Science with the VLT in the ELT Era (Astrophysics and Space Science Proceedings)

Science with the VLT in the ELT Era Astrophysics and Space Science Proceedings Advisory Editors: W.B. Burton, Charlot...

Author: Alan F.M. Moorwood

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Science with the VLT in the ELT Era

Astrophysics and Space Science Proceedings Advisory Editors: W.B. Burton, Charlottesville, VA, USA L.L. Christensen, Garching bei München, Germany

For other titles published in this series, go to www.springer.com/series/7395

Alan Moorwood Editor

Science with the VLT in the ELT Era

Editor: Alan Moorwood European Southern Observatory Karl-Schwarzschild-Str. 2 85748 Garching Germany [email protected] Series Advisory Editors: W.B. Burton National Radio Astronomy Observatory Charlottesville, VA USA

ISBN 978-1-4020-9189-6

Lars Lindberg Christensen ESA/Hubble Garching bei München Germany

e-ISBN 978-1-4020-9190-2

Library of Congress Control Number: 2008936095 © 2009 Springer Science + Business Media B.V. No part of this work may be reproduced, stored in a retrieval system, or transmitted in any form or by any means, electronic, mechanical, photocopying, microfilming, recording or otherwise, without the written permission from the Publisher, with the exception of any material supplied specifically for the purpose of being entered and executed on a computer system, for the exclusive use by the purchaser of the work. Printed on acid-free paper 9 8 7 6 5 4 3 2 1 springer.com

Preface

The Workshop “Science with the VLT in the ELT Era” held in Garching from 8th to 12th October 2007 was organised by ESO, with support from its Scientific and Technical Committee, to provide a forum for the astronomical community to debate the long term future of ESO’s Very Large Telescope (VLT) and its interferometric mode (VLTI). In particular it was considered useful for future planning to evaluate how its science use may evolve over the next decade due to competition and/or synergy with new facilities such as ALMA, JWST and, hopefully, at least one next generation 30–40 m extremely large telescope whose acronym appears in the title to symbolise this wider context. These discussions were also held in the fresh light of the Science Vision recently developed within ASTRONET as the first step towards a 20 year plan for implementing astronomical facilities—the first such attempt within Europe. Specific ideas and proposals for new, second generation VLT/I instruments were also solicited following a tradition set by several earlier Workshops held since the start of the VLT development. The programme consisted of invited talks and reviews and contributed talks and posters. Almost all those given are included here although, unfortunately not the several lively but constructive discussion sessions. The scientific context was set by presentations of the highlights of nearly 10 years of VLT/I operations followed by projections into the future, including themes of the ASTRONET Science Vision to which the VLT will clearly contribute for the foreseeable future such as exoplanet searches and characterisation, tests of general relativity, galaxy evolution, etc. The scope was then widened with presentations of the European ELT programme, ALMA (ESO on behalf of Europe, North America and Japan (with Taiwan) and the James Web Space Telescope (JWST) and discussions of the synergies and complementarities to be expected. Entering more then into the detailed plans for instrumentation at the VLT/I, the already approved second generation instruments under development were presented (X-Shooter, KMOS, SPHERE and MUSE) plus the proposed VLTI instruments GRAVITY, MATISSE and VSI and an overview given of the expected resources available for additional instruments within ESO’s medium and long range plans. Finally, the longest session was devoted to new instrumentation proposals and related technology which confirmed that there should be no problem in making good use of the available resources. There was also a clear consensus that the priority was for maintaining or providing new instruments which fully exploit VLT/I’s unique capability to combine four 8.2 m telescopes at both its coherent and incoherent foci. In fact the Workshop contributed directly to subsequent recommendations made by our STC committee and endorsed by Council to proceed with development of all three of the VLTI second generation instruments proposed; to issue a formal call for proposals for an ultra-stable spectrograph (ESPRESSO) at the incoherent combined focus and to upgrade operating first generation instruments including the wide v

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field, multi-object spectrograph VIMOS. ESO has also been encouraged to issue calls for additional instruments including a new multi-object spectrograph covering the widest possible field for studies related to the nature and origin of dark energy. However, because of the technical challenges and likely cost involved the need for a wider study of the options for this within Europe has been recommended by ASTRONET in its infrastructure road-map. I couldn’t have organised such a large and successful Workshop without the help and support of many people and wish particularly here to thank the members of the Scientific Organising Committee—Willy Benz, Mark Casali, Tom Herbst (co-chair), Bruno Leibundgut, Yannick Mellier and Jorge Melnick and of the Local Organising Committee—Markus Kissler-Patig, Christina Stoffer, Iris Bronnert and Pam Bristow. I am also grateful to all the authors who have contributed to this written record for posterity. Also, special thanks to the Max Planck Institut für Extraterrestriche Physik and to Linda Tacconi in particular for making available to us their large seminar room once it was clear that the interest in this Workshop had exceed the capacity of our Auditorium. ESO, Garching bei München, Germany July 2008

Alan F.M. Moorwood

Contents

Preface Part I

VLT Science Highlights

VLT Science Highlights Alvio Renzini . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

3

Pushing the VLT Spectroscopy of Distant Galaxies to the Limits and Future Prospects Andrea Cimatti . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

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Pushing FORS to the Limit—A New Population of Faint Extended Lyα Emitters at z ∼ 3 Martin G. Haehnelt, Michael Rauch, Andrew Bunker, George Becker, Francine Marleau, James Graham, Stefano Cristiani, Matt J. Jarvis, Cedric Lacey, Simon Morris, Celine Peroux, Huub Röttgering, and Tom Theuns . . . . . . . . . . . . . . . . .

23

VIMOS Integral Field Spectroscopy of Gaseous Nebulae in Local Group Dwarf Galaxies E.V. Held, M. Gullieuszik, I. Saviane, F. Sabbadin, Y. Momany, L. Rizzi, and F. Bresolin . . . . . . . . . . . . . . . . . . . . . . .

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Near IR Integral Field Spectroscopy of a Nearby Starburst L. Vanzi, G. Cresci, J. Melnick, and E. Telles . . . . . . . . . . . . . . .

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The ESO Large Programme “First Stars” P. Bonifacio, J. Andersen, S.M. Andrievsky, B. Barbuy, T.C. Beers, E. Caffau, R. Cayrel, E. Depagne, P. François, J.I. González Hernández, C.J. Hansen, F. Herwig, V. Hill, S.A. Korotin, H.-G. Ludwig, P. Molaro, B. Nordström, B. Plez, F. Primas, T. Sivarani, F. Spite, and M. Spite . . . . . . . . . . . .

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The Contribution of UVES@VLT to the New Era of QSO Absorption Line Studies Valentina D’Odorico and Miroslava Dessauges-Zavadsky . . . . . . . .

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IMAGES: A Unique View of the Galaxy Mass Assembly Since z = 1 M. Puech, F. Hammer, H. Flores, Y. Yang, and B. Neichel . . . . . . . .

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The Metallicity Evolution at High Redshift R. Maiolino, T. Nagao, A. Grazian, F. Cocchia, and the Amaze Team . .

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Near-IR Spectroscopy of Blue Supergiants N. Przybilla, A. Seifahrt, K. Butler, M.F. Nieva, H.-U. Käufl, and A. Kaufer . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

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Integral Field Spectroscopy of Protoplanetary Disks in Orion with VLT FLAMES Y.G. Tsamis, J.R. Walsh, and D. Péquignot . . . . . . . . . . . . . . . .

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MAD@VLT: Deep into the Madding Crowd of Omega Centauri G. Bono, A. Calamida, C.E. Corsi, P.B. Stetson, E. Marchetti, P. Amico, P.G. Prada Moroni, I. Ferraro, G. Iannicola, M. Monelli, R. Buonanno, F. Caputo, M. Dall’Ora, S. Degl’Innocenti, S. D’Odorico, L.M. Freyhammer, D. Koester, M. Nonino, A.M. Piersimoni, L. Pulone, and M. Romaniello . . . . . . . . . .

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Chemical Evolution of the Galaxy and Supernova Yields after UVES G. Israelian and P. Bonifacio . . . . . . . . . . . . . . . . . . . . . . . .

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Part II

VLTI Science Highlights

VLTI Science Highlights Guy Perrin . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

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MIDI Sees Active Galactic Nuclei W. Jaffe, D. Raban, K. Meisenheimer, K. Tristram, Ch. Leinert, and H. Röttgering . . . . . . . . . . . . . . . . . . . . . . . . . . . .

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The Use of the VLTI for Studying the Asymmetric Mass Loss of Evolved Stars Olivier Chesneau . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

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Mid-infrared Interferometric Observations of Young Circumstellar Discs Th. Ratzka, Ch. Leinert, R. van Boekel, and A.A. Schegerer . . . . . . . 101

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VLTI-AMBER Observations of η Carinae with High Spatial Resolution and Spectral Resolutions of λ/λ = 1500 and 12 000∗ G. Weigelt, S. Kraus, T. Driebe, K.-H. Hofmann, F. Millour, R. Petrov, D. Schertl, O. Chesneau, K. Davidson, A. Domiciano de Souza, T. Gull, J.D. Hillier, F. Malbet, F. Rantakyrö, A. Richichi, M. Schöller, and M. Wittkowski . . . . . . . . . . . . . . . . . . . 107 Resolving the Inner Active Accretion Disk Around the Herbig Be Star MWC 147 with VLTI/MIDI + AMBER Spectro-interferometry S. Kraus, Th. Preibisch, and K. Ohnaka . . . . . . . . . . . . . . . . . . 113 Multi-epoch VLTI/MIDI Observations of the Carbon-rich Mira Star V Oph K. Ohnaka, T. Driebe, G. Weigelt, and M. Wittkowski . . . . . . . . . . 119 A Mid-infrared Interferometric Study of the Circumstellar Environment of Dusty OH/IR Stars with VLTI/MIDI T. Driebe, K. Ohnaka, K. Murakawa, K.-H. Hofmann, D. Schertl, G. Weigelt, T. Verhoelst, O. Chesneau, A. Domiciano de Souza, D. Riechers, M. Schöller, and M. Wittkowski . . . . . . . . . . . . 125 The Closest Dusty Cloud Ever Detected Around a R CrB Variable Star Using the VLTI/MIDI Instrument I.C. Leão, P. de Laverny, O. Chesneau, D. Mékarnia, and J.R. De Medeiros . . . . . . . . . . . . . . . . . . . . . . . . . . . 127 Part III Future VLT and VLTI Science Priorities A Twenty Year Science Vision for European Astronomy Guy Monnet and Tim de Zeeuw . . . . . . . . . . . . . . . . . . . . . . 133 Baryonic Acoustic Oscillations Gavin Dalton . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 141 Galaxy Formation and Evolution J. Bergeron . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 147 Exoplanets: The Road to Earth Twins S. Udry, F. Pepe, C. Lovis, M. Mayor, the HARPS, and ESPRESSO/CODEX Teams . . . . . . . . . . . . . . . . . . . . 155 Next Generation Deep Redshift Surveys with the VLT Olivier Le Fèvre . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 163

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GUAIX: The UCM Group of Extragalactic Astrophysics and Astronomical Instrumentation J. Gallego, N. Cardiel, S. Pascual, M.C. Eliche-Moral, A. Castillo-Morales, R. Guzmán, A. Gil de Paz, P.G. Pérez-González, J. Gorgas, J. Zamorano, and GUAIX Team . . . . . . . . . . . . . . . . . . . . . . . . . . . . 169 VISTA Public Surveys and VLT followup Will Sutherland . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 171 Probing Dark Energy with Cosmological Redshift Surveys at the VLT L. Guzzo and the VVDS Consortium . . . . . . . . . . . . . . . . . . . 177 The First Galaxies and Galaxy Clusters Eelco van Kampen . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 183 Narrow Band Surveys and the Epoch of Reionization B.P. Venemans, R.G. McMahon, I.R. Parry, D.J. King, J. Bland-Hawthorn, and A.J. Horton . . . . . . . . . . . . . . . . . 187 Stellar Archaeology and Galaxy Genesis: The Need for Large Area Multi-Object Spectrograph on 8 m-Class Telescopes Mike J. Irwin and Geraint F. Lewis . . . . . . . . . . . . . . . . . . . . 193 Near-field Cosmology with the VLT Steffen Mieske and Helmut Jerjen . . . . . . . . . . . . . . . . . . . . . 199 Chemical Evolution of Local Group Galaxies Gražina Tautvaišien˙e, Doug Geisler, and George Wallerstein . . . . . . . 205 The VLTI as a Tool to Study Eclipsing Binaries for an Improved Distance Scale K. Shabun, A. Richichi, U. Munari, A. Siviero, and B. Paczynski . . . . 211 Part IV VLT and VLTI Synergy with ELTs Status of the European ELT Roberto Gilmozzi and Jason Spyromilio . . . . . . . . . . . . . . . . . 217 The Science Case for the European ELT Isobel Hook . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 225 GRB Afterglows in the ELT Era David Alexander Kann and Sylvio Klose . . . . . . . . . . . . . . . . . 233 On the Way to an E-ELT Instrumentation Plan Sandro D’Odorico, Mark Casali, and Vincenzo Mainieri . . . . . . . . . 235

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From ESPRESSO to CODEX J. Liske, L. Pasquini, P. Bonifacio, F. Bouchy, R.F. Carswell, S. Cristiani, M. Dessauges, S. D’Odorico, V. D’Odorico, A. Grazian, R. Garcia-Lopez, M. Haehnelt, G. Israelian, C. Lovis, E. Martin, M. Mayor, P. Molaro, M.T. Murphy, F. Pepe, D. Queloz, R. Rebolo, S. Udry, E. Vanzella, M. Viel, T. Wiklind, M. Zapatero, and S. Zucker . . . . . . . . . . . . . . . . . . . . . . . . . . . . 243 First Results of AQuEye, a Precursor ‘Quantum’ Instrument for the E-ELT C. Barbieri, G. Naletto, E. Verroi, C. Facchinetti, T. Occhipinti, A. Di Paola, E. Giro, P. Zoccarato, G. Anzolin, M. D’Onofrio, F. Tamburini, G. Bonanno, S. Billotta, C. Pernechele, P. Bolli, V. Da Deppo, and S. Fornasier . . . . . . . . . . . . . . . . . . . . 249 The E-ELT: A Chance to Measure Cosmic Magnetic Fields K.G. Strassmeier and I.V. Ilyin . . . . . . . . . . . . . . . . . . . . . . 255 The Experience from VISIR and the Design of an ELT Mid-infrared Instrument E. Pantin, R. Siebenmorgen, H.U. Käufl, and M. Sterzik . . . . . . . . . 261 HARMONI: A Narrow Field Near-infrared Integral Field Spectrograph for the E-ELT Matthias Tecza, Niranjan Thatte, Fraser Clarke, and David Freeman . . . 267 Which Synergies Between LBT/LINC Nirvana and Future ELTs? L. Labadie, T.M. Herbst, S. Egner, M. Brix, and M. Kürtser . . . . . . . 273 TMT Science and Instruments David Crampton, Luc Simard, and David Silva . . . . . . . . . . . . . . 279 Part V

VLT Synergies with ALMA and JWST

The Atacama Large Millimeter/Submillimeter Array Leonardo Testi . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 289 Observational Cosmology with the ELT and JWST Massimo Stiavelli . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 295 Integral Field Spectroscopy of (U)LIRGs. From VLT to JWST L. Colina, S. Arribas, A. Bedregal, A. Monreal-Ibero, M. García-Marín, A. Alonso-Herrero, and J. Alfonso . . . . . . . . . . . . . . . . . 301

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Part VI Second Generation VLT and VLTI Instrument Programme VLT and VLTI Second Generation Instrument Overview and Resources Alan Moorwood . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 309 HAWK-I and Infrared Imaging on the VLT M. Casali, N. Ageorge, C. Alves de Oliveira, P. Biereichel, M. Casali, B. Delabre, R. Dorn, R. Esteves, G. Finger, D. Gojak, G. Huster, Y. Jung, F. Koch, M. Kiekebusch, M. Kissler-Patig, M. Le Louarn, J.-L. Lizon, L. Mehrgan, A. Moorwood, J. Pirard, E. Pozna, A. Silber, B. Sokar, and J. Stegmeier . . . . . . . . . . . . . . . . 315 X-Shooter: A Medium-resolution, Wide-Band Spectrograph for the VLT L. Kaper, S. D’Odorico, F. Hammer, R. Pallavicini, P. Kjaergaard Rasmussen, H. Dekker, P. Francois, P. Goldoni, I. Guinouard, P.J. Groot, J. Hjorth, M. Horrobin, R. Navarro, F. Royer, P. Santin, J. Vernet, and F. Zerbi . . . . . . . . . . . . . . . . . . . . . . . . 319 KMOS and KMOS++ Ray Sharples and KMOS Consortium . . . . . . . . . . . . . . . . . . . 325 New Science Opportunities Offered by MUSE R. Bacon, S. Bauer, S. Brau-Nogué, P. Caillier, L. Capoani, M. Carollo, T. Contini, E. Daguisé, B. Delabre, S. Dreizler, J.P. Dubois, M. Dupieux, J. Dupin, E. Emsellem, P. Ferruit, M. Francois, M. Franx, G. Gallou, J. Gerssen, B. Guiderdoni, G. Hansali, D. Hofmann, A. Jarno, A. Kelz, C. Koehler, W. Kollatschny, J. Kosmalski, F. Laurent, S. Lilly, J. Lizon, M. Loupias, C. Monstein, J. Moultaka, H. Nicklas, L. Parés, L. Pasquini, A. Pecontal, R. Pello, C. Petit, A. Manescau, R. Reiss, A. Remillieux, E. Renault, M. Roth, J. Schaye, M. Steinmetz, S. Ströbele, R. Stuik, P. Weilbacher, L. Wisotzki, and H. Wozniak . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 331 SPHERE: A ‘Planet Finder’ Instrument for the VLT D. Mouillet, J.-L. Beuzit, M. Feldt, K. Dohlen, P. Puget, F. Wildi, A. Boccaletti, T. Henning, C. Moutou, H.M. Schmid, M. Turatto, S. Udry, F. Vakili, R. Waters, A. Baruffolo, J. Charton, R. Claudi, T. Fusco, R. Gratton, N. Hubin, M. Kasper, M. Langlois, J. Pragt, R. Roelfsema, and M. Saisse . . . . . . . . . . . . . . . . . . . . 337

