&A, 414, 409 (2004) 4. Woodall, J., Agundez, M., Markwick-Kemper, A.J., & Millar, T.J.: A&A, 466, 1197 (2007)

Survival of H2 and CO in MHD Disk Winds of Class 0, Class I and Class II Stars Despina Panoglou, Paolo J.V. Garcia, Sylvie Cabrit, and Guillaume Pineau des Forˆets

Abstract We present a model of warm self–similar disk winds [1], aimed at constraining the origin of molecular jets from young stars. We computed the thermal properties, ionization structure and chemical evolution, after imposing an extended molecular network of species and reactions [2]. Here we report results for typical class 0, class I and class II stars. In particular, no H2 and CO destruction occurs for younger stars (higher accretion rates), since the temperature, the ionization fraction and the X-ray photoreaction rate are lower. In general, there is as less molecular gas remaining at the recollimation point as older is the protostar.

1 Test Cases We chose to examine the differences in the molecular abundances of disk winds in three different classes of stars: (a) A typical class 0 star of mass equal to Mˇ =10 and with accretion rate MP acc D 5  106 Mˇ =yr, describing typical submillimeter sources. (b) A typical class I star of mass Mˇ =2 and accretion rate of MP acc D 106 Mˇ =yr, that resembles IR protostars. (c) A typical class II star of mass Mˇ =2 as well but with the accretion rate being lower, MP acc D 107 Mˇ =yr.

D. Panoglou () Faculdade de Ciˆencias, Universidade do Porto, Portugal Universit´e Pierre et Marie Curie – Paris 6, France e-mail: [email protected] P.J.V. Garcia Faculdade de Engenharia, Universidade do Porto, Portugal e-mail: [email protected] S. Cabrit LERMA, Observatoire de Paris-Meudon, France e-mail: [email protected] G.P. des Forˆets Institut d’Astrophysique Spatiale, Orsay, France e-mail: [email protected]

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2 Results With the inclusion of an appropriate radiation field and the calculation of the UV and X-ray components, the ionization fraction is lower by more than one order of magnitude in the class 0 star, and so does the drift speed, resulting to lower temperatures (500 K in the class 0 case, against 1 000 K and 2 000 K in the class I and class II stars, respectively). The main observed molecules, H2 and CO (Fig. 1) do not dissociate at all in the younger class 0 star: The attenuated radiation flux is much lower than the other two cases, preventing photoreactions from being effective; and temperature is also lower resulting into low rates for reactions in the gas phase. No matter the age, the mass or the accretion rate of a protostar, the main mechanisms that govern the thermochemical evolution in disk winds are the same: Reactions between various species in the gas phase and photodissociation appear to govern the chemical profile of wind flow line; drag heating and molecular cooling determine the temperature; ionization by the X-rays and the UV field account for the ionization fraction. Nonetheless, we see H2 varying from no dissociation at all in the class 0 case, to 40% dissociation in the class I star, and 50% in the class II star. Acknowledgements This work was supported by the European Community’s Marie Curie Actions Human Resource and Mobility within the JETSET (Jet Simulations, Experiments and Theory) network under contract MRTN-CT-2004 005592.

References 1. Casse, F., and J. Ferreira, 2000, A&A 361, 1178. 2. Flower, D. R., G. Pineau des Forˆets, and T. W. Hartquist, 1985, MNRAS 216, 775.

Three-Fluid Magnetohydrodynamics in Star Formation Cecilia Pinto and Daniele Galli

Abstract Interstellar magnetic fields influence all stages of the process of star formation, from the collapse of molecular cloud cores to the formation and evolution of circumstellar disks and protostellar jets. This requires us to have a full understanding of the physical properties of magnetized plasmas of different degrees of ionization for a wide range of densities and temperatures. We derive general equations governing the magneto-hydrodynamic evolution of a three-fluid medium of arbitrary ionization, also including the possibility of charged dust grains as the main charge carriers ([4]). We complement this analysis computing accurate expressions of the collisional coupling coefficients for a variety of gas mixtures relevant for the process of star formation ([3]).

C. Pinto () Dipartimento di Astronomia e Scienza dello Spazio,Universit`a di Firenze, Largo E. Fermi 5, 50125 Firenze, Italy e-mail: [email protected] D. Galli INAF-Osservatorio Astrofisico di Arcetri, Largo E. Fermi 5, 50125 Firenze, Italy e-mail: [email protected]

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1 Magneto-Hydrodynamic Equations and Momentum Transfer Rate Coefficients We derive an advection-diffusion equation for the evolution of the magnetic field and general expressions for the resistivities, the diffusion time scales and the heating rates in a three-fluid medium and we use them to estimate the evolution of the magnetic field in molecular cloud cores and protostellar jets ([4]). As example in Fig. 2 we show our results evaluating ambipolar diffusion heating in protostellar jets, which can provide a good fraction of the required energy input in the outer layers of weakly-ionized jets. These evaluations require us to compute the collisional coupling coefficients. Accurate numerical values for momentum transfer rates are obtained as function of the gas temperature and for various values of the drift velocities. An example is displayed in Fig. 1, where we note that the often used polarization approximation (Langevin rate) fails by order of magnitude ([3]).

Fig. 1 The momentum transfer rate coefficient for e–H2 collisions as a function of the temperature T for vd D 0 (thick solid curve); and vd D 50 km s1 ; vd D 100 km s1 (thin solid curves, top to bottom), compared with the values by [1] for the same values of the drift velocity (dotted curve).The dashed line shows the Langevin rate ([2])

Fig. 2 Ambipolar diffusion heating rate as function of the distance from the jet’s axis x D R=R0 (R0  103 AU) compared with the mechanical heating rate empirically determined by [5] for three values of the distance from the central star (top to bottom: 0.01, 0.1 and 1 pc, dot-dashed lines). The dotted and solid curves are for an ion number fraction of ni =.ni C nn / D 0:5 and 0.01

Three-Fluid Magnetohydrodynamics in Star Formation

References 1. B.T. Draine, W.G. Roberge, A. Dalgarno, ApJ. 264, 485 (1983) 2. P. Langevin, Annales de Chimie et de Physique, Series 8, 5, 245 (1905) 3. C. Pinto, D. Galli, A&A, 484, 17-28 (2008) 4. C. Pinto, D. Galli, F. Bacciotti, A&A, 484, 1-15 (2008) 5. H. Shang, S. Lizano, A. Glassgold, F. Shu, ApJ. 612, 69 (2004)

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Physical Conditions of the Shocked Regions of Planetary Nebulae Angels Riera, Alejandro C. Raga, Garrelt Mellema, Alejandro Esquivel, and Pablo F. Vel´azquez

Abstract We have computed a series of axisymmetric simulations of a cloudlet travelling away from a photoionizing source to reproduce the emission arising from the compact knots observed in some Planetary Nebulae. The predicted spectra agree approximately with the observed spectra when shown in two-line ratio diagnostic diagrams. The predicted and observed spatial distributions of the emission (with high ionization lines extending more towards the source than lower ionization lines) agree in a qualitative way.

1 Introduction The unexpected discoveries of collimated outflows and shock-excited features in Planetary nebulae (PNe) have led over the years to different terminologies. Small scale structures found in a fraction of PNe, such as strings of knots appearing as symmetrical pairs, point-symmetrical features or jet structures, which move supersonically with respect to the main body of the nebula, are often referred to as FLIERs (Fast Low Ionization Emission Regions). The word “jet” is generally restricted to thin radial features. FLIERs are characterized by low ionization spectra, arcsecond scale sizes (or 1016 cm at a typical distance of 1 kpc) and supersonic velocities (with Doppler shifts of ˙25-200 km s1 ). FLIERs have been identified with either

A. Riera () Departament de F´ısica i Enginyeria Nuclear, Universitat Polit`ecnica de Catalunya (Spain) e-mail: [email protected] A.C. Raga, A. Esquivel, and P.F. Vel´azquez Instituto de Ciencias Nucleares, Universidad Nacional Aut´onoma de M´exico (M´exico) e-mail: raga,esquivel,[email protected] G. Mellema Stockholm University (Sweden) e-mail: [email protected]

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a working surface of a jet or a shocked cloudlet. The small sizes and high outflowing velocities of the FLIERs strongly support that these microstructures are associated to outward-flowing bullets or clumps.