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Milli-arcsecond Astrophysics with VSI, the VLTI Spectro-imager in the ELT Era F. Malbet, D. Buscher, G. Weigelt, P. Garcia, M. Gai, D. Lorenzetti, J. Surdej, J. Hron, R. Neuhäuser, P. Kern, L. Jocou, J.-P. Berger, O. Absil, U. Beckmann, L. Corcione, G. Duvert, M. Filho, P. Labeye, E. Le Coarer, G. Li Causi, J. Lima, K. Perraut, E. Tatulli, E. Thiébaut, J. Young, G. Zins, A. Amorim, B. Aringer, T. Beckert, M. Benisty, X. Bonfils, A. Chelli, O. Chesneau, A. Chiavassa, R. Corradi, M. de Becker, A. Delboulbé, G. Duchêne, T. Forveille, C. Haniff, E. Herwats, K.-H. Hofmann, J.-B. Le Bouquin, S. Ligori, D. Loreggia, A. Marconi, A. Moitinho, B. Nisini, P.-O. Petrucci, J. Rebordao, R. Speziali, L. Testi, and F. Vitali . . . . . . . . . . . . . . . . . . . . . . . . . 343 Prospects for Near-infrared Characterisation of Hot Jupiters with the VLTI Spectro-imager (VSI) S. Renard, O. Absil, J.-P. Berger, X. Bonfils, T. Forveille, and F. Malbet . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 349 MATISSE B. Lopez, S. Lagarde, S. Wolf, W. Jaffe, G. Weigelt, P. Antonelli, P. Abraham, J.-Ch. Augereau, U. Beckman, J. Behrend, N. Berruyer, Y. Bresson, O. Chesneau, J.M. Clausse, C. Connot, W.C. Danchi, M. Delbo, K. Demyk, A. Domiciano, M. Dugué, A. Glazenborg, U. Graser, H. Hanenburg, Th. Henning, M. Heininger, K.-H. Hofmann, Y. Hugues, S. Jankov, S. Kraus, W. Laun, Ch. Leinert, H. Linz, A. Matter, Ph. Mathias, K. Meisenheimer, J.-L. Menut, F. Millour, L. Mosoni, U. Neumann, A. Niedzielski, E. Nussbaum, R. Petrov, Th. Ratzka, S. Robbe-Dubois, A. Roussel, D. Schertl, F.-X. Schmider, B. Stecklum, E. Thiebaut, F. Vakili, K. Wagner, L.B.F.M. Waters, O. Absil, J. Hron, A. Matter, N. Nardetto, J. Olofsson, B. Valat, M. Vannier, B. Goldman, D. Schertl, S. Hönig, and W.D. Cotton . . 353 MATISSE Science Cases S. Wolf, B. Lopez, W. Jaffe, G. Weigelt, J.-Ch. Augereau, N. Berruyer, O. Chesneau, W.C. Danchi, M. Delbo, K. Demyk, A. Domiciano, Th. Henning, K.-H. Hofmann, S. Kraus, Ch. Leinert, H. Linz, Ph. Mathias, K. Meisenheimer, J.-L. Menut, F. Millour, L. Mosoni, A. Niedzielski, R. Petrov, Th. Ratzka, B. Stecklum, E. Thiebaut, F. Vakili, L.B.F.M. Waters, O. Absil, J. Hron, S. Lagarde, A. Matter, N. Nardetto, J. Olofsson, B. Valat, M. Vannier, and MATISSE Science team . . . . . . . . . . . . . . . . . . . . . . . 359

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GRAVITY: Microarcsecond Astrometry and Deep Interferometric Imaging with the VLT F. Eisenhauer, G. Perrin, W. Brandner, C. Straubmeier, A. Böhm, H. Baumeister, F. Cassaing, Y. Clénet, K. Dodds-Eden, A. Eckart, E. Gendron, R. Genzel, S. Gillessen, A. Gräter, C. Gueriau, N. Hamaus, X. Haubois, M. Haug, T. Henning, S. Hippler, R. Hofmann, F. Hormuth, K. Houairi, S. Kellner, P. Kervella, R. Klein, J. Kolmeder, W. Laun, P. Léna, R. Lenzen, M. Marteaud, V. Naranjo, U. Neumann, T. Paumard, S. Rabien, J.R. Ramos, J.M. Reess, R.-R. Rohloff, D. Rouan, G. Rousset, B. Ruyet, A. Sevin, M. Thiel, J. Ziegleder, and D. Ziegler . . . . . . . . . . . 361 Part VII New Instrument Concepts and VLT/I Operating Modes Smart Focal Plane Technologies for VLT Instruments C.R. Cunningham and C.J. Evans . . . . . . . . . . . . . . . . . . . . . 369 Applications of Digital Micromirror Devices to Astronomical Instrumentation M. Robberto . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 375 FORS in the Era of Second Generation VLT Instrumentation Kieran O’Brien . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 377 Wide Field Options on the VLT Stephen Todd . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 379 A Few Degrees Very Wide Field of View Camera for VLT as a Finder for ELT Roberto Ragazzoni, Jacopo Farinato, Emiliano Diolaiti, Giorgia Gentile, Carmelo Arcidiacono, Renato Falomo, and Emanuele Giallongo . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 385 Science with a 16 m VLT: The Case for Variability of Fundamental Constants Paolo Molaro . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 389 ESPRESSO: A High Resolution Spectrograph for the Combined Coudé Focus of the VLT Luca Pasquini, A. Manescau, G. Avila, B. Delabre, H. Dekker, J. Liske, S. D’Odorico, F. Pepe, M. Dessauges, C. Lovis, D. Megevand, D. Queloz, S. Udry, S. Cristiani, P. Bonifacio, P. Dimarcantonio, V. D’Odorico, P. Molaro, E. Vanzella, M. Viel, M. Haehnelt, B. Carswell, M. Murphy, R. Garcia-Lopez, J.M. Herreros, J. Perez, M.R. Zapatero, R. Rebolo, G. Israelian, E. Martin, F. Zerbi, P. Spanò, S. Levshakov, N. Santos, and S. Zucker . . . . . 395

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Feeding Optics for the ESPRESSO Spectrograph G. Avila, P. Dimarcantonio, and F. Zerbi . . . . . . . . . . . . . . . . . 401 New Design Approach for a Very-High Resolution Spectrograph for the VLT Combined Focus Paolo Spanò and Hans Dekker . . . . . . . . . . . . . . . . . . . . . . . 403 ESPRESSO Optomechanics J. Pérez, H. Dekker, R.J. García López, J.M. Herreros, R. López, F. Pepe, J.L. Rasilla, P. Spanò, and M.R. Zapatero Osorio . . . . . 405 ESPRESSO Science Software D. Mégevand, V. D’Odorico, and C. Lovis . . . . . . . . . . . . . . . . 409 High Resolution Wavelength Calibration: Advancements with the Laser Frequency Comb Development A. Manescau, C. Araujo-Hauck, L. Pasquini, M.T. Murphy, Th. Udem, T.W. Hänsch, R. Holzwarth, A. Sizmann, H. Dekker, and S. D’Odorico . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 411 Precision Radial Velocities in the Infrared Hugh R.A. Jones, John Rayer, Larry Ramsey, Bill Dent, Andy Longmore, Bill Vacca, Mike Liu, Adrian Webster, Alex Wolscznan, and John Barnes . . . . . . . . . . . . . . . . . . 415 Very Large Spectroscopic Surveys with the VLT I.R. Parry . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 417 New Developments in Integral Field Spectroscopy Anthony Horton, Joss Bland-Hawthorn, and Simon Ellis . . . . . . . . . 423 ULTRAPHOT Françoise Roques, Isabelle Guinouard, Jean-Tristan Buey, Alain Doressoundiram, David Horville, and Michel Marteaud . . . 429 Super-GIRAFFE: The Next Generation High Multiplex Optical Spectrograph with d-IFUs M.D. Lehnert, I. Guinouard, D. Horville, P. Jagourel, F. Chemla, J.-P. Amans, P. Bonifacio, C. Babusiaux, F. Hammer, V. Hill, F. Royer, and M. Puech . . . . . . . . . . . . . . . . . . . . . . . 431 FLEX (The First Light Explorer)—The Science Case for a Fully OH Suppressed IFU Spectrograph Simon Ellis, Joss Bland-Hawthorn, Anthony Horton, and Roger Haynes . . . . . . . . . . . . . . . . . . . . . . . . . . . . 437 An N -Band Integral Field Spectrometer Survey Instrument for the VLT A.C.H. Glasse, D.M. Henry, and D. Lee . . . . . . . . . . . . . . . . . . 443

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Contents

High Resolution Visible Imaging on the VLT Craig Mackay . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 449 Life on the Fast Lane: The Burst Mode at the VLT at Present and in the Future Andrea Richichi, Octavi Fors, Elena Mason, Marco Delbó, Jörg Stegmaier, and Gert Finger . . . . . . . . . . . . . . . . . . . 455 High Resolution Near Infrared Spectroscopy: Prospects for 10 and 40 m Class Telescopes E. Oliva and L. Origlia . . . . . . . . . . . . . . . . . . . . . . . . . . . 461 Prospects and Needs of Micro-arcsecond Astrometry Andreas Seifahrt, Tristan Röll, and Ralph Neuhäuser . . . . . . . . . . . 469 CASIS: Cassegrain Adaptive-Optics Simultaneous Imaging System for the VLT M. Kissler-Patig, M. Casali, B. Delabre, N. Hubin, H.U. Käufl, P. Jolley, M. Le Louarn, S. Oberti, and J. Pirard . . . . . . . . . . . . . . . . 475 The Need for a General Purpose Diffraction Limited Imager at the VLT Thomas Ott, Richard Davies, Frank Eisenhauer, Reinhard Genzel, Reiner Hofmann, and Stefan Gillessen . . . . . . . . . . . . . . . 481 Exploring the Time Axis—High Resolution Timing Observations with Present and Future Instrumentation V.D. Ivanov, C. Caceres, E. Mason, D. Naef, F. Selman, C. Melo, D. Minniti, and G. Pietrzynski . . . . . . . . . . . . . . . . . . . . 487 Advanced Calibration for Quantitative Astrophysics: 2nd Generation VLT Instruments and Beyond Florian Kerber, Paul Bristow, and Michael R. Rosa . . . . . . . . . . . . 493 Quantitative Near-IR Spectroscopy of OB Stars M.F. Nieva, N. Przybilla, A. Seifahrt, K. Butler, H.U. Käufl, and A. Kaufer . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 499 The Very Large Telescope Interferometer in the ELT Era M. Schöller, F. Delplancke, A. Glindemann, and A. Richichi . . . . . . . 501 VLTI and Beyond: The Next Steps in AGN Research with Interferometers Klaus Meisenheimer . . . . . . . . . . . . . . . . . . . . . . . . . . . . 507

Participants List

ARCIDIACONO, Carmelo INAF – Osservatorio Astronomico di Padova (Italy) [email protected] ARDEBERG, Arne Lund Observatory (Sweden) [email protected] AVILA, Gerardo ESO – Garching [email protected]

BRYSON, Ian UK Astronomy Technology Centre, Edinburgh (UK) [email protected] CASALI, Mark ESO – Garching [email protected] CHELLI, Alain LAOG (Laboratoire d’Astrophysique de Grenoble) (France) [email protected]

BAADE, Dietrich ESO – Garching [email protected]

CHESNEAU, Olivier Observatoire de la Côte d’Azur (OCA), Grasse (France) [email protected]

BACON, Roland CRAL – Observatoire de Lyon (France) [email protected]

CIMATTI, Andrea Università di Bologna, Dip. di Astronomia (Italy) [email protected]

BARBIERI, Cesare Università di Padova, Dip. di Astronomia (Italy) [email protected] BERGERON, Jacqueline Institut d’Astrophysique de Paris (France) [email protected] BONIFACIO, Piercarlo GEPI – Observatoire de Paris, Meudon (France) [email protected]

CIRASUOLO, Michele Institute for Astronomy, Edinburgh University (UK) [email protected] CLARKE, Fraser University of Oxford (UK) [email protected] COLINA, Luis DAMIR/IEM/CSIC, Madrid (Spain) [email protected]

BONO, Giuseppe INAF – Osservatorio Astronomico di Roma (Italy) [email protected]

CRAMPTON, David Dominion Astrophysical Observatory (Canada) [email protected]

BOUTSIA, Konstantina ESO – Garching [email protected]

CUBY, Jean-Gabriel Laboratoire d’Astrophysique de Marseille (France) [email protected] xvii

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Participants List

CUNNINGHAM, Colin UK Astronomy Technology Centre, Edinburgh (UK) [email protected]

EVANS, Chris UK Astronomy Technology Centre, Edinburgh (UK) [email protected]

D’ODORICO, Sandro ESO – Garching [email protected]

FALOMO, Renato INAF – Osservatorio Astronomico di Padova (Italy) [email protected]

D’ODORICO, Valentina INAF – Osservatorio Astronomico di Trieste (Italy) [email protected] DALTON, Gavin University of Oxford/CCLRC Rutherford Appleton Laboratory (UK) [email protected] DE ZEEUW, Tim ESO – Garching [email protected] DEKKER, Hans ESO – Garching [email protected] DENNEFELD, Michel Institut d’Astrophysique de Paris (France) [email protected] DRIEBE, Thomas Max-Planck-Institut für Radioastronomie, Bonn (Germany) [email protected] EISENHAUER, Frank Max-Planck-Institut für Extraterrestrische Physik, Garching (Germany) [email protected] ELLIS, Simon Anglo-Australian Observatory, Eastwood (Australia) [email protected]

FARINATO, Jacopo INAF – Osservatorio Astronomico di Padova (Italy) [email protected] FORSTER SCHREIBER, Natascha Max-Planck-Institut für Extraterrestrische Physik, Garching (Germany) [email protected] FRANX, Marijn Leiden Observatory (Netherlands) [email protected] GALLEGO, Jesus Universidad Complutense, Madrid (Spain) [email protected] GENTILE, Giorgia INAF – Osservatorio Astronomico di Padova (Italy) [email protected] GENZEL, Reinhard Max-Planck-Institut für Extraterrestrische Physik, Garching (Germany) [email protected] GIALLONGO, Emanuele INAF – Osservatorio Astronomico di Roma (Italy) [email protected]

ENARD, Daniel ELT [email protected]

GILMORE, Gerry Institute of Astronomy, University of Cambridge (UK) [email protected]

ERGENZINGER, Klaus Astrium GmbH, Friedrichshafen (Germany) [email protected]

GILMOZZI, Roberto ESO – Garching [email protected]

Participants List GLASSE, Alistair UK Astronomy Technology Centre, Edinburgh (UK) [email protected] GLINDEMANN, Andreas ESO – Garching [email protected] GRATTON, Raffaele INAF – Osservatorio Astronomico di Padova (Italy) [email protected]

xix HOOK, Isobel University of Oxford (UK) [email protected] HORROBIN, Matthew Astronomical Institute, University of Amsterdam (Netherlands) [email protected] HORTON, Anthony Anglo-Australian Observatory, Eastwood (Australia) [email protected]

GROENEWEGEN, Martin Institute for Astronomy, K.U. Leuven (Belgium) [email protected]

HUBIN, Norbert ESO – Garching [email protected]

GUZZO, Luigi INAF – Osservatorio Astronomico di Brera (Italy) [email protected]

HURTADO, Norma Service d’Astrophysique CEA-Saclay (France) [email protected]

HÄHNELT, Martin Institute of Astronomy, University of Cambridge (UK) [email protected]

IRWIN, Mike Institute of Astronomy, University of Cambridge (UK) [email protected]

HALL, Donald Institute for Astronomy, University of Hawaii (USA) [email protected]

IVANOV, Valentin ESO – Chile [email protected]

HAMMER, François GEPI – Observatoire de Paris, Meudon (France) [email protected]

IVISON, Rob Royal Observatory Edinburgh (UK) [email protected]

HANSEN, Camilla ESO – Garching [email protected]

IZAN DE CASTRO, Leao Universidade Federal do RN (Brazil) [email protected]

HANUSCHIK, Reinhard ESO – Garching [email protected]

JAFFE, Walter Leiden Observatory (Netherlands) [email protected]

HELD, Enrico V. INAF – Osservatorio Astronomico di Padova (Italy) [email protected]

JONES, Hugh University of Hertfordshire (UK) [email protected]

HERBST, Thomas Michael Max-Planck-Institut für Astronomie, Heidelberg (Germany) [email protected]

KAPER, Lex Astronomical Institute, University of Amsterdam (Netherlands) [email protected]

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Participants List LIGORI, Sebastiano INAF – Osservatorio Astronomico di Torino (Italy) [email protected]

KÄUFL, Hans Ulrich ESO – Garching [email protected] KERBER, Florian ESO – Garching [email protected] KISHIMOTO, Makoto Max-Planck-Institut für Bonn (Germany) [email protected]

LIND, Karin ESO – Garching [email protected] Radioastronomie,

KISSLER-PATIG, Markus ESO – Garching [email protected] KJAER, Karina ESO – Garching [email protected] KLOSE, Sylvio Thüringer Landessternwarte Tautenburg (Germany) [email protected] KORHONEN, Heidi ESO – Garching [email protected] KRAUS, Stefan Max-Planck-Institut für Radioastronomie, Bonn (Germany) [email protected] LABADIE, Lucas Max-Planck-Institut für Astronomie, Heidelberg (Germany) [email protected] LE FÈVRE, Olivier Laboratoire d’Astrophysique de Marseille (France) [email protected] LEHNERT, Matt GEPI – Observatoire de Paris, Meudon (France) [email protected] LEIBUNDGUT, Bruno ESO – Garching [email protected]

LISKE, Jochen ESO – Garching [email protected] LONGMORE, Andrew UK Astronomy Technology Centre, Edinburgh (UK) [email protected] LOPEZ, Bruno Observatoire de la Côte d’Azur (OCA), Nice (France) [email protected] LORENZETTI, Dario INAF – Osservatorio Astronomico di Roma (Italy) [email protected] LOVIS, Christophe Geneva Observatory, Versoix (Switzerland) [email protected] MACKAY, Craig Institute of Astronomy, University of Cambridge (UK) [email protected] MAINIERI, Vincenzo ESO – Garching [email protected] MAIOLINO, Roberto INAF – Osservatorio Astronomico di Roma (Italy) [email protected] MALBET, Fabien LAOG/CNRS/UJF (France) [email protected] MANDEL, Holger Landessternwarte Heidelberg (ZAH), Heidelberg (Germany) [email protected]

Participants List MANESCAU, Antonio ESO – Garching [email protected] MARCONI, Gianni ESO – Chile [email protected] MARSH, Thomas Dep. of Physics, Univ of Warwick (UK) [email protected] MCCAUGHREAN, Mark University of Exeter (UK) [email protected] MCLURE, Ross Institute for Astronomy, Edinburgh University (UK) [email protected] MÉGEVAND, Denis Geneva Observatory, Sauverny (Switzerland) [email protected] MEISENHEIMER, Klaus Max-Planck-Institut für Astronomie, Heidelberg (Germany) [email protected] MELNICK, Jorge ESO – Garching/Chile [email protected] MIESKE, Steffen ESO – Garching [email protected] MOLARO, Paolo INAF – Osservatorio Astronomico di Trieste (Italy) [email protected]

xxi MOUILLET, David LAOG (Laboratoire d’Astrophysique de Grenoble) (France) [email protected] NEUHÄUSER, Ralph Astrophysikalisches Institut und UniversitätsSternwarte, Jena (Germany) [email protected] NICHOL, Bob Institute of Cosmology and Gravitation, Portsmouth (UK) [email protected] NIEVA, Maria Fernanda Dr. Remeis Sternwarte Bamberg (Germany) [email protected] O’BRIEN, Kieran ESO – Chile [email protected] OLIVA, Ernesto INAF TNG, Santa Cruz de La Palma, TF (Spain) [email protected] ORIGLIA, Livia INAF – Osservatorio Astronomico di Bologna (Italy) [email protected] ORTOLANI, Sergio Università di Padova, Dip. di Astronomia (Italy) [email protected] OTT, Thomas Max-Planck-Institut für Extraterrestrische Physik, Garching (Germany) [email protected] PARRY, Ian Institute of Astronomy, University of Cambridge (UK) [email protected]

MONNET, Guy ESO – Garching [email protected]

PASQUINI, Luca ESO – Garching [email protected]

MOORWOOD, Alan ESO – Garching [email protected]

PERCIVAL, Will University of Portsmouth (UK) [email protected]

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Participants List

PERRIN, Guy LESIA – Observatoire de Paris, Meudon (France) [email protected]

RENZINI, Alvio INAF – Osservatorio Astronomico di Padova (Italy) [email protected]

PRADA, Francisco Instituto de Astrofisica de Andalucia (CSIC), Granada (Spain) [email protected]

RICHICHI, Andrea ESO – Garching [email protected]

PRIMAS, Francesca ESO – Garching [email protected]

ROBBERTO, Massimo Space Telescope Science Institute, Baltimore (USA) [email protected]

PRZYBILLA, Norbert Dr. Remeis Sternwarte Bamberg (Germany) [email protected]

ROCHE, Patrick University of Oxford (UK) [email protected]

PUECH, Mathieu ESO – Garching [email protected]

ROMANIELLO, Martino ESO – Garching [email protected]

QUELOZ, Didier Geneva Observatory, Sauverny (Switzerland) [email protected] RAGAZZONI, Roberto INAF – Osservatorio Astrofisico di Arcetri (Italy) [email protected] RASMUSSEN, Per Kjaergaard Niels Bohr Institute, Copenhagen (Denmark) [email protected] RATZKA, Thorsten Astrophysikalisches Institut, Potsdam (AIP) (Germany) [email protected] REINSCH, Klaus Institut für Astrophysik, Göttingen (Germany) [email protected] REJKUBA, Marina ESO – Garching [email protected] RENARD, Stephanie LAOG (Laboratoire d’Astrophysique de Grenoble) (France) [email protected]

ROQUES, Françoise LESIA – Observatoire de Paris, Meudon (France) [email protected] RUPPRECHT, Gero ESO – Garching [email protected] SAITTA, Francesco ESO – Garching [email protected] SANTOS, Nuno Centro de Astrofisica da Universidade do Porto (Portugal) [email protected] SCHÖLLER, Markus ESO – Chile [email protected] SEIFAHRT, Andreas Institut für Astrophysik, Göttingen (Germany) [email protected] SHABUN, Klara ESO – Garching [email protected]

Participants List SHARPLES, Ray University of Durham (UK) [email protected] SIEBENMORGEN, Ralf ESO – Garching [email protected] SPANÒ, Paolo INAF – Osservatorio Astronomico di Brera (Italy) [email protected] SPYROMILIO, Jason ESO – Garching [email protected] STERZIK, Michael ESO – Chile [email protected] STIAVELLI, Massimo Space Telescope Science Institute, Baltimore (USA) [email protected] STRASSMEIER, Klaus G. Astrophysikalisches Institut, Potsdam (AIP) (Germany) [email protected] SUTHERLAND, Will Queen Mary College, University of London (UK) [email protected] TACCONI, Linda Max-Planck-Institut für Extraterrestrische Physik, Garching (Germany) [email protected] TACCONI-GARMAN, Lowell ESO – Garching [email protected] TAMAI, Roberto ESO – Chile [email protected] TANAKA, Masayuki ESO – Garching [email protected]

xxiii TAUTVAISIENE, Grazina Institute of Theoretical Physics and Astronomy of Vilnius University (Lithuania) [email protected] TECZA, Matthias University of Oxford (UK) [email protected] TESTI, Leonardo ESO – Garching [email protected] TODD, Stephen UK Astronomy Technology Centre, Edinburgh (UK) [email protected] TOLSTOY, Eline Kapteyn Astronomical Institute, Groningen (Netherlands) [email protected] TSAMIS, Yiannis University College London (UK) [email protected] UDRY, Stephane Geneva Observatory, Versoix (Switzerland) [email protected] VAN KAMPEN, Eelco Universität Innsbruck (Austria) [email protected] VANKO, Martin Astrophysikalisches Institut und UniversitätsSternwarte, Jena (Germany) [email protected] VANZI, Leonardo ESO – Chile [email protected] VENEMA, Lars ASTRON, Dwingeloo (Netherlands) [email protected] VENEMANS, Bram Institute of Astronomy, University of Cambridge (UK) [email protected]

xxiv VERMA, Aprajita University of Oxford (UK) [email protected] VETTOLANI, Giampaolo INAF HQ, Rome (Italy) [email protected] VICK, Andrew UK Astronomy Technology Centre (STFC), Edinburgh (UK) [email protected] WALSH, Jeremy ESO – Garching [email protected]

Participants List WEIGELT, Gerd Max-Planck-Institut für Radioastronomie, Bonn (Germany) [email protected] WILSON, Tom ESO – Garching [email protected] ZINNECKER, Hans Astrophysikalisches Institut, Potsdam (AIP) (Germany) [email protected]

Participants List

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Part I

VLT Science Highlights

VLT Science Highlights Alvio Renzini

1 Introduction Over eight years have passed since April first, 1999, when the UT1 of the VLT started regular scientific observations. These have been incredibly exciting years, with all the four VLT telescopes coming progressively into play, and its ten scientific instruments starting to deliver high quality data. I have been asked to start this timely meeting on how best use the VLT over the next ten years and beyond, by mentioning some of the most exciting results so far achieved with the VLT. So many and in so many disparate areas have been the outstanding VLT results, that it is a great embarrassment to be forced to make a choice. I have then decided to show one result for each of the ten VLT instruments, preferentially when the instrument was pushed to its limits. This exercise may show what fascinating science is just beyond such limits, and draw indications for possible upgrades or new VLT instruments. Inevitably, important instrument modes and science areas have been left out of this brief introductory review. The choice is certainly biased depending on my more or less direct experience, and focuses in particular on two areas, exoplanets and galaxy evolution at 1.4 < z < 3, where progress has been most spectacular in recent years.

2 NACO and the Mass of the BH at the Galactic Center I wish to start this review with what—to my taste—is the single most beautiful result so far achieved with the VLT: the NACO astrometric orbits of stars around the black hole (BH) at the center of our Galaxy (see Fig. 1). So, NACO, classical astrometry and celestial mechanics have allowed Schödel et al. [15] to measure the BH mass as (3.7 ± 1.5) × 106 M . By adding radial velocities measured with SINFONI, Eisenhauer et al. [4] were then able to cut down the error further, getting (3.61 ± 0.32) × 106 M . It was then proven, beyond any reasonable doubt, that there is indeed a supermassive BH at the center of our Galaxy, the nearest to us, and therefore an object worth of further study. One says that BHs are fully characterized by just three numbers: mass, angular momentum and electric charge. The next tantalizing step would then be to measure the BH angular momentum, another fundamental step, and at this conference we are going to hear more on how we could get there. A. Renzini () Osservatorio Astronomico di Padova, INAF, Padova, Italy e-mail: [email protected] A. Moorwood (ed.), Science with the VLT in the ELT Era, Astrophysics and Space Science Proceedings, doi: 10.1007/978-1-4020-9190-2_1, © Springer Science + Business Media B.V. 2009

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Fig. 1 The latest NACO mapping of the stellar orbits around the supermassive BH at the galactic center (Gillessen et al., in preparation). A colour version of this figure is available at dx.doi.org/10.1007/978-1-4020-9190-2_1

3 FLAMES + GIRAFFE Demonstrates Super-Helium-Rich Population With its ∼ 3 × 106 M , ω Centauri is the most massive globular cluster in the Galaxy, and since the early ’seventies it was known as the single globular having a clear spread in metallicity. Thus, it was not such a big surprise when accurate photometry with HST/WFC2 showed that its main sequence splits in two parallel sequences (see Fig. 2, left panel). The big surprise came when FLAMES + GIRAFFE spectroscopy demonstrated that stars on the bluer main sequence are a factor ∼ 2 more metal rich than those on the redder main sequence, contrary to the expectation from the well established theory of stellar structure [11]. The only way to explain the oddity has been to appeal to helium: stars on the blue main sequence must have a helium abundance by mass Y ∼ 0.37, well above the cosmic value pertaining to the red main sequence (∼ 0.24). How ω Cen managed to generate this helium enriched population remains a puzzle. A puzzle that this year has grown bigger, as helium-enriched and other sub-populations have been reported to exist in others among the most massive globulars [8, 12]. The FLAMES facility is the ideal tool to fully characterize the multiple stellar populations in globular clusters, and therefore it promises to provide unique insight and help solving the puzzle.

VLT Science Highlights

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Fig. 2 The double main sequence of ω Cen as seen by WFPC2 (left panel, [1]), and (right panel) its ACS version with highlighted the stars for which FLAMES + GIRAFFE has measured the metallicity [11]. Also plotted are theoretical main sequences with various helium contents. A colour version of this figure is available at dx.doi.org/10.1007/978-1-4020-9190-2_1

4 UVES: The Metallicity of Planet Hosts Finding extrasolar planets using medium-size telescopes has become a prosperous industry. But then, if one wants to know more about their hosts, the VLT is a unique resource. I take the study of Santos et al. [14] on the metallicity of planet hosts as representative of one of the cleanest achievements so far obtained with UVES, the high-resolution spectrograph on the VLT. Figure 3 shows the metallicity distribution of 98 planet-hosting stars (from the CORALIE radial velocity program), compared to the metallicity distribution of an unbiased sample of 875 CORALIE stars. Clearly, when it comes to make planets (even the gaseous giants detected by CORALIE) a metal-rich environment helps a lot, which suggests that a rocky seed is instrumental for starting the accretion of hydrogen and helium from the circumstellar disk.

5 FLAMES + UVES: Exoplanets in the Bulge Another way of discovering exoplanets is by very accurate photometry, searching for transits in front of host stars. One week of HST/ACS was fully dedicated to a search of transiting exoplanets in a field in the Galactic bulge [13]. Figure 4, left panel, shows the ACS color-magnitude diagram of this field, having highlighted the 16 stars hosting candidate transiting planets. A UVES follow-up was soon attempted

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A. Renzini

Fig. 3 The metallicity distribution (from UVES) of planet hosting stars (shaded histogram) compared to that of an unbiased sample of stars [14]. A colour version of this figure is available at dx.doi.org/10.1007/978-1-4020-9190-2_1

for the brightest among these stars, checking whether the radial velocity curve can exclude that the transiting body is a brown dwarf, by definition more massive than ∼ 13 Jupiter masses (MJ ). Figure 4, right panel, shows that in one case the mass of the transiting body is < 3.8 MJ , and in another case the best fit mass (9.7 MJ ) is

Fig. 4 Left: the ACS color-magnitude diagram of the bulge field over which 16 transiting planet candidates have been found, with the host stars shown as empty circles. Right: the FLAMES/UVES radial velocity curves of the first and third brightest host stars, indicating the planetary nature of the transiting body [13]. A colour version of this figure is available at dx.doi.org/10.1007/978-1-4020-9190-2_1

VLT Science Highlights

7

also lower than the brown dwarf limit, yet not by much. Thus, in at least one case UVES has reasonably proven the planetary nature of the transiting body. The fainter candidate planet hosts in Fig. 4 are out of reach for UVES, but those brighter than V ∼ 22 may soon become feasible with the X-Shooter, as it is advertised as able to reach ∼ 2.5 magnitudes fainter than UVES.

6 VIMOS: Redshift Surveys The VIMOS multi-object spectrograph has been conceived as the redshift survey workhorse. To my knowledge, three main VIMOS surveys have been started at the VLT: the VVDS [6], the zCOSMOS [7] and the GOODS-South surveys (Popesso et al., in preparation). As O. Le Fèvre will present VVDS later at this meeting, I will report here on some lessons learned from the zCOSMOS project. Figure 5 shows the VIMOS LR-Blue coadded spectra of star-forming galaxies at 1.4 < z < 3 [7], that were selected via the BzK and UGR criteria (see respectively [2] and [16]). Notice that longward of Ly-α there are numerous, but quite weak absorption lines on top of the UV continuum. Figure 6 (left panel) shows the redshift

Fig. 5 The co-added VIMOS spectra of star-forming galaxies at z ∼ 2 in the COSMOS field, separately for those with Ly-α either in emission or in absorption [7]

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A. Renzini

Fig. 6 Left: the VIMOS spectroscopic redshift distribution of a first sample of zCOSMOS galaxies pre-selected with photometric criteria (BzK, BX, BM) for being at 1.4 < z < 3 [7]. Note the relatively small number of low redshift contaminants. The area of the blank rectangles refers to observed objects for which no redshift was obtained (in the same scale as the histograms). Right: the redshift distribution of galaxies observed at Keck with LRIS [16]. A colour version of this figure is available at dx.doi.org/10.1007/978-1-4020-9190-2_1

distribution for the first ∼ 1000 zCOSMOS BzK/UGR-selected galaxies. Notice the excellent efficiency at getting redshifts around 2, but the relatively few galaxies found at 1.5 < z < 1.8. This is in contrast with the fairly good success rate obtained by Steidel [16] using the LRIS spectrograph at the Keck telescope, as shown in Fig. 6 (right panel). This difference can be traced to the different UV throughput of VIMOS and LRIS, which drops to zero at λ ∼ 3500 Å in VIMOS, while it is still ∼ 30% in LRIS. Thus, for 1.5 < z < 1.8 LRIS spectra include Ly-α and VIMOS spectra do not. Missing the Ly-α handle, the only way of getting redshifts from VIMOS spectra is then via the weak absorptions shown in Fig. 5, a tough task with the noisy spectra of B ∼ 25 galaxies. The lesson is that exploring the 1.5 < z < 1.8 universe requires high UV throughput. The recently re-furbished FORS1 may help a little, but its multiplex is largely insufficient for any massive effort.