2 Numerical Simulations and Results We carried out a series of axisymmetric numerical simulations of a high velocity cloudlet moving away from an ionizing photon source through a uniform medium (for a detailed description of these numerical simulations see [2]). In the initial configuration, we have a spherical cloudlet of density D 103 cm3 , temperature D 104 K, and a radius D 1016 cm moving at a velocity of 40, 70, 100 or 150 km s1 along the axis. For the photoionizing source, we have assumed a blackbody source with a luminosity of 5000 Lı , and an effective temperature of 50000 K or 70000 K. The source is located at distances of (3, 1, 0.3) 1018 cm of the centre of the cloudlet. Simulations with low photoionization rates show that the flow evolves in a qualitatively similar way to what has been found in previous, axisymmetric simulations of shocked cloudlets, with a fragmentation of the fully shocked cloudlet through a series of vortex shedding events ([1]). The effect of the increase of the photoionization on the flow structure is quite dramatic, producing a broader bow shock structure, and less fragmentation. From all models, the [O III] emission shows the longest extension away from the bow-shock head, towards the photon source. This behaviour qualitatively agrees with the [O III] surface brightness distribution observed in FLIERs. We obtained a set of diagnostic diagrams involving several emission line ratios, commonly used to discriminate photoionized nebulae from shock-excited objects. From these diagrams, we have seen that the observed and predicted line ratios cover similar regions of the diagnostic diagrams.

References 1. Raga, A. C., Esquivel, A., Riera, A. & Vel´azquez, P. F. 2007, ApJ, 668, 310 2. Raga, A. C., Riera, A., Mellema, G., Esquivel, A. & Vel´azquez, P. F. 2008, A&A, 489, 1141

The Jets of the Proto-Planetary Nebula CRL 618 Angels Riera, Alejandro C. Raga, Pablo F. Vel´azquez, Sinhue Haro-Corzo, and Primoz Kajdic

Abstract We present here the kinematic structure and the excitation conditions of the collimated outflows of the proto-planetary nebula CRL 618 based on high spatial resolution spectroscopy obtained with STIS onboard HST. The spectra obtained show a linear increase of the radial velocity with distance to the central source. We find that the emission line ratios observed in the clumpy lobes of CRL 618 are similar to high or low-excitation HH excitation class depending on the emission line ratio.

1 Observations We have analyzed STIS data which were retrieved from the HST Data Archive. The long slit spectra were obtained with the slit oriented along the bipolar axis of CRL 618 (Cycle 11 proposal 9430), using the 5200  0:200 slit and the G430L, G750L and G750M gratings. This configuration provides continuous wavelength coverage ˚ The nominal spatial sampling was of 0.00 05 pixel1 and the from 3000 to 10250 A. 00 spatial resolution 0.2 . Spectra were obtained at two positions with a spatial offset of 0.00 894 perpendicular to the slit. The long-slit STIS spectra were calibrated at the Space Telescope Science Institute following the standard HST pipeline calibration.

A. Riera () Departament de F´ısica i Enginyeria Nuclear, Universitat Polit`ecnica de Catalunya (Spain) e-mail: [email protected] A.C. Raga, P.F. Vel´azquez, and S. Haro-corzo Instituto de Ciencias Nucleares, Universidad Nacional Aut´onoma de M´exico (M´exico) e-mail: [email protected]; [email protected]; [email protected] P. Kajdic Instituto de Geof´ısica, UNAM (M´exico) e-mail: [email protected]

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2 Kinematical Properties The observed position-velocity diagrams are reminiscent of those observed in HH objects with some knots showing sudden drops in the radial velocity (e.g., [2]). These sharp drops in velocity at some knots of the lobes are consistent with the expected kinematical behaviour for jets formed by a succession of internal working surfaces. The high spectral resolution spectra show two ramps of increasing velocity (i.e. more negative velocities). A linear increase of the radial velocity with distance is the predicted behaviour in a wake of a bullet, and may also result from a timedependent ejection velocity.

3 Emission Line Ratios and Excitation Conditions Emission line intensities have been obtained from the one-dimensional spectra obtained by adding over the spatial extent of each feature observed in the lobes. Next, we have obtained a set of dereddened .E.B  V/ D 1/ line ratios. The values observed at the bow shock-like features within the lobes of CRL 618, were compared with a compilation of HH objects (taken from [1]). Most of the observed emission line ratios seem to agree quite well with the location of the HH objects. The [O III]/Hˇ ratios of all bow-shock like feature show values just above the value defining the “high excitation” group of HHs (i.e. [O III]/Hˇ  0:1). In the [N II]/[O I] (vs. H˛/[O I] plot, CRL 618 partially overlaps the region of the “lowexcitation” HH objects. Summarizing, CRL 618 seems to belong to one or other HH excitation class depending on the emission line ratio.

References 1. Raga, A.C, B¨ohm, K.-H. & Cant´o, J. 1996, RMxAA, 32, 161 2. Raga, A.C., Noriega-Crespo, A., Reipurth, B. et al. 2002, ApJ, 565, L29

The Formation of Filamentary Structures in Radiative Cluster Winds Ary Rodr´ıguez-Gonz´ales, Alejandro Esquivel, Alejandro C. Raga, and Jorge Cant´o

Abstract We explore the dynamics of a “cluster wind” flow in the regime in which the shocks resulting from the interaction of winds from nearby stars are radiative. We show that for a cluster with low-intermedia mass stars, the wind interactions are indeed likely to be radiative. We then compute three dimensional, radiative simulations of a cluster of 75 young stars, exploring the effects of varying the wind parameters and the density of the initial ISM that permeates the volume of the cluster. These simulations show that the ISM is compressed by the action of the winds into a structure of dense knots and filaments.

1 Radiative Losses in a Cluster Wind In order to estimate the cool filaments formation, we consider a stellar cluster with a local stellar density n (Dnumber of stars per unit volume), of stars with identical, isotropic winds with a mass loss rate MP w and a terminal wind velocity vw . The typical separation between stars then is D D n1=3 . The shock interactions between nearby stars will be radiative if  dcool =D < 1 (where, dcool is the cooling distance). is function of MP w  .0:1pc=D/ and vw , where MP w is the mass loss rate and vw the wind velocity (see Fig. 1, left panel and Rodriguez-Gonzalez et al. 2008). We have carried out numerical simulations (using the “yguaz´u-a” code) in which we place 75 stars which are uniformly distributed inside a sphere of radius Rc D 0:5 pc. The stars have identical mass loss rates MP w D 106 Mˇ and wind velocities vw D 100 km/s. The enviroment is homogeneous medium with nenv D 104 cm3 and Tenv D 1000 K. Figure 1 shows a central density condensation. The main

A. Rodr´ıguez-Gonz´ales (), A. Esquivel, and A.C. Raga Instituto de Ciencias Nucleares, Universidad Nacional Aut´onoma de M´exico, Ap. 70-543, 04510 D.F., M´exico e-mail: [email protected]; [email protected]; [email protected] J. Cant´o Instituto de Astronom´ıa, Universidad Nacional Aut´onoma de M´exico, Ap. 70-264, 04510 D.F., M´exico

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Fig. 1 Left panel shows the parameter: The thick curve shows the D 1 contour, the other contours represent values of at successive factors of two. The region above the thick curve therefore represents the parameter space in which the two-wind interactions in a cluster wind flow are radiative. Right panel shows a volume rendition of the density structure, after a t D 1:9  105 yr time-integration. The grey-scale shows the density values in g cm3 , linearly scaled, The cubic domain has a size of 1.2 pc

contribution to the column density of the central condensation comes from the stellar wind material. However, we demostrate (Rodriguez-Gonzalez et al. 2008) that the ISM column density are proportional to the observed extinction.

2 Conclusions We find that if the low-intermedia mass stars have wind velocities vw < 200 km s1 , the stellar wind interactions between nearby stars are highly radiative for mass loss rates as low as MP w  3  107 Mˇ yr1 . However, for higher wind velocities, considerably higher mass loss rates are required for the wind interactions to be radiative (when D  0:1 pc between the stars in the cluster). These values indicate that in a cluster of Herbig Ae/Be stars many of the wind interactions between nearby stars are likely to be in the highly radiative regime. We show, numerically, that it does lead to the production of a structure of dense filaments and clumps. From our simulations, we obtain predicted column density which show spatial structures which are qualitatively similar to the observations of IRAS 18511C0146 (Vig et al. 2007).