7 ISAAC: Widespread Compton-Thick AGN at z ∼ 2 The ISAAC K-band imaging coverage of the GOODS-South field is possibly the largest area/depth combination completed with this instrument, and likely the most widely used dataset ISAAC has so far produced. No paper has yet been published presenting these data, but the reduced data are public through the ESO archive since 2004. It is indeed worth mentioning that much more papers have been published using public GOODS data (HST, Spitzer, VLT, etc.) then by the GOODS Team itself: undoubtedly GOODS has been a prototype service to the community. One recent result based on the GOODS-South archival data offers the opportunity of illustrating the power of the full multiwavelength approach that is now becoming

VLT Science Highlights

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Fig. 7 Left: the identification of the Mid-IR Excess galaxies over the GOODS-South field for which ISAAC has provided the deep K-band imaging. Open and closed symbols refer to photometric and spectroscopic redshifts, respectively. Right: stacked Chandra soft and hard X-Ray images of the Mid-IR Excess and of normal galaxies [3]. A colour version of this figure is available at dx.doi.org/10.1007/978-1-4020-9190-2_1

possible. Daddi et al. [3] have selected K < 22 starforming galaxies at 1.4 < z < 3 via the so-called BzK criterion, and have determined their star formation rates (SFR) by a variety of techniques (from UV, mid-IR, sub-mm, radio, etc.). With one exception, all derived SFRs agree with each other within a factor of 2–3, with no systematic offsets. The exception is the SFR derived from the Spitzer/MIPS 24 µ flux. Figure 8 (left panel) shows the ratio RSFR of two SFRs: on the numerator is the sum of the SFR derived from the UV flux before correcting for dust extinction, plus the SFR derived from the 24 µ flux; on the denominator is the SFR derived from the UV flux after correction for dust extinction. Thus, the numerator is the total star formation rate as the sum of the unextinct and extincted SFRs, and the denominator is the total SFR as derived correcting for extinction the unextinct SFR. In absence of systematic errors this ratio RSFR should show a symmetric distribution peaking at unity, whose width is a measure of the SFR error. Figure 7 shows that the distribution does indeed peak at unity, but it is not symmetric: it is skewed towards large values, with a long tail reaching out to ∼ 100. Moreover, ∼ 1/3 of all galaxies have RSFR > 3, i.e. there appears to be more 24 µ dust emission than expected from the extinction measured in the UV. Figure 7 (right panel) shows the result of stacking the Chandra X-Ray data, separately for the normal galaxies (RSFR < 3), and for the Mid-IR Excess ones (RSFR > 3): the Mid-IR Excess galaxies clearly stand out as hard X-Ray emitters, indicating that they harbor a Compton-thick AGN. This provides direct evidence for the co-evolution of galaxies and supermassive BHs,

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Fig. 8 The co-added FORS2 spectra of 10 passively evolving galaxies at 1.4 < z < 2 from the GMASS Large Programme at the VLT (PI A. Cimatti). A colour version of this figure is available at dx.doi.org/10.1007/978-1-4020-9190-2_1

at the time of their major mass growth, a discovery with a great deal of potential ramifications.

8 FORS2: Integrating for 360 Hours One of the most demanding goals in the observations of high-z galaxies is to map the evolution of passively evolving galaxies (PEG) with redshift. At redshifts lower than ∼ 1.4 the most prominent features in the FORS2 spectra of PEGs are the 4000 Å break and the CaII doublet. However, at higher redshifts these feature move to the near-IR, and the only spectral feature allowing us to get the redshift of a PEG is the so called MgUV feature, at λ 2600–2850 Å, due to a combination of MgI, MgII and Fe lines. To exploit this feature FORS2 was really forced to its limits with the GMASS project (PI A. Cimatti). Among other targets, several PEGs at 1.4 < z < 2 were observed with integration times of 30 h, and in a few cases up to 60 h. Figure 8 shows the composite spectrum of 10 PEGs at z = 1.58 (with a S/N corresponding to an integration of 360 hours!), compared to the composite spectrum of PEGs at z = 0.75. While this pilot study demonstrates the ubiquity of high-z PEGs, it also shows that with FORS2 one would need a prohibitive amount of VLT time to properly map the population of very high redshift PEGs over a field such as the COSMOS field of 1–2 square degrees. The dream of many among us would be to have, in the shortest possible time, a high multiplex zJ spectrograph, a simple, warm instrument allowing astronomers to do 4000 Å break science beyond z ∼ 1.4, and Ly-α science beyond z ∼ 7.

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Fig. 9 The SINFONI velocity map of a galaxy at z = 2.34 [5]. A colour version of this figure is available at dx.doi.org/10.1007/978-1-4020-9190-2_1

9 SINFONI: Mapping the Kinematics of z ∼ 2 Galaxies As high redshift galaxies are photometrically culled and their spectroscopic redshifts are measured, they can be fed as targets to SINFONI, the near-IR, AO-assisted, integral-field spectrograph. Figure 9 shows the first successful AO-assisted, 0 .15 spatial resolution mapping of the gas kinematics within a z ∼ 2 galaxy. It reveals a large rotating disk, with high velocity dispersion all across its face, and gas falling towards its nucleus, where an AGN is being fed, a quite intriguing object indeed. Now, having vindicated its capacities and being further powered by the new LGS facility, SINFONI is starting to provide data of similar complexity for significant samples of galaxies, as reported by N. Förster-Schreiber at this meeting. Several objects appears to be similar to that shown here, others are more likely undergoing a merging, and the still unanswered question is how such galaxies will manage to evolve into their lower-z descendants, the local early-type galaxies and the quiescent spirals.

10 VISIR: Subarcsecond Imaging in the Mid-IR The last few years have been characterized by the spectacular achievements of an 80 cm telescope: Spitzer. This small satellite has revolutionized the mid-IR astronomy, making it accessible to everybody and pushing it to the highest redshifts. This great success has largely eclipsed the Mid-IR instruments that at nearly the same time came in operation at the 8–10 m telescope on the ground. For what we are concerned, this has affected also VISIR, whose performance, especially for the spectroscopic mode, has been largely sub-optimal due to a poor detector. Yet, over Spitzer, VISIR offers subarcsecond imaging in the Mid-IR, as illustrated in Fig. 10 which

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Fig. 10 VISIR 17.6 µ (top) and 8.6 µ (bottom) images of Neptune, showing that the South Pole is currently the hottest region on the planet [10]. The resolution of these images is ∼ 0.15 arcsec. A colour version of this figure is available at dx.doi.org/10.1007/978-1-4020-9190-2_1

shows that, ironically enough, the hottest region on Neptune is its South Pole, a result of having enjoyed permanent sunshine over the last 40 year, still just midway of its ∼ 80-year long summer.

11 CRIRES The last instrument to come on line at the VLT was CRIRES, the high-resolution infrared spectrograph. As of today, only one paper has been published from CRIRES data, so there was no embarrassment when it came to select a highlight for this meeting. As an anticipation of the kind of fine abundance analyses that have become possible, Fig. 11 shows that the [S/Zn] abundance ratio of Galactic stars of all metallicities is just the same of that of Dumped Ly-α Systems at high redshift. This gives (in)direct evidence that both sulfur and zinc are produced only by massive stars, exploding as core-collapse supernovae.

12 A Few Short-Term Suggestions This meeting aims at brainstorming for the future of the VLT/VLTI well into the ELT eta. On a more modest mood, I would like to let a few very short-term suggestions to emerge from the experience gained from some of the projects and results highlighted here. This is a short wish list of what some of us would like to have now at the VLT, and would be ready to exploit.

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Fig. 11 The Sulfur-to-Zinc ratio of galactic stars from CRIRES spectroscopy, compared to dumped Ly-α systems at high redshift [9]. A colour version of this figure is available at dx.doi.org/10.1007/978-1-4020-9190-2_1

Redshift Surveys: VIMOS is a compromise instrument, that works reasonably well from the B to the I band, but its optics cuts the UV and in the red its detectors are far inferior to those of FORS2. After several years in its current configuration, VIMOS could be attractively rejuvenated by either changing detectors or optics. Yet, this would be hardly sufficient to satisfy needs that are already mature today for near-IR, high multiplex spectroscopy, needs that can only grow stronger as VISTA comes on line, generating zillions of potential targets for spectroscopic follow-up. High Resolution Spectroscopy: With the X-Shooter soon coming on line, quite a few urgent needs are going to be satisfied. But going deeper at very high resolution on the VLT may require placing a UVES/HARPS-type instrument at the VLT incoherent combined focus, if one can demonstrate that using 4 UTs simultaneously saves telescope time compared to using one UT for 4 times longer exposures. GIRAFFE is producing an enormous set of prime quality data, especially for stellar population studies. Its main limit is the very narrow spectral range covered at a time. A Super-GIRAFFE, covering a few thousand Ångstroms, as opposed to a few hundred, could be quite attractive for this kind of studies. Other Upgrades: Upgrades of existing VLT instruments are relatively cheap ways of opening new scientific avenues to the VLT. I believe that each VLT instrument should have its own Upgrade Plan, that could then be activated depending on the scientific pressure to do so, and the availability of the resources, either in the ESO budget or elsewhere. One obvious and urgent upgrade concerns VISIR, which would deserve a better detector.

References 1. L.R. Bedin et al., Astrophys. J. 605, L125 (2004)

14 2. 3. 4. 5. 6. 7. 8. 9. 10. 11. 12. 13. 14. 15. 16.

A. Renzini E. Daddi et al., Astrophys. J. 617, 746 (2004) E. Daddi et al., Astrophys. J. 670, 173 (2007) F. Eisenhauer et al., Astrophys. J. 628, 246 (2005) R. Genzel et al., Nature 442, 786 (2006) O. Le Fèvre et al., Astron. Astrophys. 439, 845 (2005) S. Lilly et al., Astrophys. J. Suppl. Ser. 172, 70 (2007) A.P. Milone et al., 0709.3762 [astro-ph], 2007 P.E. Nissen et al., Astron. Astrophys. 469, 319 (2007) G.S. Orton et al., Astron. Astrophys. 473, L5 (2007) G. Piotto et al., Astrophys. J. 621, 777 (2005) G. Piotto et al., Astrophys. J. 661, L53 (2007) K.C. Sahu et al., Nature 443, 534 (2006) N.C. Santos, G. Israelian, M. Mayor, Astron. Astrophys. 415, 1153 (2004) R. Schödel et al., Nature 419, 694 (2002) C.C. Steidel, Astrophys. J. 604, 534 (2004)

Pushing the VLT Spectroscopy of Distant Galaxies to the Limits and Future Prospects Andrea Cimatti

1 Introduction The aim of this paper is to illustrate a key scientific case in the field of galaxy formation and evolution which is relevant to highlight the current limits of 8m-class telescopes and discuss the prospects of future VLT and E-ELT instrumentation. The science case concerns the evolution of early-type galaxies (ETGs) (i.e. the ellipticals and lenticulars, E/S0). ETGs play a crucial role in cosmology. They are the most massive galaxies in the local Universe, contain most of the stellar mass and are primary probes to investigate the cosmic history and the physics of galaxy mass assembly. ETGs are also fundamental in tracing the evolution of the large scale structure and the co-evolution of spheroids and their central supermassive black holes. Although ETGs in the present-day Universe are rather simple and homogeneous systems in terms of morphology, colors, stellar population content and scaling relations [30], their formation and evolution is still a debated question. The most recent surveys suggest the most massive ETGs (stellar mass > 1011 M ) were already in place at z ≈ 0.7–0.8, with a number density consistent with the one at z = 0 whereas the evolution is more pronounced for the lower mass ETGs which may increase their mass from z ≈ 0.7–0.8 to z = 0 through the merging of disk and/or ETGs (e.g. [3, 5, 6, 9, 17, 28, 29, 33, 37]). This mass-dependent evolution is known as “downsizing” [12], i.e. with massive galaxies forming their stars earlier and faster than the low mass ones. It is unclear whether the downsizing can be extended to the stellar mass assembly evolution itself [6, 9] and if this may represent a significant problem for galaxy formation models where massive galaxies are expected to assemble their mass more gradually through hierarchical merging of CDM halos [14]. Beyond z ≈ 1 the picture is even more controversial because the spectroscopic identification and study of ETGs at these redshifts is often beyond the capabilities of 8m-class telescopes. ETGs have absorption line spectra and the most prominent spectral features (e.g. the D4000 continuum break and CaII H&K absorption lines) are redshifted at λ > 0.8–1 µm for z > 1–1.5, where ground-based spectroscopy is increasingly difficult due to the intense OH sky lines and telluric absorptions. Moreover, due to the strong k-correction, ETGs become rapidly very faint at z > 1 (e.g. A. Cimatti () Dipartimento di Astronomia, Universitá di Bologna, Via Ranzani 2, 40127, Bologna, Italy e-mail: [email protected] A. Moorwood (ed.), Science with the VLT in the ELT Era, Astrophysics and Space Science Proceedings, doi: 10.1007/978-1-4020-9190-2_2, © Springer Science + Business Media B.V. 2009

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I > 24, K > 21, AB) and it is extremely difficult to obtain high quality continuum spectroscopy either in the optical or near-infrared. Pushing the 8m-class telescopes to their limits, it has been possible to unveil some ETGs up to z ≈ 2 with deep spectroscopy done either in the optical [8, 10, 13, 24] or in the near-IR [20, 32]. However, despite the very long integration times, spectroscopy was limited only to a few brightest objects and to low spectral resolution, hence making it impossible to derive critical information like the dynamical masses through the velocity dispersion of the absorption lines. The distant ETGs spectroscopically identified to date at 1.5 < z < 2.5 are very red (R − K > 5–6), compact (re ≈ 0.1–0.2 arcsec), dominated by passively evolving old stars with ages of 1–4 Gyr, and have stellar masses typically > 1011 M , implying a star formation history characterized by strong (> 100 M /yr) and short-lived (0.1–0.3 Gyr) starbursts occurring at z > 2–3. The existence of these old, massive, passive ETGs z > 1.5 was unexpected in galaxy formation models available in 2004–2005 and, to date, no new predictions have been published and compared with the available observational data of highz ETGs. The existence of these galaxies opened the key question on how it was possible to assemble such systems when the Universe was still relatively young. It is generally thought that negative feedback from AGN might play an important role by “quenching” the star formation in ETG precursors (e.g. [25]).

2 Pushing VLT to the Limits: The GMASS Project The VLT was intensively used in the context of two ESO Large Programmes in order to place constraints on ETG evolution at z > 1 (K20 project, [7], and the GMASS project, [21]). In this paper we focus on the latter. GMASS (“Galaxy Mass Assembly ultra-deep Spectroscopic Survey”1 ) is a project based on an ESO VLT Large Program (PI A. Cimatti). The GMASS main scientific driver is to investigate the physical and evolutionary processes of galaxy mass assembly in the redshift range of 1.5 < z < 3, i.e. in the epoch when the crucial processes of massive galaxy formation took place. Photometric redshifts are not sufficient to fully address the above questions because they provide limited clues on the physical and evolutionary status of the observed galaxies. Spectroscopy is therefore essential to derive reliable spectroscopic redshifts, to perform detailed spectral and photometric SED fitting (with known spectroscopic redshift), and to characterize the nature and diversity of galaxies at 1.5 < z < 3. The GMASS sample was selected at 4.5 µm using the GOODS-South2 public image taken at that wavelength with the Spitzer Space Telescope equipped with IRAC (Dickinson et al., in preparation), and extracting from a region of 6.8 × 6.8 arcmin2 all the sources with m4.5 < 23.0 (AB). This flux-limited sample was then used to 1 http://www.arcetri.astro.it/~cimatti/gmass/gmass.html. 2 http://www.stsci.edu/science/goods.

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extract a sub-sample of target galaxies to be observed spectroscopically. To reach the GMASS scientific aims, spectroscopy was very deep in order to derive secure spectroscopic redshifts for the faintest galaxies and to obtain high-quality spectra for the brighter galaxies in order to allow detailed spectral studies. The GMASS optical multi-slit spectroscopy was done with the ESO VLT + FORS2 (MXU mode) and focused on galaxies pre-selected with a cut in photometric redshift of zphot > 1.4 in order to concentrate the study on galaxies in the critical range of 1.5 < z < 3. In order to make the spectroscopy feasible, two cuts in the optical magnitudes were adopted: B < 26.5 or I < 26.5 for spectroscopy done in the blue or in the red respectively. The integration times were very long (up to 32 hours per spectroscopic mask), and the spectroscopy was optimized by obtaining spectra in the blue (4000–6000 Å, grism 300V) or in the red (6000–10000 Å, grism 300I) depending on the colors and photometric SEDs of the targets. For both grisms, the slit width was always 1 arcsec and the spectral resolution λ/λ ≈ 600. Despite the faintness of the targets, GMASS spectroscopy provided an overall spectroscopic redshift success rate of about 85% for the targeted galaxies. The power and the novelty of the GMASS sample is the selection at 4.5 µm, which is crucial for two main reasons: (1) the peak of the stellar SEDs (λrest = 1.6 µm) is redshifted in the 4.5 µm band for z > 1.5, (2) it is sensitive to the rest-frame near-IR emission, i.e. to stellar mass, up to z ≈ 3. For m4.5 < 23.0 and a Chabrier IMF, the limiting stellar mass is log(M/M ) ≈ 9.8, 10.1, and 10.5 for z ≈ 1.4, z ≈ 2, and z ≈ 3 respectively, hence allowing to properly investigate the galaxy mass assembly evolution within a wide range of masses.

3 High-z Passive Galaxies The GMASS spectroscopic sample was mined to search for and study ETGs at z > 1.4. Thanks to the combination of the deep spectroscopy, HST + ACS imaging and multi-band photometric SEDs from 0.4 µm to 8 µm, it was possible to study a subsample of 13 ETGs at 1.4 < z < 2 [10]. However, despite the ultradeep integration times of 32 hours, the faintness of these objects prevented detailed spectral studies of individual galaxies (Fig. 1). Thus, the detailed spectral analysis was done by co-adding the 13 individual spectra to obtain a stacked spectrum with an equivalent integration time of 480 hours (Fig. 2). The main results can be summarized as follows (see [10] for more details). The GMASS ETGs have spectra and photometric SEDs dominated by old stars and very weak or absent star formation. The comparison of the stacked rest-frame UV spectrum with synthetic stellar population model spectra indicates a stellar age of ≈ 0.7–2.7 Gyr for a metallicity range of 0.2–1.5 Z . Extending the model fitting at longer wavelengths using near-infrared and IRAC photometry helps to reduce the age–metallicity degeneracy and indicates ages of ≈ 1–1.6 Gyr, Z = Z , e-folding timescales τ ∼ 0.1–0.3 Gyr, where SFR(t) ∝ exp(−t/τ ), AV ≈ 0 and stellar masses in the range of 1010–11 M . The specific star formation rates are consequently very low (≤ 10−2 Gyr−1 ).

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Fig. 1 Individual spectra of ETGs at z > 1.5 obtained with the VLT + FORS2 + grism 300I taken in the context of the GMASS project (see [10]). The typical magnitudes are I ≈ 24–25 (AB) and each spectrum was obtained with an integration time of 32 hours. The histogram line is the spectrum of the old galaxy LBDS 53w091 (z = 1.55; [16, 34]) used as a comparison. A colour version of this figure is available at dx.doi.org/10.1007/978-1-4020-9190-2_2

The HST + ACS morphological and surface brightness profile analysis indicate that the majority of the spectroscopically-selected passive galaxies have spheroidal morphologies consistent with being analogous to present-day ETGs. However, their sizes are smaller by a factor of ≈ 2–3 than at z ≈ 0, and imply that the stellar mass surface and volume internal densities are up to ≈ 10 and ≈ 30 times larger respectively [10, 22, 36, 38]. Submillimeter–selected galaxies are the only systems at z > 2 with sizes and mass surface densities similar to those of the passive galaxies at z ≈ 1–2 [35]. This suggests that an evolutionary link is present between these two galaxy populations. It is currently unclear how these superdense high-z ETGs can evolve from z ≈ 1.5–2 to z ≈ 0 and move to the present-day size–mass relation. The possible scenarios proposed so far include a size and mass growth evolution

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Fig. 2 The co-added spectrum of 13 ETGs at 1.4 < z < 2 obtained with the GMASS project compared with two synthetic spectra (histogram lines) [4, 23] which provide a good fit of the rest-frame UV spectrum for solar metallicity and age of 1 Gyr (see [10]). The question mark indicates an unidentified absorption feature at ∼3018 Å. The co-added spectrum has an equivalent integration time of 480 hours(!) and is publicly available at http://www.arcetri.astro.it/~cimatti/gmass/gmass.html. A colour version of this figure is available at dx.doi.org/10.1007/978-1-4020-9190-2_2

with mechanisms like major dry merging and/or envelope smooth accretion more or less rapidly depending on their mass and environment [1, 2, 11, 15, 18, 19, 26, 27].

4 Open Questions and VLT/E-ELT Requirements Larger samples of spectroscopically-identified ETGs in the range of 1 < z < 3 are needed to place stringent constraints on their physical, evolutionary, structural properties and to properly use them as cosmological probes. The key questions still open

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include the number density evolution, the stellar population content, the star formation and metallicity history, the evolution of scaling relations, the evolution of internal mass density, the environmental effects. Although this is very challenging from the observational point of view, 8m-class telescopes could still play a role before the advent of JWST and Extremely Large Telescopes (ELTs) if equipped with appropriate instrumentation. For instance, a multi-slit spectrograph with performances and sensitivity optimized in the zY J spectral region (possibly extended to the H -band) would provide competitive results. For 1 < z < 2, the key spectral features (e.g. the D4000 break and CaII H&K absorption lines) are redshifted in the zY J region. A spectral resolution R ≥ 1000 would be needed for adequate OH sky lines removal and measurement of absorption line velocity dispersions. The surface density of high-z ETGs is in the range of 0.1– 0.5 galaxies arcmin−2 , and a field of view of ≈ 10 × 10 arcmin−2 would provide high multiplexing. Due to the small angular sizes of these galaxies (0.1–0.2 arcsec), this instrument would benefit enormously from a ground-layer adaptive optic (GLAO) system providing an “improved seeing” with FWHM ≈ 0.1–0.3 arcsec across the field of view. This would allow to narrow the slits and strongly reduce the sky background while still keeping most of the target fluxes within the slits. With such an instrument, the VLT would be capable to obtain competitive results on the brightest envelope of the ETG population at 1 < z < 3. Moreover, if combined with more traditional optical spectroscopy (0.5–0.8 µm) like the one of GMASS discussed in this paper, the results would place even stronger constraints on the stellar population content and star formation history thanks to the coverage of both the rest-frame UV and optical regions. A similar instrument on the E-ELT, possibly extended to K-band in order to properly identify spectroscopically and study ETG candidates at higher redshifts (e.g. z > 4; [31]) would be essential, and appropriate also for many other studies of high-z galaxies. The ETG science case will also benefit from the synergy with JWST, which will provide rest-frame optical imaging of galaxies with superb spatial resolution over the widest redshift range, as well as low resolution integrated spectroscopy extended in the spectral regions unaccessible from the ground > 2.5 µm. Finally, we note that the current plans for E-ELT instrumentation do not include an optical multi-object spectrograph. This would preclude the possibility to perform spectroscopy in the rest-frame UV for high-z galaxies, hence making a 42 m telescope “blind” to a wide range of physical processes and diagnostics observable in the UV.

References 1. 2. 3. 4. 5. 6.

M. Boylan-Kolchin, C.-P. Ma, E. Quataert, Mon. Not. R. Astron. Soc. 369, 1081 (2006) F. Bournaud, C.J. Jog, F. Combes, Astron. Astrophys. (2007), in press A. Borch, K. Meisenheimer, E. Bell et al., Astron. Astrophys. 453, 869 (2005) G. Bruzual, S. Charlot, Mon. Not. R. Astron. Soc. 344, 1000 (2003) K. Bundy, R.S. Ellis, C.J. Conselice et al., Astrophys. J. 651, 120 (2003) K. Bundy, T. Treu, R.S. Ellis, Astrophys. J. 665, L5 (2003)

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Pushing FORS to the Limit—A New Population of Faint Extended Lyα Emitters at z ∼ 3 Martin G. Haehnelt, Michael Rauch, Andrew Bunker, George Becker, Francine Marleau, James Graham, Stefano Cristiani, Matt J. Jarvis, Cedric Lacey, Simon Morris, Celine Peroux, Huub Röttgering and Tom Theuns

1 Introduction We present here the results of an ambitious attempt to push searches for spatially extended emission from damped Lyα absorption sustems (DLAS) and Lyman limit systems (LLS) to fainter surface brightness limits with observations performed in service mode in periods of bad seeing. For this purpose we have initiated a program of long-slit spectroscopy with FORS2 at the VLT to search for extended Lyα emission in a mostly blank field (LP 173.A-0440). Because of the small volume covered this is not a viable solution for a search for objects as rare as Lyman break galaxies, which are hard to hit with a single randomly positioned slit, but the rate of incidence of LLS with neutral hydrogen column densities exceeding (N (HI) > 1019 cm−2 ) per unit redshift is approximately unity at redshift ∼ 3 i.e., the objects essentially cover the sky, and there should be numerous hits in a single setting for a typical long slit spectrograph. A major motivation for our deep spectroscopic search was thereby the prospect of reaching sufficiently low surface brightness levels to detect or place interesting upper limits on the the fluorescent Lyα emission from optically thick regions induced by the metagalactic ionizing UV background [1, 2]. The project resulted in (after overheads) 92 hours of on-source exposure, finally reaching a 1σ surface brightness detection threshold of 8 × 10−20 erg cm−2 s−1

−1 . The results presented here will be published in detail as Rauch et al. [3].

2 Observations The field around the QSO DMS B 2139-0405 (z = 3.32) was observed in the period 2004–2006 with the 1400 V grism at FORS2. The 2 wide slit gave a spectral resolution of λ/λFWHM = 1050. The spectrum (Fig. 1) on the detector ranges from 4457 to 5776 Å, with a midpoint at 5099 Å. The seeing conditions generally where not as bad as anticipated, with 89% of the seeing better than 1.5 . M.G. Haehnelt () e-mail: [email protected] A. Moorwood (ed.), Science with the VLT in the ELT Era, Astrophysics and Space Science Proceedings, doi: 10.1007/978-1-4020-9190-2_3, © Springer Science + Business Media B.V. 2009

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Fig. 1 Two-dimensional spectrum obtained in 92 hours of exposure time with FORS2 at VLT. The line emitter candidates for HI Lyα are enclosed by the numbered boxes. The dispersion direction is horizontal, with blue to the left and red to the right; the spatial direction along the slit is vertical. The figure is adopted from Rauch et al. [3]

3 Results The white boxes in Fig. 1 mark faint emission line objects. Figure 2 shows a close up of our sample of 27 faint candidate Lyα emitters. At least a third of the sample shows emission line profiles or an association with absorption systems in the nearby QSO, strongly suggesting identification with Lyα. Spectroscopic features and the absence of detected continua down to 3 − σ flux limits of ∼ 1.5 × 10−19 erg s−1 cm−2 make a direct identification of the other emitters (as HI Lyα, [OII] doublet, or [OIII]/HI Balmer emission lines) difficult, but comparison with known galaxy populations and other statistical arguments indicate that the majority of emitters is likely to be Lyα at mean redshift 3.2. If this identification is correct, the emitters present a steeply rising luminosity function with a total number density more than 20 times larger than the comoving density of Lyman break galaxies (MR < 25.5) at comparable redshifts (Fig. 3). About half of the profiles are extended, possibly owing to radiative transfer of Lyα

Pushing FORS to the Limit—A New Population of Faint Extended Lyα Emitters at z ∼ 3

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Fig. 2 Spectra of individual line emitters (15.12 or 116 proper kpc wide in the spatial direction and about 2266 km s−1 long in the spectral direction. The spectra have been smoothed with a 7 × 7 pixel boxcar filter. The areas within the light grey (turquoise in the colour version) contours have a flux density greater than approximately 1.5 × 10−20 erg cm−2 s−1 Å. The first 12 of the spectrs show a single central peak; the next six show a clearly asymmetric red peak with a much weaker blue counter-peak; the following three either have a stronger blue than red peak (ID 15) or emission features blueward of an absorption line (36, 37); the remaining six spectra are unclassifiable, sometimes amorphous objects. The figure is adopted from Rauch et al. [3]. A colour version of this figure is available at dx.doi.org/10.1007/978-1-4020-9190-2_3

photons from a central source, and there are candidates for both outflows and infall features. We have investigated several mechanisms for the Lyα production and find star formation to be the energetically most viable process, with a few objects being candidates for cooling radiation. The inferred star formation rates range from 7 × 10−2 to 1.5 M yr−1 for standard assumptions.

4 Conclusions The inferred low star formation rates, large line emission cross-sections, high number density, and a fitting total cross-section per unit redshift on the sky provide an excellent match to the low luminosities, low metallicities, low dust content, and rate of incidence of damped Lyα systems, the main reservoir of neutral gas at high redshift. This suggests that our objects are the long-sought counterparts of DLAS in

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Fig. 3 Cumulative UV continuum luminosity functions of Steidel et al. [4] (asterisks), the Hubble Ultra Deep Field (Bouwens et al. [5]; dashed line), and the cumulative distribution of our survey (solid line; dotted lines are ±1σ errors). The emitters are assumed to have a continuum magnitude predicted by their Lyα line flux assuming a rest-frame equivalent width of 70 Å. The space density of the detected faint emitters agree well with the space density of Bouwens et al. [5]. The absence of objects brighter than −21 is consistent with our small survey volume. The figure is adopted from Rauch et al. [3]

emission. The crucial next step will be to gain information on the rest-rame UV continuum of these objects which is expected to be extremely faint. This information will be most efficiently obtained by performing a similar spectroscopic survey in the HUDF.

References 1. 2. 3. 4.

C.J. Hogan, R.J. Weymann, Mon. Not. R. Astron. Soc. 225, 1 (1987) A. Gould, D.H. Weinberg, Astrophys. J. 468, 462 (1996) M. Rauch et al., Astrophys. J., in press, arXiv:0711.1354 (2007) C.C. Steidel, K.L. Adelberger, M. Giavalisco, M. Dickinson, M. Pettini, Astrophys. J. 519, 1 (1999) 5. R.J. Bouwens, G.D. Illingworth, M. Franx, H. Ford, Astrophys. J. 670, 929 (2007)

VIMOS Integral Field Spectroscopy of Gaseous Nebulae in Local Group Dwarf Galaxies E.V. Held, M. Gullieuszik, I. Saviane, F. Sabbadin, Y. Momany, L. Rizzi and F. Bresolin

The study of very metal-poor dwarf irregular (dIrr) galaxies is fundamental to test the cosmological scenarios of galaxy formation. Among Local Group galaxies, Leo A and SagDIG are probably the most metal-poor dwarfs, as suggested by estimates of their nebular abundances based on the empirical method [2–4]. Service mode observations of H II regions in these dwarf galaxies have been obtained in April 2006 and May 2007, using the Integral Field Unit of VIMOS. The observations consisted of 3 OBs of 50 min each for SagDIG, and 2 OBs for Leo A, taken with the medium resolution grism HR Blue, yielding an instrumental resolution R ∼ 2600. The spectral range covered by this grism includes Hβ, Hγ , the [OIII] doublet at 4959, 5007 Å, and the auroral line [OIII] 4363 Å. With this HR grism, the IFU field-of-view is 27 × 27 using the “wide-field” mode (0.67 per fibre). This is sufficient to cover the H II regions and some adjacent galaxy field free from nebular emission. Complementary observations with the grism HR Orange, needed to provide such important abundance and density diagnostic as the [NII] and [SII] lines, could not be taken. Integral Field spectroscopy provides unique spatial maps of the flux in emission lines, allowing the study of internal structure and physical conditions of the ionized nebulae in distant gas-rich, star-forming dwarfs. As an example, Fig. 1 shows our preliminary results for the nebula in SagDIG. Our reduction made use of the standard products of the ESO pipeline, while custom IDL scripts written by one of us (M.G.) were used for the analysis. Emission line relative fluxes were measured by Gaussian fitting on the strong lines in each spectrum/fibre, and the measurement were used to construct two-dimensional intensity maps as in Fig. 1. The H II region in SagDIG is nearly coincident with a peak in the H I distribution, suggesting it represents the photo-ionized section of a larger gas distribution rather than an isolated nebula. The possible sources of ionizing photons in the SagDIG nebula are discussed in [1]. By co-adding all spectra with [OIII] 5007 Å line flux above a threshold, we obtained the integrated spectrum of each H II region. For SagDIG, this led to the detection of the auroral [OIII] 4363 Å line and O abundance calculations with a direct E.V. Held () Osservatorio Astronomico di Padova, INAF, vicolo dell’Osservatorio 5, Padova 35122, Italy e-mail: [email protected] A. Moorwood (ed.), Science with the VLT in the ELT Era, Astrophysics and Space Science Proceedings, doi: 10.1007/978-1-4020-9190-2_4, © Springer Science + Business Media B.V. 2009

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Fig. 1 Relative intensity maps of the Hβ and [OIII] 5007 Å emission lines in the H II region in SagDIG. Note the different intensity distributions of the two emission lines, probably indicative of an asymmetric density distribution and stratification of the emitting regions. A colour version of this figure is available at dx.doi.org/10.1007/978-1-4020-9190-2_4

method. By combining the IFU-derived relative strength of the [OIII] 4363 Å line with previous spectroscopy [2, 3], we could measure a very low metallicity for this galaxy which confirms previous indirect estimates. More details will be given in a forthcoming paper (Saviane et al., 2008, in prep.). In the case of Leo A, the partial IFU data obtained so far are not sufficient to detect the [OIII] Å line. An upper limit to the strength of the auroral line sets a lower limit 12 + [O/H] > 6.96 to the O abundance.