References 1. Vig, S. et al. 2007, A&A, 470, 977. 2. Rodr´ıguez-Gonz´alez, A., Esquivel, A., Raga, A. C. & Cant´o, J., 2008, ApJ, 684, 1384.

Hydrodynamic Modeling of Accretion Shock on CTTSs Germano G. Sacco, Constanza Argiroffi, Salvatore Orlando, Antonio Maggio, Giovanni Peres, and Fabio Reale

High resolution (R  600) X-ray observations of some classical T Tauri stars (CTTSs) (TW Hya, BP Tau, V4046 Sgr, MP Mus and RU Lupi) have shown the presence of X-ray plasma at T  2–3  106 K and denser than ne  1011 cm3 [1, 2, 3, 4, 5], which suggests an origin different from the coronal one (ne  1010 cm3 ). Stationary models demonstrated that X-ray emission from CTTSs could also be produced by the accreting material [6]. We address this issue with the aid of a time-dependent hydrodynamic numerical model describing the impact of an accretion stream onto the chromosphere of a CTTS (see [7] for more details). Our simulations include the effects of gravity, radiative losses from optically thin plasma, the thermal conduction and a detailed modeling of the stellar chromosphere. The

G.G. Sacco () Consorzio COMETA, Via S. Sofia, 64, 95123, Catania, Italy e-mail: [email protected] C. Argiroffi, G. Peres, and F. Reale DSFA, Universit`a di Palermo, Piazza del Parlamento, 1, Palermo Italy S. Orlando and A. Maggio INAF-Osservatorio Astronomico di Palermo, Piazza del Parlamento, 1, Palermo, Italy

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a

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gravity and the parameters describing the accretion stream (density ne D 1011 cm3 and velocity v D 450 km s1 ) were chosen in order to match the X-ray properties of the CTTSs MP Mus [4]. The evolution of temperature and density, during the first part of our simulation is shown in Fig. 1. During the first 300 s, the shock heats up the accretion column creating a slab with density ne D 3  7  1011 cm3 and temperature T  3  106 K Subsequently, due to thermal instabilities [8] triggered by radiative cooling at the base of the hot slab, plasma temperature and pressure decrease of more than two orders of magnitude in few seconds, leading to the collapse of the accretion column and the cooling of the whole system in 100 s. After the accretion column is completely cooled, a new hot slab is generated and the system starts a quasi-periodic evolution with alternating heating and cooling phases lasting 400 s. This work is described in detail in [7]. Acknowledgements We thank J.J. Drake for useful discussions. This work was supported in part by the Italian Ministry of University and Research (MIUR), by Istituto Nazionale di Astrofisica (INAF) and by the European Community’s Marie Curie Research and Training Network JETSET (Jet Simulation, Experiments and Theory) under contract MRTN-CT-2004 005592. S.O., G.P. and F.R. acknowledge support from Marie Curie Fellowship No. MTKD-CT-2005-029768. The software used in this work was in part developed by the DOE-supported ASC / Alliance Center for Astrophysical Thermonuclear Flashes at the University of Chicago. This work makes use of results produced by the PI2S2 Project managed by the Consorzio COMETA, a project co-funded by the Italian Ministry of University and Research (MIUR) within the Piano Operativo Nazionale “Ricerca Scientifica, Sviluppo Tecnologico, Alta Formazione” (PON 2000–2006).

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References 1. J.H. Kastner, D.P. Huenemoerder, N.S. Schulz, C.R. Canizares, D.A. Weintraub, ApJ 567, 434 (2002) 2. J.H.M.M. Schmitt, J. Robrade, J.U. Ness, F. Favata, B. Stelzer, A&A 432, L35 (2005) 3. H.M. G¨unther, C. Liefke, J.H.M.M. Schmitt, J. Robrade, J.U. Ness, A&A 459, L29 (2006) 4. C. Argiroffi, A. Maggio, G. Peres, A&A 465, L5 (2007) 5. J. Robrade, J.H.M.M. Schmitt, A&A 473, 229 (2007) 6. H.M. G¨unther, J.H.M.M. Schmitt, J. Robrade, C. Liefke, A&A 466, 1111 (2007) 7. G.G. Sacco, C. Argiroffi, S. Orlando, A. Maggio, G. Peres, F. Reale, A&A, in press (2008) 8. G.B. Field, ApJ 142, 531 (1965)

MRI and Outflows: Angular Momentum Transport in Protoplanetary Disks Raquel Salmeron

Abstract Angular momentum in protostellar disks can be transported radially by MHD turbulence induced by the magnetorotational instability (MRI); or vertically, via magnetocentrifugal outflows. We have studied these processes under realistic fluid conditions and explored the possibility that they operate at the same radial location. The results suggest that they may be unlikely to coexist in real disks.

1 Formulation Magnetocentrifugally driven winds are likely to be the dominant angular momentum transport mechanism in protostellar disks when the magnetic field is strong (so that the ratio a of the Alfv´en speed to the sound speed is .1; [2, 4]). On the other hand, the MRI [1] is thought to dominate when a 1 [e.g. [6, 7]]. Both mechanisms are likely play a role in real disks, but previous studies have considered them

R. Salmeron () Research School of Astronomy & Astrophysics and Research School of Earth Sciences, The Australian National University, Canberra, Australia e-mail: [email protected]

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separately and under simplified assumptions. We have studied these processes under more realistic conditions, incorporating the ionization structure of the disk and exploring the possibility that they operate at different heights of the same radial location. We used the methodology in [9] to obtain radially localised disk-wind solutions from strongly magnetised disks. They are found to satisfy 2a2 > 1 and  & 1 (where a / 1=2 and  measures the field-matter coupling). We then developed a prescription to incorporate radial transport in our wind models, using the fact that when the condition 2a2 > 1 is violated, the surface layers of the disk become super-Keplerian [9]. This is associated with outward streaming motion that is unphysical for a pure wind solution, but is characteristic of the two-channel mode that underlies the MRI. We, therefore, differentiate the MRI-unstable section of the disk (taken to be where 2a2 < 1) from the region where only vertical angular momentum transport occurs (2a2 > 1), quantify the turbulent transport expected to develop in the former [3] using the results of published numerical simulations [8], and add this term to the disk angular-momentum conservation equation (see [5]).

2 Results Figure 1 displays a pure wind solution (solid lines) and how it is modified when radial angular-momentum transport is incorporated via the prescription outlined above (dashed lines). Note that the addition of radial transport results in an increased inflow speed (") and a higher mass inflow rate (MP in ). Although the disc thickness is unchanged and the height of the sonic point decreases, the stronger density stratification leads to a lower outflow rate and vertical torque in the combined solution. We are using this scheme to estimate the fractions of angular momentum transported radially and vertically in protoplanetary disks, and assess the implications of the resulting disk structure to other processes (e.g. planet formation and migration).

Fig. 1 Density, magnetic field (left) and velocity components (right) for a pure wind solution (solid lines) and how it is modified when radial transport is included where 2a2 < 1 (dashed lines). The curves terminate at the respective sonic points. The parameters are (1) a0 , the magnetic field strength; (2)  D 0:6  0:04z= hT , the magnetic coupling; (3) ", the inward radial speed at z D 0; (4) cs =vK , the disc geometric thickness; and (5) "B , the azimuthal component of the electric field

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References 1. Balbus S., Hawley J., 1998, Rev. Mod. Phys., 70, 1 2. Blandford R., Payne D., 1982, MNRAS, 199, 883 3. Goodman J., Xu G., 1994, ApJ, 432, 213 4. Pudritz R., Ouyed R., Fendt Ch., Brandenburg A., 2007, in Reipurth, V., Jewitt, D., Keil, K. eds, Protostars & Planets V. Univ. Arizona Press, Tucson, p. 277 5. Salmeron R., K¨onigl A., Wardle M., 2007, MNRAS, 375, 177 6. Salmeron R., Wardle M., 2005, MNRAS, 361, 45 7. Sano T., Stone J., 2002, ApJ, 577, 534 8. Sano T., Inutsuka S., Turner N., Stone J., 2004, ApJ, 605, 321 9. Wardle, M., K¨onigl, A. 1993, ApJ, 410, 218

Analysis of the Central X-ray Source in DG Tau ¨ P. Christian Schneider and Jurgen H.M.M. Schmitt

As a stellar X-ray source DG Tau shows two rather unusual features: A resolved X-ray jet [2] and an X-ray spectrum best described by two thermal components with different absorbing column densities, a so called “two-absorber X-ray (TAX)” morphology [1, 2]. In an effort to understand the properties of the central X-ray source in DG Tau a detailed position analysis was carried out. Our work is based on the archival Chandra data of DG Tau (90 ks in total) described in detail in [2]. We use the same energy ranges for the soft (0.3–1.1 keV) and the hard spectral component (1.7–7.0 keV) as in [2]. The excellent angular resolution of the Chandra telescope allows to measure relative positions with an accuracy of 0:1 arcsec. We tentatively identify the hard, strongly absorbed, X-ray component with coronal emission of DG Tau. The TAX property of the spectrum suppresses the mutual contamination of the components almost completely. If the

P.C. Schneider () Hamburger Sternwarte, Gojenbergsweg 112, 21029 Hamburg, Germany e-mail: [email protected] J.H.M.M. Schmitt Hamburger Sternwarte e-mail: [email protected]

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Fig. 1 The derived separations of the two components. All four observations are aligned so that the hard component is located at (0, 0). The grey areas indicate the 95% error ranges. Note that the position of the soft component might vary in the course of time

soft component is emitted not from the star itself but rather from – say – a shocked wind, one might find a position difference between the soft and the hard X-ray components. In Fig. 1 we show the derived separations between the soft and the hard X-ray centroids. All observations reveal an offset between both central X-ray components at the order of 0.2 arcsec 45 AU. The position angles of the offsets differ only slightly from the one of the optically observed jet. In [3] we showed that the offset is neither caused by statistical fluctuations nor by instrumental effects. None of the almost 30 point-like comparison sources with a compatible data quality showed an offset as large as the average separation found between the soft and the hard photon centroids in DG Tau. We therefore believe that the offset of 0:2 arcsec is physical and cannot be attributed to an instrumental or statistical effect. Assuming that the soft X-ray emission is produced by singly shocked material within the inner part of the jet, we can estimate the required mass-loss by MP Xray ' mH .T / EM=.3kB T /;

(1)

with mH : Hydrogen mass, .T /: cooling function, EM : emission measure, kB T : thermal energy density. The derived value is orders of magnitude smaller than the total mass-loss. If we further assume a cooling distance of 0.5 arcsec, the density of the X-ray emitting material turns out to be 106 cm3 . Our analysis shows that the X-ray emission arises co-spatial with the emission in forbidden emission line regions and probably has compatible densities. This favors an interpretation of the X-rays as being produced by internal shocks, which are also the favored heating mechanism of the optically observed emission [4]. However, only a small fraction of the outflowing material is needed to explain the observed soft X-rays.

Analysis of the Central X-ray Source in DG Tau

References 1. G¨udel M., Telleschi A., Audard M. et al. 2007, A&A 468, 515–528 2. G¨udel M., Skinner S. L., Audard M. et al. 2008, A&A, 478, 797–807 3. Schneider P. C., Schmitt J. H. M. M. 2008, A&A, 488, L13–L16 4. Lavalley-Fouquet C., Cabrit S. and Dougados C. 2000, A&A 356, L41–L44

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Verification of Candidate Protostellar Outflows in GLIMPSE Bringfried Stecklum, Alessio Caratti o Garatti, Chris Davis, Hendrik Linz, Thomas Stanke, and Hans Zinnecker

Abstract Using the 4.5 m excess as a tracer of shocked H2 , we identified 160 candidate protostellar outflows in the Spitzer GLIMPSE survey. In order to verify their nature, a follow-up campaign for H2 1–0 (S1) 2.12 m imaging was initiated. Here we present first results for southern objects obtained with ESO-SOFI. About half of the observed targets could be detected at 2.12 m. The non-detections point to large dust column densities and/or different excitation conditions. The comparison with the IRAC images reveals the diversity of the emission.

1 Scientific Objective, Observations and Data Reduction Jets and bipolar outflows are clear signs of stellar growth by accretion from circumstellar disks [1]. This holds for stars of .8 : : : 10 Mˇ [2] while the formation mechanism of high-mass stars still needs investigation [3]. For this aim the GLIMPSE survey is well suited since it permits to identify very young massive stars [4]. Excess emission at 4.5 m hints at shock-excited H2 [5] and was used by us to establish a sample of 160 candidate outflows. The majority coincides with methanol/and or water masers which are signs of luminous YSOs. SOFI 2.12 m and Ks imagery for

B. Stecklum () and A. Caratti o Garatti TLS Tautenburg, Sternwarte 5, 07778 Tautenburg, Germany e-mail: [email protected] C. Davis JACH, Hilo, Hawaii, USA H. Linz MPIA, K¨onigstuhl, Heidelberg, Germany T. Stanke ESO, Garching, Germany H. Zinnecker AIP, Potsdam, Germany

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32 targets was obtained in March 20081 . After basic image processing the astrometry was established using 2MASS positions. Before the mosaicing the continuum was subtracted for each tile separately after scaling with the 2.12 m to Ks total flux ratio. The photometric calibration is based on SEXTRACTOR magnitudes, with the zero point tied to 2MASS, and accounting for the different filter band widths. The median 3 line-flux sensitivity is 2  105 erg/s/cm2 /sr for the central tile.

2 Results and Discussion For about two half of the imaged regions H2 emission could be detected. Generally, it is concentrated towards the prime target but for a few fields distant weaker flows were found as well. The peak line fluxes are in the range of 0:7 : : : 3  104 erg/s/cm2 /sr. In most cases the emission has a complex morphology which is typical for outflows from massive YSOs. The comparison of the IRAC and SOFI images revealed a large diversity which can be attributed to spatially varying extinction and excitation conditions (shocks vs. fluorescence). Fig. 1 displays a typical example. Spectroscopy of the jet knots is foreseen to derive their physical properties

Fig. 1 IRAC image of Caswell OH 332.35200.117 (cross) based on 5.8, 4.5, and 3.6 m frames (left) and continumm-subtracted SOFI image (right). The central H2 emission is not seen at 2.12 m presumably due to large extinction. Its morphology resembles a photo-dissociation region rather than an outflow. However, the 2.12 m knots likely trace a bended jet from the YSO. A color version of this plot can be found in the electronic version of this book and in Appendix A (Fig. A.39)

1 Based on observations collected at the European Organisation for Astronomical Research in the Southern Hemisphere, Chile, Prop. Id. 082.C-0130.

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using the approach from [6]. Kinematic distances will be applied in tandem with data from the GLIMPSE/MIPSGAL/BOLOCAM catalogues to derive luminosities and assess the possible masses of the driving sources.

References 1. B. Reipurth and J. Bally, ARA&A, 39, 403 (2001). 2. H. G. Arce, D. Shepherd, F. Gueth et al., in PPV, eds. B. Reipurth, D. Jewitt & K. Keil (2007), pp. 245–260. 3. H. Zinnecker and H. W. Yorke, ARA&A, 45, 481, (2007). 4. E. Churchwell, B. Babler, B. Benjamin et al., in 36th COSPAR Sci. Assembly, 146 (2006). 5. A. Noriega-Crespo, P. Morris, F. R. Marleau et al., ApJS, 154, 352 (2004). 6. A. Caratti o Garatti, T. Giannini, B. Nisini, & D. Lorenzetti, A&A, 449, 1077 (2006).

Young Stellar Jets and Outflows in the Massive Star Forming Complex W5 Guy S. Stringfellow, John Bally, and Adam Ginsburg

Abstract W5 is an active region of star formation located at 1.9 kpc in the Perseus spiral arm within Cassiopeia. We have surveyed W5 in the optical (H˛ , [S II], and i filters), and at 1.1 mm. Our surveys cover several square degrees. Several highly collimated stellar jets have been discovered, along with other interesting outflows and Herbig-Haro objects. We have obtained optical spectra with the ARC APO 3.5 m telescope at various position angles for several of these jets and outflows, revealing jet velocities of 200 km s1 .

1 Introduction W5 is an active region of star formation located at 1.9 kpc [1, 5], in the Perseus spiral arm within Cassiopeia. The western region of W5 (W5W) is defined by a large 1.5ı diameter HII region rimmed by bright molecular clouds (e.g., [2]). The eastern component W5E is similarly characterized, extending about 1ı in diameter [2]. W5E is predominately ionized by a single O7 star, while W5W has several “O-stars”, each being multiple O-star systems, that together ionize the region [1]. We have surveyed W5 in the optical (H˛ , [S II], and i filters) using the wide-field Mosaic I camera on the KPNO 4 m, and dust continuum emission at 1.1 mm using Bolocam on the CSO on Mauna Kea [3, 4]. These surveys cover several square degrees (and include other nearby star forming complexes). The narrow-band imaging survey of the W5 HII region shows it is actively forming new stars. Our observations have discovered over a dozen candidate protostellar jets and bipolar outflows. We are conducting a program of visual- wavelength spectroscopy to measure their radial velocities to confirm their Herbig-Haro nature.