References 1. 2. 3. 4.

Y. Momany, E.V. Held, I. Saviane et al., Astron. Astrophys. 439, 111 (2005) I. Saviane, L. Rizzi, E.V. Held, F. Bresolin, Y. Momany, Astron. Astrophys. 390, 59 (2002) E.D. Skillman, R. Terlevich, J. Melnick, Mon. Not. R. Astron. Soc. 240, 563 (1989) L. van Zee, E.D. Skillman, M.P. Haynes, Astrophys. J. 637, 269 (2006)

Near IR Integral Field Spectroscopy of a Nearby Starburst L. Vanzi, G. Cresci, J. Melnick and E. Telles

1 Introduction We present integral field spectroscopy in the near infrared of II Zw 40, one of the most interesting nearby starburst galaxies. The radiation emitted by the galaxy is dominated at all wavelengths by very young stars, in particular one giant HII region in the core of the galaxy appears to be powered by a single young massive cluster. We observed this region with SINFONI at the VLT, using the H and K grisms and the 8 × 8 arcsec field of view. In Fig. 1 we show a HST ACS image of the galaxy in the F814W filter compared with the K continuum image obtained with SINFONI.

2 Results We analyzed the spectra of the brightest sources detected in the field and derived maps in the most prominent emission lines as Brγ , [FeII] and H2 , in addition we derived radial velocity and velocity dispersion maps, Fig. 2. We found that the young massive cluster powering the giant HII region has a mass of 1.7 × 106 M and an age of 3 Myr at most, a second fainter cluster turned to be older and less massive. From the ratios of the emission lines we found that the [FeII] and H2 are photoexcited, no evidence of shock excitation is observed consistently with the very young age and lack of SN. However while the [FeII] follows the geometry of Brγ , H2 has a quite different morphology. In addition the H2 is shifted in velocity by 90 km s−1 with respect to the other lines, so that we must be observing a giant molecular cloud not related to the giant HII region. We did not find any velocity feature that could be related to the large scale dynamics of the galaxy as depicted in the literature, so the dynamics of the ISM in the star forming region does not seem to follow or to be determined by a larger scale pattern, and as SNe are virtually absent, it must be fully determined by the star formation. We detected instead structures in the ISM, both in space and velocity, that L. Vanzi () ESO, Alonso de Cordova 3107, Vitacura, Santiago, Chile e-mail: [email protected] A. Moorwood (ed.), Science with the VLT in the ELT Era, Astrophysics and Space Science Proceedings, doi: 10.1007/978-1-4020-9190-2_5, © Springer Science + Business Media B.V. 2009

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Fig. 1 HST ACS image of II Zw 40 in F814W, in the insert at the botom left the SINFONI K continuum with a FOV of 8 arcsec. North is up east to the left

Fig. 2 Radial velocity maps of II Zw 40 in Brγ , [FeII] and H2 . The H2 has a significantly different velocity respect to the other components. The position of the main compact sources are marked by crosses, the cross to the north marks the position of the main cluster. A colour version of this figure is available at dx.doi.org/10.1007/978-1-4020-9190-2_5

are consistently oriented and we speculate that they are the effects of a common cause, as for instance a compression of the ISM that eventually triggered the star formation.

The ESO Large Programme “First Stars” P. Bonifacio, J. Andersen, S.M. Andrievsky, B. Barbuy, T.C. Beers, E. Caffau, R. Cayrel, E. Depagne, P. François, J.I. González Hernández, C.J. Hansen, F. Herwig, V. Hill, S.A. Korotin, H.-G. Ludwig, P. Molaro, B. Nordström, B. Plez, F. Primas, T. Sivarani, F. Spite and M. Spite

1 Introduction In ESO period 65 (April–September 2000) the large programme 165.N-0276, led by Roger Cayrel, began making use of UVES at the Kueyen VLT telescope. Known within the Team and outside as “First Stars”, it was aimed at obtaining high resolution, high signal-to-noise ratio spectra in the range 320 nm–1000 nm for a large sample of extremely metal-poor (EMP) stars identified from the HK objective prism survey [3, 4]. The goal was to use these spectra to determine accurate atmospheric parameters and chemical composition of these stars which are among the oldest objects amenable to our detailed study. Although these stars are not the first generation of stars they must be very close descendants of the first generation. One may hope to gain insight on the nature of the progenitors from detailed information on the descendants. The extremely metal-poor stars are very rare objects and finding them in large numbers requires specially designed surveys. All of the proponents of the large programme had been actively working on the medium-resolution follow-up of the HK survey (results still to be published), from either ESO La Silla, Kitt Peak or Roque de los Muchachos. Such a follow-up is mandatory in order to obtain a good list of candidates on which one can invest the time of an 8 m telescope. The programme was allocated a total of 39 nights between periods 65 and 68, these were split into 8 observational runs of unequal length. The observations were carried out in visitor mode because UVES was used in non-standard settings. The settings selected were Dic1 396 + 573 and Dic1 396 + 850, typically with a 1 slit for a resolution R ∼ 43 000. These settings were preferred over the standard Dic1 390 + 580 and Dic1 390 + 860, because you gain the Ba II 455.4 nm line in the blue, Zn I 471 nm in the red and Li I 670.8 nm in the reddest setting. The main results of the large programme are published on a series of papers “First Stars” on A&A, so far 10 papers have been published, one is in press, a few more in preparation. In addition a number of papers not in the “series” have been P. Bonifacio () CIFIST Marie Curie Excellence Team, 61, avenue de l’Observatoire, 75014 Paris, France e-mail: [email protected] A. Moorwood (ed.), Science with the VLT in the ELT Era, Astrophysics and Space Science Proceedings, doi: 10.1007/978-1-4020-9190-2_6, © Springer Science + Business Media B.V. 2009

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published, up to know there are 19 refereed papers published, which make use of the data acquired in the course of this large programme. In this contribution we highlight the main results of the large programme.

2 Uranium in a EMP Star The first surprise came quite early in the programme, Vanessa Hill was conducting the observations in August 2000 when she realised, from the quick look data, that the giant CS 31082-001 had an exceptional spectrum, characterised by extremely low metallicity and a large enhancement of the r-process elements. She was in fact able to identify immediately the Th II 401.9 nm line, which displayed a remarkable strength. This induced her in the following nights to acquire blue spectra of higher resolution with slicer #2 in the hope of being able to identify and perhaps measure uranium in this star. The star is now often colloquially dubbed as “Hill’s star”, and in fact her intuition proved correct since this was the first metal-poor star for which it was possible to measure the uranium abundance, opening up a new possibility for nucleochronology [6, 12]. This star actually showed how little we knew on the r-process. While Th/U proved to be a reliable chronometer Eu/Th and Eu/U provided unrealistically small ages, using the then available production ratios. Also Pb in this star is a real puzzle, in fact the majority of lead in this star is what you expect from the decay of Th and U, leaving very little space for Pb production during the r-process [13].

3 The Spite Plateau at the Lowest Metallicities Ever since Monique and François Spite discovered that warm metal-poor stars share the same Li abundance (the Spite plateau) [17, 18], there has been an active research on this field. What we wish to understand is if this plateau indicates the primordial Li abundance, as initially proposed [17, 18], or not. The “First Stars” large programme allowed to explore the Spite plateau at the lowest known metallicities. There are no known dwarf stars with a metallicity (meant as Z, total metallicity, not [Fe/H]) lower than the stars shown in Fig. 1. The data are those of [5] and [11], the picture which emerges is that the plateau seems to continue at the lowest metallicities. It is possible that there is a larger scatter, however the impact on this picture of stars in which lithium may have been partially depleted is yet unclear. The difference in lithium content between the two components of the binary system CS 22876-32 has no clear explanation. The cooler component (star B) has an effective temperature of 5900 K and should not display Li depletion according to standard models. From the cosmological point of view there is a tension between the value of the Spite plateau, A(Li) ∼ 2.1 and the value predicted by standard big bang nucleosynthesis, when the baryonic density derived from the power spectrum of the

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Fig. 1 The Spite plateau at the lowest metallicities as portrayed by the “First Stars” data. The stars whose names are labelled are binaries, for CS 22876-32 an orbital solution is available and the analysis has been done taking properly into account the veiling and Li in both components has been measured, for CS22957-15 this has not been possible, due to the lack of the necessary data, however the correction for the veiling is likely not very large

fluctuations of the cosmic microwave background is used [5, 16], A(Li) = 2.64. Several ways to explain this discrepancy have been suggested, and generally they go in two possible directions: (a) the Spite plateau does not represent the primordial abundance or (b) primordial nucleosynthesis did not proceed as assumed in the “standard” model. At present both solutions are possible and further observations of EMP stars, to understand if there is an excess scatter of Li at the lowest metallicities, could give useful indications.

4 Abundance Ratios, what Did We Learn? When we started the large programme, several of us, were expecting that at the lowest metallicities we would begin to see the effects of the pollution of very few supernovae (SNe), possibly a single supernova. As a consequence we were expecting considerable scatter in the abundance ratios, which would be the signature of the different masses of the polluting SNe and incomplete mixing of the gas in the early Galaxy. To the surprise of several of us we found instead that the majority of elements C to Zn display a remarkable uniformity, with well defined trends with metallicity [7]. The scatter in these trends can be totally explained by observational error. One explanation of this low scatter is an efficient mixing of the early Galaxy. Alternatively one could argue in favour of a narrow range of masses of SNe actually contributing to chemical enrichment. The exceptions, among lighter elements, were Na and Al, that displayed a star to star scatter larger than observed for other elements. This excess scatter also made the

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definition of trends somewhat ambiguous. A reanalysis of both elements using full NLTE line formation was in fact able to solve the problem [1, 2]. Sodium appears to be constant with metallicity among EMP stars, with [Na/Fe] = 0.21 ± 0.13, and the same is true for aluminium with [Al/Fe] = 0.08 ± 0.12. The observations of C and N, showed that in a [C/Fe] vs. [N/Fe] diagram the giants split nicely into two groups, one with high [N/Fe] and low [C/Fe], which we call “mixed”, the other with lower [N/Fe] and higher [C/Fe] which we call “unmixed” [19, 20]. As expected from the theory of stellar evolution “mixed” stars are typically the more luminous giants, although there are a few exceptions. At variance with lighter elements the n-capture elements display a large scatter, which cannot be explained by observational errors [10]. Such scatter, coupled with the very uniform ratios of the lighter elements demands an inhomogeneous chemical evolution. Already from the results of CS 31082-001 it was clear that the r-process is not “universal” and several r-processes may be needed. The data on n-capture elements clearly indicates that a second r-process is the main production channel at [Ba/H] < −2.5. Among the dwarf stars we found four which were C enhanced [14, 15]. In all of them the C-enhancement has come about as a consequence of mass-transfer from an AGB companion. The abundance pattern of n-capture elements is rather diverse among the stars, suggesting nucleosynthesis taking place under different physical conditions. The giant CS 22949-037 is one of the most extraordinary found in the course of the large programme [8]. With [Fe/H] ∼ −4.0 it is one of the most iron poor stars found in the sample, however its high abundances of CNO ([O/Fe] ∼ +2, [C/Fe] ∼ +1.2, [N/Fe] ∼ +2.6) make its global metallicity Z not so extreme as that of the four giants with [Fe/H] ∼ −4 [9], which are, so far, the star with the lowest Z known. There is no totally satisfactory model to explain the abundance pattern in CS 22949-037, however it is clear that some special kind of SN is needed to explain such an extraordinary pattern.

5 Needs for the Future One would like to extend the work done so far, with high resolution, high S/N ratio spectra of stars of even lower metallicities. Such stars should be found by on-going and projected surveys (SEGUE, LAMOST, SkyMapper. . . ). Most of these are however expected to be around 18th magnitude or fainter, UVES can work at these faint magnitudes, but. . . slowly. The proposed high resolution spectrograph ESPRESSO (see Pasquini this meeting) in the mode combining the 4 UTs, would be ideal for these targets. According to the preliminary estimates ESPRESSO 4 UTs, at a resolution of R ∼ 45 000 should beat UVES in efficiency for all targets fainter than V = 17.5.

The ESO Large Programme “First Stars”

References 1. 2. 3. 4. 5. 6. 7. 8. 9. 10. 11. 12. 13. 14. 15. 16. 17. 18. 19. 20.

S.M. Andrievsky et al., Astron. Astrophys. 464, 1081 (2007) S.M. Andrievsky et al., Astron. Astrophys., accepted T.C. Beers, G.W. Preston, S.A. Shectman, Astron. J. 90, 2089 (1985) T.C. Beers, G.W. Preston, S.A. Shectman, Astron. J. 103, 1987 (1992) P. Bonifacio et al., Astron. Astrophys. 462, 851 (2007) [Paper VII] R. Cayrel et al., Nature 409, 691 (2001) R. Cayrel et al., Astron. Astrophys. 416, 1117 (2004) [Paper V] E. Depagne et al., Astron. Astrophys. 390, 187 (2002) [Paper II] P. François et al., Astron. Astrophys. 403, 1105 (2003) [Paper III] P. François et al., Astron. Astrophys. 476, 935 (2007) [Paper VIII] J.I. González Hernández et al., Astron. Astrophys., accepted [Paper XI] V. Hill et al., Astron. Astrophys. 387, 560 (2002) [Paper I] B. Plez et al., Astron. Astrophys. 428, L9 (2004) T. Sivarani et al., Astron. Astrophys. 459, 125 (2006) [Paper X] T. Sivarani et al., Astron. Astrophys. 413, 1073 (2004) [Paper IV] D.N. Spergel et al., Astrophys. J. Suppl. Ser. 170, 377 (2007) M. Spite, F. Spite, Nature 297, 483 (1982) F. Spite, M. Spite, Astron. Astrophys. 115, 357 (1982) M. Spite et al., Astron. Astrophys. 430, 655 (2005) [Paper VI] M. Spite et al., Astron. Astrophys. 455, 291 (2006) [Paper IX]

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The Contribution of UVES@VLT to the New Era of QSO Absorption Line Studies Valentina D’Odorico and Miroslava Dessauges-Zavadsky Abstract We briefly review the main results obtained in the field of QSO absorption line studies with the UVES high resolution spectrograph mounted on the Kueyen unit of the ESO Very Large Telescope (Paranal, Chile).

1 Introduction Over the past decade, our understanding of the intergalactic medium (IGM) at high redshift, z = 2–5, the main baryonic component of the cosmic web, has advanced considerably. We are now able to measure properties of the diffuse baryons, among them the temperature, metallicity, kinematics and radiation field, and obtain their distributions as functions of time, spatial scale, and density. These improvements are primarily determined by a rich amount of new high-resolution observational data (mainly from UVES at the VLT and HIRES at Keck) and by new theoretical hydrodynamical simulations, that incorporate the relevant physical processes. These two factors have determined a paradigm shift in the study of absorption line systems: they are now considered as tracers of the entire cosmic structure formation process over the cosmic history, and not simple probes of physical processes taking place at the Jeans length (see [16] for a recent review).

2 Why UVES Made the Difference? The UV-Visual Echelle Spectrograph [6] is the high-resolution (R ∼ 40 000 with 1-arcsec slit) optical spectrograph of the ESO VLT. It started operation in fall 1999. The high efficiency from the atmospheric cut-off at 300 nm to the long wavelength limit of the CCD detectors (about 1100 nm) was determinant for UVES to excel among the equivalent instruments (HIRES at Keck and HDS at Subaru, see Fig. 1). In our view, two other features contributed to the excellent scientific results obtained with UVES: the dedicated data reduction pipeline [3], working since the beginning of operations, and the availability of observed data in the public archive of ESO which, in the specific case of the data rich QSO spectra, allowed different users to ‘squeeze’ all the possible science out of them. V. D’Odorico () INAF, Osservatorio Astronomico di Trieste, via G.B. Tiepolo 11, 34143 Trieste, Italy e-mail: [email protected] A. Moorwood (ed.), Science with the VLT in the ELT Era, Astrophysics and Space Science Proceedings, doi: 10.1007/978-1-4020-9190-2_7, © Springer Science + Business Media B.V. 2009

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Fig. 1 Efficiency curves for UVES, HIRES (before August 2004), and HDS including telescope, instrument and detector as a function of wavelength. Courtesy of S. D’Odorico. A colour version of this figure is available at dx.doi.org/10.1007/978-1-4020-9190-2_7

3 The IGM Probed by Single Lines of Sight A remarkable contribution to the study of the IGM with the Ly-α forest was given by the ESO Large Programme (LP) “Cosmic Evolution of the IGM” [4]. A sample of 19 QSO spectra at R ∼ 45 000 with S/N ∼ 35 and 70 per pixel at 350 and 600 nm, respectively, was collected covering the Ly-α forest between z ∼ 1.7 and 3.5. Data were immediately released to the public and, as of August 2007, they were used in 29 refereed publications generating 658 citations in 3 years.1 The following important results were obtained using this sample. The investigation of the metal content of the low density IGM with the “pixel optical depth statistics” (see details in [1]), to discriminate between early and late enrichment scenarios. Metal ions are detected down to the mean cosmic density, but the information in the under-dense regions, which is critical for the studied issue, is still poor (see Fig. 2 and [2]). The nature of the Ly-α forest was exploited to compute the transmitted flux [14] and the dark matter power spectra, allowing to tighten the values of the cosmological parameters derived from the CMB [26]. The statistical and physical properties of Ly-α (e.g. [13, 21]) and C IV (e.g. [22]) absorbers were refined and assessed. 1 Source:

“Telescope Bibliography” maintained by the ESO Library.