G.S. Stringfellow (), J. Bally, and A. Ginsburg Center for Astrophysics and Space Astronomy, Department of Astrophysical and Planetary Sciences, University of Colorado, Boulder, Colorado e-mail: [email protected]

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2 Observations The wide-field imaging survey conducted in narrow-band optical H˛ and [S II] filters with the Mosaic I camera on the KPNO Mayall 4 m telescope has revealed the presence of numerous jets and outflows in the W5 massive star forming region. When these optical images are correlated with the BGPS 1.1 mm CSO data, these jets are found to be closely associated with the bright cold star forming dust cores. Figure 1 displays one of the highly collimated jets discovered from the optical imaging survey. The jet emanates from the tip of an opaque finger-column in which the young forming star is embedded. Inspection of the optical images suggests the jet emanates from both sides of the star, though the red-shifted component is the more extended. Optical spectroscopy confirms this geometry (see below). At 1.9 kpc distance, the optically visible extent of the red-shifted component of the jet extends beyond 600 AU. Optical spectra have been obtained with the Dual Imaging Spectrograph on the ARC Apache Point Observatory 3.5 m telescope at various position angles for several of these jets and outflows. The instrument setup utilized provides an angular ˚ per pixel, corresponding to resolution of 0.4200 per pixel and a dispersion of 0.56 A a velocity resolution of 26 km s1 near H˛ . Figure 1 also shows the slit alignments for the two positions – aligned along the jet, and an orthogonal cut across the jet. The resulting 2-D spectra near H˛ are also shown in Fig. 1. The spectrum with the slit aligned along the jet shows a continuous accelerating jet originating at or close to the star itself, and achieving a maximum velocity 200 km s1 . The spectrum where the slit crosses the outer region of the jet displays a narrow velocity range near the maximum jet velocity. Similar velocities and profiles are also seen in the [S II] lines redward of H˛ .

Fig. 1 Left two panels: H˛ image of a new collimated jet in W5. The young star producing the jet is embedded at the tip of the finger-column. The slit is aligned parallel along the jet, and the resulting 2-D spectrum is shown in the adjacent panel to its right. Right two panels: Slit now makes a cut across the jet away from the star. The resulting 2-D spectrum is shown in the adjacent panel

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3 Conclusions Imaging surveys of dust continuum emission at 1.1 mm and optical emission lines (H˛ , [S II]) of the massive star-forming region W5 have revealed a strong correlation between heated dust continuum emission radiating from cold prestellar cores with other young stellar objects (YSOs) forming in these regions. These YSOs are identified by the visible jets and outflows they produce, while optical spectroscopy confirms that these features are indeed jets or outflows, rather than filamentary nebulosity within the intra-cloud medium that is often also found in star forming regions. Acknowledgements GSS would like to thank the organizers and the AAS for financial support that enabled his attendance at the meeting. Observations were obtained with the NOAO-KPNO Mayall 4m telescope, operated by AURA with support through NSF, and the Apache Point Observatory 3.5 m telescope, owned and operated by the Astrophysical Research Consortium.

References 1. Hillwig, T.C., Gies, D.R., Bagnuolo, Jr., W.G., Huang, W., McSwain, M.V., & Wingert, D.W.: Binary and Multiple O-Type Stars in the Cassiopeia OB6 Association. ApJ 639, 1069–1080 (2006) 2. Karr, J.L., & Martin, P.G.: Triggered Star Formation in the W5 H II Region. ApJ 595, 900–912 (2003) 3. Stringfellow, G.S., Bally, J., & Bolocam Galactic Plane Collaboration: Bolocam 1.1mm and Optical Observations of the W5 Star Forming Region in the Perseus Arm. BAAS, 40, 198 [#7.06] (2008) 4. Stringfellow, G.S., Bally, J., & Allen, L.: Irradiated Cometary Gas-Free Dust Tails and New Stellar Jets in the Massive Star Forming Complex W5. BAAS, 39, 1006 [#154.04] (2007) 5. Xu, Y., Reid, M.J., Zheng, X.W., & Menten, K.M.: The Distance to the Perseus Spiral Arm in the Milky Way. Science, 311, 54–57 (2006)

Water Masers and Radio Continuum Emission Tracing Thermal Radio Jets M.A. Trinidad

Abstract We present interferometric observations of radio continuum and water maser emission carried out with the Very Large Array (VLA1 ) toward five massive star-forming regions. Radio continuum sources were detected toward all starforming regions and some of them could be thermal radio jets. In addition, although water maser emission was found in all massive star-forming regions, the water masers are not spatially associated with all detected continuum sources. Based on the analysis of the distribution of water masers and the characteristics of the continuum emission, we suggest that some water masers are tracing jets.

1 Introduction At present, only few luminous protostars show evidence of thermal radio jets (e.g. Cepheus and HH 80–81; [1,2]). Then, it is not clear whether high-mass outflows are driven by collimated jets or if another, different mechanism is required. Therefore, studies of individual high-mass protostars are essential to address this issue.

2 Results and Discussion We present radio continuum and water maser observations carried out with the VLA in its A configuration during several runs toward five star-forming regions. Main results of the observations are given in Fig. 1. Radio continuum emission was detected

M.A. Trinidad () Department of Astronomy, University of Guanajuato e-mail: [email protected] 1 The NRAO is operated by Associated Universities Inc., under cooperative agreement with the National Science Foundation.

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Fig. 1 Multiples sources are common in massive star-forming regions. In addition, binary and triple systems are identified, which suggest that high-mass stars are not formed isolated, instead, they are formed in small groups. In addition, several radio continuum sources, associated with highmass protostars, could be interpreted as thermal jets. The relative positions of the water masers (circles and crosses) with respect to the radio continuum emission have been established with ten milliarcseconds of accuracy. Almost all water maser features are spatially associated with some known radio continuum source and that a few maser features seem to be isolated (see [3, 4, 5, 6, 7] for details)

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toward all regions. Moreover, some of them could be interpreted as thermal jets, which are aligned with the large scale molecular outflows. Interferometric water maser observations are a powerful tool to study the circumstellar regions of the young stellar objects. In this way, we have studied the spatio-kinematical distribution of the water masers and their relation with the radio continuum emission with an angular resolution of 0.100 in all massive star-forming regions. We find that water masers are tracing the base of the large scale bipolar outflows in almost all massive star-forming regions. Then, based on the analysis of the distribution of water masers and the characteristics of the continuum emission, we suggest the thermal jets are common in massive star-forming regions. Acknowledgements M.A.T. acknowledges the support from CONACyT grant 46157-E and Programa de Fortalecimiento Institucional from Universidad de Guanajuato.

References 1. Mart´ı, J., et al. 1999, A&A, 345, L5 1918 2. Rodr´ıguez, L.F., et al. 1994, ApJ, 430, L65 3. Trinidad, M.A., et al. 2003, ApJ, 598, 386 4. Trinidad, M.A., et al. 2004, ApJ, 613, 416 5. Trinidad, M.A., et al. 2005, AJ, 130, 2206 6. Trinidad, M.A., et al. 2006, AJ, 132, 1918 7. Trinidad, M.A., et al. 2007, AJ, 134, 1870

Effects of Flaring Activity on Dynamics of Accretion Disks in YSOs Tatiana G. Yelenina, Salvatore Orlando, Fabio Reale, Giovanni Peres, Andrea Mignone, and Titos Matsakos

Abstract We investigate the effects of strong flares on the accretion phenomena in YSOs. Among all classical assumptions, the model accounts magnetic-fieldoriented thermal conduction. We study the global dynamics of the system for two positions of the heating release triggering the flare.