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Fig. 2 C IV (top panel) and O VI (bottom panel) optical depth plotted versus the HI optical depth for the LP QSO sample (from [2]). A colour version of this figure is available at dx.doi.org/10.1007/978-1-4020-9190-2_7

Another topic that was extensively investigated with UVES spectra was the timevariability of fundamental constants (see the contribution by P. Molaro to these proceedings).

4 High Redshift Galaxies Traced by Damped Ly-α Systems Damped Ly-α systems (DLA) are the quasar absorption line systems with the highest HI column densities (N (HI) > 2 × 1020 cm−2 ) and they arise in galactic disks or halos. They are invaluable tools to study the chemical abundances in the interstellar medium of objects in the very young Universe. At present, there are more than 60 QSO spectra with a DLA available in the ESO public archive. Some of the notable results obtained in this field are: the estimate of the temperature of the CMB radiation at z > 2, which is in agreement with the hot Big Bang cosmology predictions [17, 23]; the hints on the nucleosynthesis of Nitrogen in young objects (see Fig. 3, [5]), the properties of molecular hydrogen in DLA [15, 20, 24], and the star formation histories of individual DLA absorbers [7, 8]. The contribution of the sub-DLA population (N (HI) ≥ 1019 cm−2 ) to the total HI gas mass and to the missing metal problem was also investigated [18, 19].

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Fig. 3 [N/α] ratio as a function of Nitrogen abundance for DLA absorbers (from [9]). A colour version of this figure is available at dx.doi.org/10.1007/978-1-4020-9190-2_7

5 The Tomography of the IGM To reconstruct the 3D distribution of diffuse matter with its physical and chemical properties and investigate the small-scale IGM-galaxy interactions, many highredshift bright sources, close in the sky, are needed to use as background lights for spectroscopic studies. Several QSO pairs and groups were identified by the recent large surveys (2dFQRS, SDSS), however, most of them are too faint to be observed with UVES (V > 18). A great effort was put in collecting a statistically significant sample of UVES spectra of close QSO pairs that was used to investigate the transverse clustering properties of strong metal absorbers [10] and of the transmitted flux in the Ly-α forest [11]. The agreement of the transmitted flux correlation functions along and across the line of sight (see Fig. 4) implied that distortions in redshift space due to peculiar velocities are small (< 100 km s−1 ) and confirms the validity of the concordance cosmological model.

6 Future Perspectives In the second half of 2008, first light is planned for X-shooter at the VLT [25]: a single target, intermediate resolution (5000 at UVB and NIR, 7000 at VIS for

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Fig. 4 Comparison of the cross-correlation function for the sample of QSO pairs (squares) with the auto-correlation function computed for the LP QSO sample (crosses) as a function of comoving spatial separation across and along the line of sight, respectively. The cross-correlation function is slightly shifted in r for clarity (from [11])

1 arcsec slit) wide wavelength range (UV to K bands) spectrograph. This instrument was conceived to study QSO pairs and groups and will finally allow to carry out the Alcock–Paczy´nski test, sensible to the cosmological energy content (e.g. [12]). For the next decade, two new spectrographs are under study: CODEX at the ELT (see the contribution by J. Liske to these proceedings) and its precursor at the VLT, ESPRESSO (see the contribution by L. Pasquini to these proceedings). They will be characterised by a very high resolution (R ∼ 150 000), high efficiency, high stability and new design concepts. A second revolution in QSO absorption line studies is approaching fast!

References 1. 2. 3. 4. 5. 6. 7. 8. 9. 10.

A. Aguirre, J. Schaye, T. Theuns, Astrophys. J. 576, 1 (2002) B. Aracil, P. Petitjean, C. Pichon, J. Bergeron, Astron. Astrophys. 419, 811 (2004) P. Ballester, A. Modigliani, O. Boitquin et al., Messenger 101, 31 (2000) J. Bergeron, P. Petitjean, B. Aracil et al., Messenger 118, 40 (2004) M. Centurión, P. Molaro, G. Vladilo et al., Astron. Astrophys. 403, 55 (2003) H. Dekker, S. D’Odorico, A. Kaufer et al., SPIE 4008, 534 (2000) M. Dessauges-Zavadsky, F. Calura, J.X. Prochaska et al., Astron. Astrophys. 416, 79 (2004) M. Dessauges-Zavadsky, F. Calura, J.X. Prochaska et al., Astron. Astrophys. 470, 431 (2007) V. D’Odorico, P. Molaro, Astron. Astrophys. 415, 879 (2004) V. D’Odorico, P. Petitjean, S. Cristiani, Astron. Astrophys. 390, 13 (2002)

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V. D’Odorico, M. Viel, F. Saitta et al., Mon. Not. R. Astron. Soc. 372, 1333 (2006) L. Hui, A. Stebbins, S. Burles, Astrophys. J. 511, L5 (1999) T.-S. Kim, R.F. Carswell, S. Cristiani et al., Mon. Not. R. Astron. Soc. 335, 555 (2002) T.-S. Kim, M. Viel, M.G. Haehnelt et al., Mon. Not. R. Astron. Soc. 347, 355 (2004) C. Ledoux, P. Petitjean, R. Srianand, Mon. Not. R. Astron. Soc. 346, 209 (2003) A.A. Meiksin, submitted to Rev. Mod. Phys., arXiv:0711.3358 (2007) P. Molaro, S.A. Levshakov, M. Dessauges-Zavadsky, S. D’Odorico, Astron. Astrophys. 381, L64 (2002) C. Péroux, M. Dessauges-Zavadsky, S. D’Odorico et al., Mon. Not. R. Astron. Soc. 363, 479 (2005) C. Péroux, M. Dessauges-Zavadsky, S. D’Odorico et al., Mon. Not. R. Astron. Soc. 382, 177 (2007) P. Petitjean, C. Ledoux, P. Noterdaeme, R. Srianand, Astron. Astrophys. 456, L9 (2006) F. Saitta, V. D’Odorico, M. Bruscoli et al., Mon. Not. R. Astron. Soc., accepted, arXiv:0712.2452 (2007) E. Scannapieco, C. Pichon, B. Aracil et al., Mon. Not. R. Astron. Soc. 365, 615 (2006) R. Srianand, P. Petitjean, C. Ledoux, Nature 408, 931 (2000) R. Srianand, P. Petitjean, C. Ledoux et al., Mon. Not. R. Astron. Soc. 362, 549 (2005) J. Vernet, H. Dekker, S. D’Odorico et al., Messenger 130, 5 (2007) M. Viel, M.G. Haehnelt, V. Springel, Mon. Not. R. Astron. Soc. 354, 684 (2004)

18. 19. 20. 21. 22. 23. 24. 25. 26.

IMAGES: A Unique View of the Galaxy Mass Assembly Since z = 1 M. Puech, F. Hammer, H. Flores, Y. Yang and B. Neichel

Abstract The Large Program IMAGES is near 2/3 of its completion. It provides us with kinematics (GIRAFFE deployable IFUs), gas chemistry (FORS2), detailed morphologies (HST/ACS) and IR photometry (Spitzer) for a set of 70 galaxies representative of intermediate mass galaxies (MJ ≤ −20.3 or 1.5 × 1010 M ) at z = 0.4–0.75. We discover that, 6 Gyr ago, a significant fraction of galaxies (≥ 40%) had anomalous kinematics, i.e. kinematics significantly discrepant from those of rotational or dispersion supported galaxies. The anomalous kinematics cause the observed large dispersion of the Tully–Fisher relation at large distances. IMAGES will soon allow us to study distant galaxies at a level of detail almost comparable to that of nearby galaxies.

1 Introduction: Intermediate-Mass Galaxies It is now relatively well-established that ∼ 50% of the present-day stellar mass has been formed since z = 1. Most of this stellar mass has been formed in intermediatemass galaxies (3 × 1010 to 3 × 1011 M , i.e., ∼ L∗ galaxies), as a result of strong star formation episodes during which galaxies take the appearance of luminous infrared galaxies (LIRGs, see [3]). However, it is still unclear what physical processes have driven this evolution. To this respect, internal kinematics of distant galaxies is a powerful tracer of the major processes governing star-formation and galaxy evolution in the early universe such as merging, accretion, and feedback related to starformation and AGN. Robustly measuring the internal kinematics of distant galaxies is thus crucial for understanding how galaxies formed and evolved.

2 A Representative Sample of Intermediate Mass Galaxies A Large Program at VLT entitled IMAGES (Intermediate MAss Galaxy Evolution Sequence) is aiming to derive both resolved kinematics and integrated properties from VLT/GIRAFFE and FORS2 for galaxies selected in the Chandra Deep Field (CDFS), and to combine these observations with deep and high quality images M. Puech () ESO, Karl-Schwarzschild-Strasse 2, 85748 Garching bei München, Germany e-mail: [email protected] A. Moorwood (ed.), Science with the VLT in the ELT Era, Astrophysics and Space Science Proceedings, doi: 10.1007/978-1-4020-9190-2_8, © Springer Science + Business Media B.V. 2009

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from HST/ACS as well as with deep mid-IR photometry from SPITZER/MIPS. This Large Program has now reached about 2/3 of its completion, and recently, 39 additional galaxies observed with GIRAFFE have been analyzed [9]. Combined with previous observations during the GTO [1], it leads to a sample of 74 galaxies, which represents so far the largest existing sample of resolved kinematics for distant galaxies. Because in IMAGES we have deliberately selected MJ (AB) < −20.3 galaxies (i.e., with stellar masses larger than 1.5 × 1010 M ), we assume a similar limit for the combined sample. It let us with a sample of 63 galaxies which is well representative of the luminosity function at z = 0.4–0.75. Notice that the combined sample includes galaxies from 4 independent fields of view [9] and is then unaffected by field-to-field variations within Poisson statistics. Within this redshift range, GIRAFFE is able to recover the kinematics of almost all galaxies with W0 ([OII]) ≥ 15 Å.

3 Kinematics of Distant Galaxies At large distances the spatial resolution is not sufficient to resolve the central regions of the galaxies. It implies that the observed velocity dispersion (σ ) of a rotational body is the convolution of the actual random motions with the rotation. For a rotationally supported galaxy, it unavoidably leads to a well defined peak in the center (see Fig. 1). We have developed a classification scheme which allows us to compare any dispersion map to what it could be if it was a rotational disk [1]. Discrepancy from a rotational body can be measured from differences in amplitude and in position of the σ peak between the pseudo rotational σ map and the observed one, which leads to a very robust diagnostic diagram for the 63 observed galaxies [9]. Among the 63 galaxies of the representative sample, we find 20 rotating disks (32%), 16 rotating disks with perturbations (25%) and 27 galaxies with complex kinematics (43%, see Fig. 1 and [9]). Within this classification, perturbed rotations correspond to a discrepant σ peak from expectations for a pure rotation. The complex kinematics class corresponds to objects for which the large scale motions are not aligned to the optical major axis, and show σ map very discrepant from expectations for a rotation (see Fig. 1).

4 Comparison with Morphology Among these 63 galaxies, 52 have multi-band HST/ACS imaging. Neichel et al. (in prep.) have constructed a new procedure to classify the morphology of distant galaxies. This procedure relies on two important ingredients which are the used of color maps (which helps to overcome limitations due to k-morphological corrections) and the use of a visual decision tree. This last ingredient is particularly important, as they have shown that automatic methods can overestimate the fraction of spirals by

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Fig. 1 Examples of kinematics of z = 0.4–0.75 galaxies; each row corresponds to one galaxy. From left to right: HST ACS F755W/F814W image, observed velocity field and σ -map. The two top rows show regular rotating disk; the three bottom rows show galaxies with anomalous kinematics, one with just a shift of the peak in the σ -map, e.g. a perturbed rotation, the two other with dynamical axis misaligned relatively to the main optical axis. A colour version of this figure is available at dx.doi.org/10.1007/978-1-4020-9190-2_8

a factor of ∼ 2. They find a relatively good agreement between the two classifications: most of Rotating Disks are classified as Spirals, while most of galaxies having Complex Kinematics show a peculiar morphology. Combining morphology with spatially resolved kinematics allows us to study distant galaxies in very fine details: for the first time, we have recently detected a minor merger with a mass ratio of ∼ 1 : 18 in a z ∼ 0.6 galaxy classified as a perturbed rotator [7] (see Fig. 2). Such processes, which are predicted to be much more numerous than major mergers by numerical models, could provide us with a very plausible mechanism for explaining the kinematics of perturbed rotators.

5 Dynamical State of Distant Galaxies Let us now consider a representative sample of intermediate mass galaxies: at z = 0.6, 60% of galaxies have W0 ([OII]) ≥ 15 Å [2]. Let us assume that quiescent (W0 ([OII]) < 15 Å) galaxy are either rotationally or dispersion supported, thus

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Fig. 2 Direct identification of a minor merger in a z ∼ 0.6 galaxy classified as a perturbed rotator [9]. Left: 3-band B-V-z HST/ACS image, with the blue ellipse indicating the position of the infalling satellite. Middle: 3-band image with GIRAFFE isovelocities superimposed. The GIRAFFE IFU bundle is shown in white dashed-lines. Right: 3-band image superimposed with GIRAFFE velocity dispersion map isocontours. The dispersion peak is shifted by one GIRAFFE pixel compared to the position of the stellar continuum of the satellite, which is due to shocks between the gas stripped out of the satellite during the interaction, and the gas of the main progenitor [6]. A colour version of this figure is available at dx.doi.org/10.1007/978-1-4020-9190-2_8

minimizing the fraction of galaxies with anomalous kinematics. Then our results imply that 42 ± 7% of z = 0.4–0.75 galaxies have anomalous kinematics, including 26 ± 7% possessing complex kinematics. Because up to 97% of local intermediate mass galaxies are either E, S0 or spirals [3], it is likely that the fraction of anomalous kinematics is close to a few percents today. This leads to an extremely rapid evolution of kinematical properties of galaxies, with about 10 times more complex kinematics about 6 Gyr ago. The observed evolution of the Tully–Fisher Relation (TFR) provides us with a strong and independent confirmation. It has been suggested that galaxies with nonrelaxed kinematics (PR and CK) are responsible for the very large dispersions of the TFR at high redshift [1]. This result is confirmed by Puech et al. (in prep.) using the new sample of 63 galaxies: all the dispersion of the distant TFR can be accounted for by galaxies having non-relaxed kinematics. So it is beyond doubt that kinematics is among the most rapidly evolving properties of galaxies. Which physical process could explain such a dramatic evolution? Anomalous kinematics are linked with strong variations of the specific angular momentum consistent with a random walk evolution due to merging between galaxies, as predicted by the hierarchical scenario of galaxy formation [6]. Indeed, during a merger, and especially a major merger, galaxies pass through various stages during which the disk may be destroyed, generating significant discrepancies to the general behavior of isolated rotating disks.

6 Conclusion We do find a strong evolution of the galaxy kinematics since z = 0.6, with a significant fraction of galaxies with complex kinematics. Observations presented here tell us that major mergers could have played an important role in shaping galaxies as we observe them today. Indeed many estimates of the merger rate found that a

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typical intermediate mass galaxy should have experienced 0.5 to 0.75 major merger since z = 1 (see [5, 8] and references therein), and most of them since z = 2–3. A scenario in which merging is the dominant physical process explains many evolutionary features since z = 1, including the number density evolution of LIRGs and the emergence of galaxies with complex morphologies and with blue cores at z > 0.4 [3]. It is moreover particularly in agreement with the fact that the Milky Way has had an exceptionally quiet merger history [4]. Present-day M ∗ galaxies being mostly spiral galaxies, it is unlikely that they have all escaped a major merging since z = 2–3.

References 1. H. Flores, F. Hammer, M. Puech, P. Amram, C. Balkowski, Astron. Astrophys. 455, 107 (2006) 2. F. Hammer et al., Astrophys. J. 481, 49 (1997) 3. F. Hammer, H. Flores, D. Elbaz, X.Z. Zheng, Y.C. Liang, C. Cesarsky, Astron. Astrophys. 430, 115 (2005) 4. F. Hammer, M. Puech, L. Chemin, H. Flores, M. Lehnert, Astrophys. J. 662, 322 (2007) 5. J.S. Kartaltepe, D.B. Sanders et al. (2007), arXiv:0705.2266 [astro-ph] 6. M. Puech, F. Hammer, M.D. Lehnert, H. Flores, Astron. Astrophys. 466, 83 (2007) 7. M. Puech, F. Hammer, H. Flores, Y. Yang, B. Neichel, M. Rogrigues, Astron. Astrophys. Lett. (2007), accepted, arXiv:0711.0611 [astro-ph] 8. A. Rawat, F. Hammer, A.K. Kembhavi, H. Flores, Astrophys. J. Lett. (2007), submitted 9. Y.B. Yang, H. Flores, F. Hammer et al., Astron. Astrophys. (2007), accepted, arXiv:0711.2305 [astro-ph]

The Metallicity Evolution at High Redshift R. Maiolino, T. Nagao, A. Grazian, F. Cocchia and the Amaze Team

We present preliminary results of an ESO large programme (AMAZE) aimed at determining the evolution of the mass-metallicity relation in star forming galaxies at z > 3, by means of deep near-IR spectroscopy with SINFONI at the VLT. When compared with previous surveys, the mass-metallicity relation inferred at z ∼ 3.3 shows an evolution significantly stronger than observed at lower redshifts, even in massive systems. There are also indications that in low mass galaxies the metallicity evolution is stronger relative to high mass systems, an effect which can be considered the chemical version of the galaxy downsizing. When compared to theoretical models, the mass-metallicity relation observed at z ∼ 3.3 is difficult to reconcile with the predictions by simulations of hierarchical evolution of galaxies.