T.G. Yelenina () and S. Orlando INAF Osservatorio Astronomico di Palermo, Italy, Piazza del Parlamento 1, 90134 (PA) e-mail: [email protected]; [email protected] F. Reale and G. Peres INAF Universit`a di Palermo, Italy, V.le delle Scienze 90128 (PA) e-mail: [email protected]; [email protected] A. Mignone and T. Matsakos Universit`a di Torino, Italy, Via P.Giuria 1, 10125 (TO) e-mail: [email protected]; [email protected]

K. Tsinganos et al. (eds.), Protostellar Jets in Context, Astrophysics and Space Science Proceedings, DOI 10.1007/978-3-642-00576-3 97, c Springer-Verlag Berlin Heidelberg 2009 

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Fig. 1 Distribution of (upper panels, log scale) and T (lower panels, log scale). Left four frames: flare at the stellar surface. Right four frames: flare on the disk. Magnetic lines are shown by solid lines. Arrows show velocity field. A color version of this plot can be found in the electronic version of this book and in Appendix A (Fig. A.40)

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1 Problem Statement We consider a rotating star with R D 2Rˇ and M D 0:8Mˇ surrounded by Keplerian disk with the inner radius at 5R . The magnetic field of the star is dipolelike and aligned with rotational axis, its intensity is 103 G at star’s surface. The initial conditions are chosen to be in mechanical equilibrium [3]. At the boundary we consider symmetry conditions at rotation axis and equatorial plane and free flow conditions at the radial boundaries. Conventional form of axi-symmetrical MHD equations is used with energy equation extended with thermal conduction (including heat flux saturation) [2], and external heating term. The heating release, triggering the flare, is assumed located either close to the stellar surface or on the disk. The flare lasts for 5 min. The total released energy 3 1035 erg is in agreement with observations [4]. The implementation is based on the MHD code PLUTO [1].

2 Results and Discussion For heating release close to the stellar surface, the flare causes a rapid increase of the temperature in corona to typical flare values of 107 K. After 0.65 days, the conduction front reaches the inner portion of the disk starting to heat it, leading to its rapid expansion and, as a result, to an accretion flow in the equatorial plane lasting for 1 day. For heating release on the disk’s surface, the flare determines a rapid increase of the temperature up to the 6 MK. The shock and conduction fronts propagating downward to the disk heat it, leading to the strong injection of the disk’s matter in the corona. We note an accretion flow forming at the equatorial plane. We conclude that flaring activity can contribute significantly to the stellar accretion. Acknowledgements The authors are grateful to the support of PHOENIX Marie Curie Host Fellowship Programme on YSOs, their Surroundings and Jets, contract No. MTKD-CT-2005-029768.

References 1. A. Mignone, G. Bodo, S. Massaglia et al, ApJS, 170, 228–242 (2007) 2. S. Orlando, F. Bocchino, F. Reale et al, ApJ, 678, 274–286 (2008) 3. M.M. Romanova, G.V. Ustyugova, A.V. Koldoba et al, ApJ, 578, 420–438 (2002) 4. S.J. Wolk, F.R. Harden Jr., E. Flaccomio et al, ApJS, 160, 423–449 (2005)

Index

[FeII] lines, 216, 487 H2 2.122 m line, 487

accretion, 34, 515, 519 accretion disk, 4, 137, 154, 166, 186, 310, 448, 498, 519, 552 accretion phenomena, 631 accretion shocks, 39 magnetospheric accretion, 35 Adaptive Optics, 232 AGN, 186, 515 ambi-polar diffusion, 454 ambient density, 196 ambient gas, 454 ambient medium, 195 ambipolar diffusion, 63, 70 AMR, 448 analytical solutions, 205, 210, 441 angular momentum, 132, 236, 242 angular momentum transport, 515 astrophysical jets, 112, 205 atomic, 478

Baroclinic instability, 59 BDF scheme, 455 BE technique, 234, 577 Bernoulli equation, 206 bipolar outflow, 623 black hole candidates, 186 Blandford-Znajek effect, 64 bow shock, 113, 456, 524, 568 bremsstrahlung, 449 brown dwarf, 259

C-shocks, 454, 524 causality, 210 centrifugally driven winds, 67

Chandra, 615 channel maps, 233 chemical network, 455 chemistry, 455 Class 0, 242 Class 0 jets, 221 Class 0 objects, 215 Class I, 242 Class I objects, 215 Classical T Tauri stars, 154, 543, 577 cocoon, 398, 454 collimation, 132 collisional excitation, 455 collisional ionisation, 455 compact objects, 515 conical outflows, 155 contact discontinuity, 456 cooling, 351 atomic cooling, 455 cooling processes, 448 cooling zones, 324 cores, 597 corotation radius, 168 CRL 618, 603 current driven instability, 290 CW Tau, 334, 349 cyclotron, 322

DG Tau, 334, 348, 543, 577, 615 jet rotation, 235 jet with AO, 233 jet with HST/STIS, 233 diagnostic lines, 486 diffusivity, 132 direct numerical simulation, 515 disk magnetic field, 133 disk wind, 166, 552 disk–wind connection, 69 disk-magnetosphere boundary, 154

635

636 Doppler shifts, 451 DP Tau, 349 dust, 455 continuum emission, 623 optical properties, 460 sublimation, 460 temperature, 460 dust evaporation radius, 310 dust reprocessing, 306 Dust streaming instability, 59 dynamical age, 396

echelle spectroscopy, 523 effective magnetic Prandtl number, 118 electron density, 487 electron fraction, 456 energy equation, 448 episodic, 196 excitation temperature, 400 experiments, 195, 360, 567

FEL regions, 219 filling factor, 351 flare, 135, 631 footpoint radius, 236 forbidden emission lines, 112, 260, 553 maps, 115 regions, 487 FU Orionis outbursts, 69

gamma-ray bursts, 4 gas depletion, 305 global disk models, 64 Gravitational instabilities, 59 GRMHD simulation, 589 growth rates, 467 GV Tau, 349 gyrosynchrotron, 322

H2 lines, 218 Hall effect, 63 Hamiltonian, 319 Herbig-Haro objects, 448, 512, 623 HH 30, 234, 235 HH211, 454 HH212, 523 HH240/241, 396, 397 HH30 jet, 314 HH34, 486, 487 HH46/47, 454, 486

Index HH 1/2, 330, 333 HH 110, 335 HH 111, 330 HH 32, 331 HH 34, 330, 333 HH 46/47, 330 HH 52, 336 HH 54, 335, 337 HN Tau, 349 hoop stress, 236 HST, 233, 448 Hubble Space Telescope, 233 HVC, 487 hydrodynamic jet, 196

imaging, 623 infrared, 486 infrared: ISM, 330 instabilities, 59, 548 integral field spectrograph, 478 intensity maps, 448 ion source term, 448 ionisation, 456 ionisation network, 448 IR lines, 216 IR spectral diagnostic, 216 IRAS 05358+3543, 272 IRAS 11101-5829, 272 IRAS 16547-4247, 272 IRAS 18151-1208, 272 IRAS 20126+4104, 267, 269 ISAAC, 486

J-shocks, 454 stationary, 457 jet-driven outflows, 454 jet-emitting disk, 123, 519 jets, 154, 205, 285, 478, 515, 519, 552, 597 asymmetries, 525 bending, 335 chemistry, 459, 596 deflection, 335 diagnostics, 312 excitation conditions, 234 heating/cooling, 460, 596 ionization, 461, 596 kinematics, 267 kinematics and dynamics, 329 knots, 135, 523 molecular, 459, 595 physical properties, 486 precession, 335

Index radiation field, 461, 462 rotation, 235, 241, 244, 245, 333, 516, 524 rotation survey, 236 shocks, 460 stability, 465 stellar, 451 temperature, 461, 596 variability, 335 width, 125 jets and outflows, 186, 205, 268 jump conditions, 457

Kelvin-Helmholtz instability, 59, 285, 332, 398, 465, 547, 590 Keplerian frequency, 43 kinematics, 486 kink instability, 188 knots, 442

laser, 567 lateral bow-shocks, 456 launching region, 478 line broadening, 451 line of sight, 450 Lk H˛ 233 jet with HST/STIS, 234 LkH˛ 233 jet with AO, 233 LVC, 487

Mach disk, 524 magnetic acceleration, 205 magnetic acceleration efficiency, 207 magnetic backbone, 115 magnetic braking, 70 magnetic bubbles, 196 magnetic cavity, 195 magnetic fields, 205, 317, 515, 589 advection, 118 enhancement, 469 magnetic flux, 119 magnetic islands, 468 magnetic pressure, 324 magnetic tower, 190, 195 magnetically threaded disks, 70 Magnetised Accretion-Ejection Structures, 520 magnetization, 134 magneto-centrifugal wind, 64 magneto-centrifugally acceleration, 112 magneto-rotational instability, 187 magneto-thermal instability, 59

637 magneto-viscous instability, 59, 61 magnetospheric accretion, 35, 51 magnetospheric ejections, 169 masers, 321 mass ejection flux, 489 mass flux rates, 219 mass loss, 34 mass-loss rate, 135, 480 mass-velocity relation, 398, 401 MHD, 124, 132, 186, 205, 448, 455, 498, 515, 547, 597, 633 jet formation, 132, 205 MHD conserved quantities, 206 MHD simulations, 112, 154, 186, 465, 552 MHD turbulence, 60 MHEL regions, 220 MINEq, 448 model, 544 molecular, 478 molecular cloud turbulence, 395 molecular cooling, 455 molecular core, 454 molecular emission, 454 molecular hydrogen, 455, 523 molecular outflows, 289, 292, 408, 454 momentum equation, 206 MRI, 58, 59, 515 multifluid MHD, 454, 547 multipole stellar fields, 51 MYSO, 267

NGC 2264 G outflow, 573 NIR, 486 non-ideal MHD, 63 non-linear shear instability, 59 numerical codes, 448 numerical simulations, 498, 531

Ohmic resistivity, 63 onion-like kinematics, 234 optical jets, 190, 408 outbursts, 374 outer disk radius, 123 outflows, 5, 154, 205, 285, 478, 498, 624

Parker instability, 59 parsec scale flows, 334 Planetary Nebulae hydrodynamics, 602 outflows, 601 shock waves, 602

638 plasma jets, 195, 567 PLUTO, 124, 448, 455, 515 PLUTO code, 498 polarization, 321 position analysis, 615 position-velocity, 450 power law index, 322 precessing jet, 573 propeller regime, 154 proto-planetary nebulae jets, 603 protostars, 186, 196, 215, 267 protostellar jets, 242, 267, 329, 447, 486, 552, 623 protostellar outflow, 523 PSF, 450 PV diagrams, 232, 306, 396, 487

radial foil, 195 radially self-similar solution, 124 radiative cooling, 449 radiative recombination, 455 radiative shocks, 360 radiative transfer, 360 radiatively cooled, 195 rate equations, 455 Rayleigh-Taylor instability, 59 reconnection, 132 recurrent novae, 374 relativistic jets, 589 retrograde rotation, 517 RMHD simulation, 590 RS Ophiuchi, 373 RW Aur, 334 jet rotation, 235 jet with HST/STIS, 234 RY Tau jet, 233

shearing box approximation, 60 shock waves, 324 shocks, 196, 306, 330, 450 simulations, 360, 442 SiO emission, 221 slitless spectroscopy, 235 source term, 449 Space Telescope Imaging Spectrograph, 233 spectral diagnostics, 305 spectro-imaging, 478 spectroastrometry, 259 spectroscopy, 623 spine-sheath structure, 590 stability, 124, 498

Index stability of disk/wind systems, 71 standard accretion disk, 119 star formation, 512, 623 stars class 0, 595 class I, 595 class II, 595 radiation field, 459 steady state, 498 stellar black holes, 515 stellar jets, 305 stellar magnetic field, 131 stellar winds, 167 stiff system, 455 STIS, 233 SU Aur, 35 survey, 623 sweeping magnetic twist, 58 synchrotron, 322 synthetic observations, 449 emission maps, 125, 532 spectra, 448 Sz 102, 348

T Tau, 36 T Tauri jets, 115 T Tauri-S, 324 Th 28, 334 thermal conduction, 631, 633 thermal instability, 288 time variability, 531 Tomographic techniques, 312 toroidal magnetic field, 195 transfield equation, 208 truncation, 124 TTSs, 232 turbulence, 515 turbulent magnetic resistivity, 118 turbulent viscosity, 118 TW Hya, 36 two-component jet, 441 two-component outflow, 133 T Tauri phase, 43

velocity, 623 velocity variable jet, 325 vertical angular momentum transport, 67 VLA 05487, 397

W5, 623 weakly ionized protostellar disks, 71

Index

639

wind-driving disks, 67 winds, 154, 286 chemistry, 459 disk, 459, 595 photoreactions, 461

young stars, 154 young stellar objects, 515 YSO, 267 YSOs, 631

X-ray, 348, 543, 615 flares, 189 jet, 615 X-wind, 51, 134, 154, 400

Zeeman effect, 319 ZEUS, 132 ZEUS-MP, 112

Appendix A

Color Figures

Fig. A.1 J. Bally, Fig. 1: Irradiated jets in the Carina Nebula illuminated by the Trumpler 14 cluster

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Fig. A.2 J. Bally, Fig. 2: The HH 666 irradiated jet emerging from dust pillars located southwest of  Car. The image was obtained with the ACS camera on the Hubble Space Telescope (see Smith, Bally, & Brooks 2004)

Wind from disk Magnetically-heated accretion columns add angular momentum coronal loops? dusty disk

gas disk inside dust evaporation radius

Connection to outer diskoutside co-rotation?

Fig. A.3 L. Hartmann, Fig. 1: Schematic view of disk-stellar magnetosphere interaction. The magnetosphere truncates the disk at some point, probably inside the radius at which dust sublimates, and then material accretes supersonically along magnetic flux tubes onto the star, where it creates excess continuum emission as it shocks. The accreting material adds angular momentum to the star, which must be taken away to prevent spinup, either by coupling to the disk outside of corotation and/or driving a wind (the stellar wind is probably unimportant)

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Fig. A.4 C. M. Johns–Krull, Fig. 3: Emission line profiles from TW Hya testing the wind hypoth˚ observed with HST STIS. The black curve esis. The upper panel shows the C II doublet at 1,335 A is the observed profile and the red curve is the observed profile shifted so that the red member of the doublet overlays the blue member. The doublet lines should have very similar shapes and the figure shows that the red wing of the blue member is significantly extincted by the wind component in the blue wing of the red member of the doublet. The middle panel shows the same test for the higher temperature C IV line. The two members have the same shape, indicating no such wind absorption at the characteristic temperature of C IV. Taken from Johns–Krull & Herczeg (2007)

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Fig. A.5 J.M. Stone, Fig. 2: Mechanism of the MRI. See the text for details

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Fig. A.8 J. Ferreira, Fig. 3: Star-disc interaction where the stellar magnetic moment is parallel to the disc magnetic field. There are three distinct types of ejection: a stellar wind on the axis, a disc wind (shown in colors) and intermittent bullets launched at the interface (Reconnection X-wind), braking down the protostar and channeled by the outer disc wind

Fig. A.9 M. Stute et al., Fig. 1: Structure of the flow (logarithmic density plots) for models SC1a–SC1e (from left to right) at timesteps t D 0 (top), t D 25 (middle) and t D 50 t0 (bottom). Also plotted is the magnetic field line anchored in the lower boundary where ˛ D ˛trunc (white line). At t D 50 t0 , we also plot the two field lines with half and twice the radius of that of the truncation field line used for investigating the integrals of motion

646

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Fig. A.10 M. Stute et al., Fig. 2: Jet shape derived from synthetic [SII] images as a function of distance from the source in the untruncated model ADO (top left) and the truncated models SC1f (top right) and SC1g (bottom right); overlaid are data points from CFHT/PUEO and HST/STIS observations of DG Tau (diamonds), HN Tau (plus signs), CW Tau (squares), UZ Tau E (crosses), RW Aur (circles), HH 34 (one triangle), HH 30 (black solid line) and HL Tau (red dashed line); data are taken from [14] for distances below 200 AU and [6] beyond this distance

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Fig. A.11 M. Cemeljic et al. Fig. 1: The initial setup, which is slightly modified analytical solution, is shown in the Left panel. The solid lines represent logarithmically spaced isocontours of density. It is also shown in color grading, in red to violet color, for the logarithm of density 1 to 4, respectively. In the Right panel shown is, in the same grading, the solution with large magnetic diffusivity. It does not reach stationary state, and shows some periodicity in time evolution. The dashed lines depict poloidal magnetic field lines, and the dotted lines depict the fast magnetosonic, Alfven and slow magnetosonic critical surfaces, top to bottom, respectively

Fig. A.12 M. Romanova et al., Fig. 1: Snapshots from axisymmetric simulations of conical winds. The background shows the matter flux with light color corresponding to higher flux. The lines are magnetic field lines

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Fig. A.13 M. Romanova et al., Fig. 2: Typical flow in conical winds (at t D 380 days). The background shows matter flux, lines are selected field lines, arrows are proportional to velocity. The numbers show poloidal vp and total vt ot velocities and number density at sample places of the simulation region

Fig. A.14 M. Romanova et al., Fig. 3: Two components of winds from slowly rotating star are labeled

Fig. A.15 M. Romanova et al., Fig. 5: Conical winds obtained in 3D MHD simulations for  D 30ı . Left panel: density distribution and sample field lines in the ˝ – plane. Right panel: same but in the perpendicular plane

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Fig. A.16 M. Romanova et al., Fig. 7: Outflows in the propeller regime. The background shows matter flux, lines are selected field lines, arrows are proportional to velocity. Labels show total velocity and density at sample points

Fig. A.17 M. Romanova et al., Fig. 8: Two components of outflows in the propeller regime

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Fig. A.18 B. Nisini, Fig. 2: Spitzer H2 S(5) map of the HH211 molecular jet superimposed over the high velocity CO J D 2  1 map of [22]. From [14]

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Fig. A.19 A. Raga et al., Fig. 1: Models with single or double-mode sinusoidal ejection variabilities. The density stratifications obtained for a t D 400 (left) and a t D 800 yr integration time (right) are shown. The frames have a horizontal extent of 4  1017 cm. The density stratifications are depicted (in g cm3 ) with the logarithmic color scale given by the bar on the bottom right. The parameters of the model M1 (top) through M5 (bottom) are described in Sect. 2.2 (also see Table 1)

652 Fig. A.20 A. Raga et al., Fig. 3: Density stratifications (corresponding to a t D 800 yr integration time) resulting from simulations with the ejection velocity time-variabilities shown in Fig. 2 (also see Sect. 2.3). The displayed domains have an axial extent of 4  1017 cm. The density stratifications are displayed (in g cm3 ) with the color scheme given by the bar on the bottom right

Fig. A.21 A. Raga et al., Fig. 4: Column density time-sequence computed from the variable + precessing jet model described in Sect. 2.4. The horizontal extent of the frames is of 4  1017 cm. The column densities are given (in cm2 ) by the bar on the top

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Fig. A.22 L. Podio et al., Fig. 1: Position–velocity diagrams of the forbidden lines ratios used in the diagnostics technique. From top to bottom panel: the [SII]6731 line intensity; the [SII]6731/6716 ratio increases with the electron density, ne ; the [NII]/[OI] ratio is mainly dependent on the ionization fraction, xe , and increases for increasing xe ; the [OI]/[SII] ratio increases for increasing temperature Te . The PV diagrams of the ratios show that ne peaks in the HVC of the most luminous knots (F, H, J and L) and a clear peak of xe and Te in the HVC of knot L. These are typical shock signatures

Fig. A.23 R. Bonito et al., Fig. 1: Difference image Hff [SII] of HH 154 as observed with HST in 2005 (adapted from [3])

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Fig. A.26 R. Bonito et al., Fig. 4: Comparison between the X-ray emission maps derived from the pulsed jet model and the ACIS observation of HH 154 in 2005

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Fig. A.27 J.H. Kastner, Fig. 1: HST optical (left) and Chandra (middle) X-ray images of NGC 7027. At upper right is the 0.3–2.5 keV spectral energy distribution extracted from the Chandra X-ray CCD data; at lower right, contours of high-velocity Br emission are superimposed on the Chandra X-ray image. Adapted from [9] and [6]

Fig. A.28 J.H. Kastner, Fig. 2: Color composites of HST and X-ray (Chandra and XMM, respectively) images of the bipolar nebulae Menzel 3 (left) and Hubble 5 (right). The HST images are color-coded blue-green, and the X-ray images are color-coded red. North is up and east is to the left in both images; the fields of view are approximately 2000  2000 (Mz 3) and 7500  7500 (Hb 5). Adapted from [10] and [13]

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Fig. A.29 A. Frank, Fig. 2: Three dimensional volume rendered realization of the logarithm of the gas density, for the continuously driven jet in a turbulent environment with semi-transparent isosurfaces rendered at D 2:33; 3:67; and 316 cm3

A Color Figures

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Fig. A.30 R. Banerjee et al., Fig. 1: Density (top) and velocity (bottom) evolution of a Mach 5 jet (run M5c) at two different times, t D 3:0 (left) and t D 5:0 (right). The jet is continuously powered and runs into a homogeneous medium. The jet develops knots from reflections off the jet edge. This structure propagates also into the ambient media. Additionally, Kelvin–Helmholtz instabilities develop at the edge of the jet. The turbulent flow outside the jet is mainly subsonic [from [3]]

6.0221e+04 yr

Boxsize 0.4 pc

6.0217e+04 yr

Boxsize 0.4 pc

Fig. A.31 R. Banerjee et al., Fig. 3: Shows the comparison of the collapse of a turbulent cloud core without (left) and with (right) mechanical feedback (i.e. collimated outflows). The snapshots are taken at an early time into the evolution (t  0:6 tff , where the free-frall time tff  105 year). Shown is the column density of the cloud and the star formation sites (i.e. sink particles marked as black dots). The overall structure in both cases is the same, differences are only noticeable at small scales

658

A Color Figures 7.5736e+04 yr

−

Boxsize 0.4 pc

Fig. A.32 R. Banerjee et al., Fig. 4: The left panel shows the velocity PDFs in the case without (red) and with (blue) outflows from protostellar regions at the time t  603 years (see also Fig. 3). The fraction of outflow powered velocity fluctuations is small compared to the overall turbulent motions. In the right panel we show the further evolved cloud where outflows are launched from protostellar regions

Fig. A.33 S. Brinkmann and M. Camenzind, Fig. 1: Toroidal velocity. Colors show the phicomponent of the velocity in units of c for different setups at different times. From left to right: dipole at 1;000 lct , dipole at 1;600 lct , sextupole at 400 lct and sextupole at 1000 lct

A Color Figures

659

Fig. A.34 S. Correia et al., Fig. 1: Overview of the inner part of the protostellar outflow HH 212 in the 1–0 S(1) line of H2 (left), and position of the slits on the first H2 northern knot NK1 (up) and southern knot SK1 (down)

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Fig. A.35 O. Dionatos et al. (poster), Fig. 1: (left) High velocity CO J D 3–2 outflow map around SMM1 (detail of the area observed) superimposed on an IRAC 8 m Spitzer image. Solid (blue) contours delineate blueshifted gas (integrated over 30 km s1 < VLSR < 0 km s1 ), while dashed (red) contours delineate redshifted gas (integrated over 18 km1 < VLSR < 1 ). Filled and dashed contours start at 0.8 K km1 with an 1 K km1 increment. (right) Mass flux versus bolometric luminosity diagram; details for the various points presented can be found in the text

660

A Color Figures

Fig. A.36 K. Cai et al., Fig. 1: left: Snapshot of log. / of the D 0:25 simulation at t D 1,900. The numbers on the colorbar is in units of logΠ=g cm3 , while the axes of the plot is in units of ri . right: The poloidal component of the magnetic field Bp at the same time. The numbers shown on the colorbar is in code units

Fig. A.37 C. McCoey et al., Fig. 1: IRAC image of the NGC 2264 G outflow region. The jet extends eastwards over 1.1 pc, assuming a distance of 800 pc [9] from VLA 2. The vertical lines mark the locations where the jet changes direction, and the horizontal arrows he location of the red and blue CO lobes [7]

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Fig. A.38 B. Loupias et al., Fig. 1: Interferogram for a 50 mg cm3 30% brominated foam density target

Fig. A.39 B. Stecklum et al., Fig. 1: IRAC image of Caswell OH 332.35200.117 (cross) based on 5.8, 4.5, and 3.6 m frames (left) and continumm-subtracted SOFI image (right). The central H2 emission is not seen at 2.12 m presumably due to large extinction. Its morphology resembles a photo-dissociation region rather than an outflow. However, the 2.12 m knots likely trace a bended jet from the YSO

662

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Fig. A.40 T. Yelenina et al., Fig. 1: Distribution of (upper panels, log scale) and T (lower panels, log scale). Left four frames: flare at the stellar surface. Right four frames: flare on the disk. Magnetic lines are shown by solid lines. Arrows show velocity field