1 Introduction The correlation between galaxy mass and metallicity has been known for a long time, and it has been recently confirmed and refined with a sample of more than 50,000 star forming galaxies thanks to the SDSS survey [8]. The origin of this relation is ascribed to various possible processes. One possibility is that outflows, originated by the starburst winds, are responsible for ejecting metal enriched gas out of their host galaxies; the efficiency of this effect is lower in massive galaxies (because of higher gravitational potential) yielding higher effective enrichment in these systems. Another possibility is that low mass systems are little evolved: they have still to convert most of their gas into stars, and therefore have little enriched their ISM yet; instead, massive systems have converted most of their gas into stars, therefore reaching maturity from the chemical point of view. The latter scenario is commonly referred to as “downsizing”. Finally, it has been proposed that the IMF may change depending on the level of star formation, so that the effective yield of metals is higher during the evolution of galaxies with larger, final stellar masses. The relative role of these different processes in shaping the mass-metallicity relation is debated. However, it is likely that each of them contributes at least to some R. Maiolino () INAF – Osservatorio Astronomico di Roma, Via di Frascati 33, 00040 Monte Porzio Catone, Italy e-mail: [email protected] A. Moorwood (ed.), Science with the VLT in the ELT Era, Astrophysics and Space Science Proceedings, doi: 10.1007/978-1-4020-9190-2_9, © Springer Science + Business Media B.V. 2009

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extent, since observational evidence has been found for each of these phenomena. Each of these factors (outflows/feedback, downsizing, IMF) has profound implications on the evolution of galaxies. Therefore, it is clear that the mass-metallicity relation contains a wealth of information useful to constrain models of galaxy formation and evolution. Indeed, any model of galaxy evolution is now required to match the mass-metallicity relation observed locally [1, 4]. However, different models predict different evolutionary patterns of the mass-metallicity relation as a function of redshift, and observational data are required to test and discriminate among them. Observational constraints of the mass-metallicity relation have been obtained up to z ∼ 2.2 thanks to various deep surveys (e.g. [2, 7]). However, at z > 3, which is a crucial redshift range, the mass-metallicity relation has been little explored yet.

2 The AMAZE Project We have started a project (AMAZE, Assessing the Mass-Abundance redshift [Z] Evolution) specifically aimed at determining the mass-metallicity relation at z > 3. This is an ESO large programme that has been awarded 180 hours of observations with SINFONI, the VLT near-IR integral field spectrometer. The goal is to obtain near-IR spectra of a sample of about 30 galaxies at 3 < z < 5. At these redshifts the near-IR spectra allow us to measure the fluxes of the nebular emission lines [OII] 3727, [NeIII] 3869, Hβ, [OIII] 4959, [OIII] 5007, whose relative ratios can be used to constrain the gas metallicity [6]. In particular, by combining all of the diagnostics that can be inferred from these lines [6], we can estimate the gas metallicity by also accounting for the effect of possible dust reddening, and we can also control possible variations of the excitation conditions of the gas (details of this procedure are given in [5]). The sample has been selected among Lyman Break Galaxies for which SpitzerIRAC data are available, which are required to obtain a good determination of the stellar mass M∗ at z > 3 (where the rest-frame near-IR stellar light is redshifted to λ > 3.5 µm). AGN contamination must be absolutely avoided (since it would affect the line ratios), therefore we also required that the galaxies have hard X-ray and Spitzer-MIPS data, which allow us to identify the presence of even Compton thick AGNs [3]. The observing programme is currently in progress. Here we summarize preliminary results from the first 9 sources at 3 < z < 3.7 for which spectra have been obtained and reduced (a more detailed discussion is given in [5]). Figure 1 shows the rest frame, stacked spectrum of these 9 sources, where the nebular lines used to constrain the gas metallicity are indicated.

3 The Mass-Metallicity Relation at z ∼ 3.3 Figure 2 (rightmost panel) shows the inferred mass-metallicity relation at z ∼ 3.3 (red points with errorbars) compared with the local mass-metallicity relation ([8],

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Fig. 1 Rest frame stacked spectrum of the first nine sources at 3 < z < 3.7 observed within the AMAZE programme. [OII] 3727, [NeIII] 3869, Hβ, [OIII] 4959, [OIII] 5007 are the nebular lines used to determine the gas metallicity. The expected location of HeII (4686), typical of AGNs, is shown to highlight the absence of AGN contribution. A colour version of this figure is available at dx.doi.org/10.1007/978-1-4020-9190-2_9

blue thin lines, adapted to match the IMF adopted by us). The other panels show the mass-metallicity relation at lower redshifts inferred by using results of previous works [2, 7], and where metallicities have been re-determined by using the same set of intercalibrated diagnostics adopted by us [6] to ensure consistency between the various methods. Note that at z ∼ 2.2 the data points are obtained from stacked spectra [2], while at z ∼ 0.7 and at z ∼ 3.3 metallicities are inferred from the spectra of individual galaxies. The black, thick solid lines indicate quadratic fits to the data at each epoch.

Fig. 2 Mass-metallicity relation observed at various redshifts (red points with errorbars) compared with the local relation [8] (thin, blue solid and dashed lines). At z ∼ 0.7 (leftmost panel) we use the data from [7], recalibrated with our relations. At z ∼ 2.2 (central panel) we use the data from [2], also recalibrated with our relations. At z ∼ 3.3 (rightmost panel) we show the preliminary results for the first nine galaxies in our AMAZE program (additional details are given in [5]). Black, solid lines show quadratic fits to the relations at each redshift. A colour version of this figure is available at dx.doi.org/10.1007/978-1-4020-9190-2_9

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The mass-metallicity relation evolves significantly from z = 0 to z = 2, but only by a factor of about two in metallicity (at high masses), which is not a very strong evolution if one considers that this redshift range embraces 75% of the age of the universe. Between z = 2.2 and z = 3.3 the average metallicity decreases by another factor of ∼ 3, but here the temporal evolution is much stronger, since the time elapsed within this redshift range is much shorter (∼ 1 Gyr). This result indicates that at z ∼ 3 galaxies are experiencing a period of major action in terms of star formation and of metal enrichment, even in massive systems. Figure 2 also suggests that the evolution rate is not constant with mass. At low stellar masses the evolution is stronger than in massive systems. This can be regarded as the chemical version of the galaxy “downsizing”: massive systems reach chemical maturity at higher redshift, while low-mass systems chemically evolve more slowly and over a period of time possibly extending to the present epoch. However, confirming this effect requires more statistics at low masses. When interpreting Fig. 2 it is important to bear in mind that at different redshifts surveys are sampling different populations of galaxies. Local star forming galaxies sampled by SDSS in [8] are mostly spirals with modest star formation rates, while LBGs used to investigate the mass-metallicity relation at z ∼ 2–3 are characterized by enhanced star formation (and will likely evolve into massive, quiescent local galaxies). Therefore, the patterns shown in Fig. 2 should not be interpreted as the evolution of individual galaxies, but as the evolution of the average mass-metallicity relation of galaxies representative the density of star formation at each epoch. Figure 3 shows the comparison between the evolution of the mass-metallicity relation inferred by the observations (black solid lines) and the evolution expected by the simulations within a hierarchical scenario as obtained by [4] (green points). The dashed lines show the observed mass-metallicity relation with an offset of 0.2 dex, to account for a possible offset in the absolute metallicity scale (see [5] for a discussion about this possible issue). While at low redshift there is a fair agreement between model and observations, at high redshift an increasing discrepancy emerges. In particular, at z ∼ 3.3 observations appear inconsistent with the prediction provided by the theoretical simulations. Similar discrepancies are found by comparing observations with the hierarchical simulations obtained by [1]. The origin of the discrepancy between these hierarchical models and the observations at high redshift is unclear. One possibility, discussed in [1], is that SN feedback is not properly accounted for, and that the removal of metals by starburst winds is not effective enough in models. Another possibility is that in these hierarchical models most of the chemical evolution occurs rapidly in small units, at low masses (M∗ < 109 M ). Therefore, according to these models, at high redshift galaxies are mostly assembled with units that are already chemically evolved, yielding a relatively flat relation and relatively high metallicity at M∗ > 109 M . The lower metallicity found in the observations at z ∼ 3.3 suggests instead that galaxies at high redshift are made through the assembly of relatively unevolved sub-units.

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Fig. 3 Comparison of the evolution of the mass-metallicity relation expected by the hierarchical simulations in [4] (green points). The black, dashed lines show the observed mass-metallicity relation shifted by 0.2 dex to account for a possible offset of the metallicity scale [5]. Note the strong discrepancy between the model and observations at z = 3. A colour version of this figure is available at dx.doi.org/10.1007/978-1-4020-9190-2_9

References 1. 2. 3. 4. 5. 6. 7. 8.

M.E. de Rossi, P.B. Tissera, C. Scannpieco, Mon. Not. R. Astron. Soc. 374, 323 (2007) D.K. Erb et al., Astrophys. J. 646, 107 (2006) F. Fiore et al., arXiv:0705.2864 (2007) C. Kobayashi, V. Springel, S.D.M. White, Astrophys. J. 376, 1465 (2007) R. Maiolino et al., Astron. Astrophys. (2007) submitted T. Nagao, R. Maiolino, A. Marconi, Astron. Astrophys. 459, 85 (2006) S. Savaglio et al., Astrophys. J. 635, 260 (2005) C.A. Tremonti et al., Astrophys. J. 613, 898 (2004)

Near-IR Spectroscopy of Blue Supergiants N. Przybilla, A. Seifahrt, K. Butler, M.F. Nieva, H.-U. Käufl and A. Kaufer

Abstract Diffraction-limited observations with the E-ELT will make blue supergiants accessible to intermediate-resolution spectroscopy in galaxies out to the Virgo and Fornax clusters. BA-type supergiants (BA-SGs) will be primary targets for the study of the young stellar populations because of their enormous luminosities, their rich metal line spectra and their comparatively simple atmospheres. Quantitative analyses will allow stellar and galactochemical evolution to be investigated in different environments: in field galaxies, in galaxy groups and in clusters. As the use of adaptive optics techniques will restrain observations to the nearIR we are faced with the challenge of extending quantitative spectral analyses of BA-SGs to a wavelength regime that has so far been neglected. We report first results from a pilot spectroscopic investigation of Galactic BA-SGs at high resolution using CRIRES on the VLT. This is in order to establish quantitative near-IR spectroscopy of BA-SGs in preparation for the extragalactic stellar science programme with the E-ELT.

1 Introduction Massive blue supergiants are among the most luminous stars in spiral and irregular galaxies. Supergiants of late B and early A-type (BA-type supergiants, BA-SGs) are of particular interest for ground-based observations as they are the visually brightest normal stars. At absolute visual magnitudes up to MV −9.5 they can rival with entire globular clusters and even dwarf spheroidal galaxies. Consequently, the present generation of 8–10 m telescopes and efficient instrumentation facilitates high-resolution spectroscopy of individual BA-SGs in the galaxies of the Local Group, e.g. [1–4]. The rich metal-line spectra of BA-SGs allow abundances for a wide variety of chemical elements to be studied, for the light elements (He, CNO), iron group members and α- and s-process elements. This makes them invaluable for deriving observational constraints on stellar and galactochemical evolution, by tracing chemical signatures of mixing with nuclear-processed material or the study of abundance gradients/patterns in galaxies. Moreover, BA-SGs can act as distance indicators by the application of the wind momentum–luminosity or the flux-weighted gravity– luminosity relationship: WLR/FGLR [5, 6]. N. Przybilla () Dr. Remeis-Sternwarte Bamberg, Sternwartstr. 7, 96049 Bamberg, Germany e-mail: [email protected] A. Moorwood (ed.), Science with the VLT in the ELT Era, Astrophysics and Space Science Proceedings, doi: 10.1007/978-1-4020-9190-2_10, © Springer Science + Business Media B.V. 2009

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First studies of the blue supergiant populations in galaxies beyond the Local Group have been performed based on low-resolution spectra obtained with FORS on the VLT, e.g. [7–9]. This can potentially be done even in systems out to distances of the Virgo and Fornax cluster of galaxies [10] with existing large telescopes. However, crowding becomes an increasing problem beyond ∼ 7 Mpc for seeing-limited observations. Diffraction-limited observations using adaptive optics (AO) will overcome the problem, in particular with the advent of the next generation of extremely large telescopes (E-ELT, TMT, GMT). Use of AO-systems implies, however, that observations will be restricted to the near-IR wavelength range because of technological constraints. Fluxes of blue stars drop rapidly towards longer wavelengths. However, BA-SGs can still easily compete with red giants and even with some red supergiants in near-IR brightness. An advantage for the analysis are the comparatively simple (radiative) atmospheres of BA-SGs. Nonetheless, quantitative spectroscopy of BA-SGs in the near-IR—so far an essentially uncharted territory—needs to be established to prepare for the cuttingedge science to be done with e.g. the E-ELT. Only then, will highly detailed studies of the cosmic cycle of matter acting in all classes of late-type galaxies in the Hubble sequence become feasible—in the field, in galaxy groups and in galaxy cluster environments. Here, we report on the first investigation of Galactic BA-SGs at high spectroscopic resolution in the near-IR, using CRIRES [11] on the VLT/UT1. Preliminary results from the analysis of the near-IR spectra are presented after an overview of the status of BA-SGs modelling is given.

2 Quantitative Spectroscopy of BA-SGs The modelling of BA-SG spectra has relied for the most part on LTE techniques in the past. A more comprehensive non-LTE approach has been applied only recently [12–14]. In brief, our hybrid non-LTE methodology [13, 14] relies on line-blanketed model atmospheres and non-LTE line formation calculations using D ETAIL / S URFACE and sophisticated model atoms. Atmospheric parameters are determined iteratively by matching several independent indicators simultaneously: Starkbroadened hydrogen lines, multiple metal ionisation equilibria and fits to the spectral energy distribution. As a consequence, parameters and elemental abundances can be derived with unprecedented accuracy when high-S/N and high-resolution spectra are available. The 1σ -uncertainties in effective temperature Teff are of the order 1–2% and ∼ 0.05–0.10 dex in surface gravity log g. Non-LTE abundance uncertainties amount to typically 0.05–0.10 dex (random) and ∼ 0.10 dex (systematic 1σ -errors). The non-LTE computations reduce random errors and remove systematic trends in the analysis. It is found that LTE analyses tend to systematically underestimate iron group abundances and overestimate the light and α-process element abundances by up to factors of 2–3 on the mean. This applies to analyses of the visual spectra.

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Little information on quantitative analyses of near-IR spectra of BA-SGs can be found in the literature. Problems with the models were found in the pioneering study of the hydrogen Paschen and in particular the Pfund series in the A-supergiant prototype Deneb [12], which were attributed to deficiencies in our understanding of the physics of stellar atmospheres. It was shown later [15] that this was in fact a result of inaccurate electron-impact excitation cross-sections used in the hydrogen model atom by [12]. Non-LTE effects can become amplified in the Rayleigh-Jeans tail of the energy distribution, i.e. in the near-IR in hot stars. This is because the line source functions can be highly sensitive to even small variations of the departure coefficients of the lower/upper levels involved in a transition. Consequently, spectral lines are extremely sensitive to the atomic input data. An example is shown in Fig. 1, where three model calculations are compared to observation (the stellar parameters were determined from an analysis of the optical spectrum). The LTE approach underestimates the strengths of the hydrogen lines in the near-IR, while non-LTE calculations using approximation formulae for collisional bound-bound transitions lead to an overestimation. Consistency—also for all other observed hydrogen lines—can only be achieved when data from ab-initio computations are used [15]. Similar sensitivities to details of the model calculations are also found for helium lines in hot stars [16, 17]. Beyond this, little evidence for the presence of metal lines is found in the available low- and intermediate-resolution spectra for BA-SGs. This motivated us to obtain high-resolution near-IR data of BA-SGs in order to (I) identify the metal-line spectra and (II) investigate the adequacy of the existing analysis methodology for quantitative near-IR spectroscopy of BA-SGs.

Fig. 1 Comparison of an CFHT/FTS observation (R ≈ 10 000) of Br11 in β Ori with spectrum synthesis using different model assumptions [15]. See text for details. A colour version of this figure is available at dx.doi.org/10.1007/978-1-4020-9190-2_10

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Fig. 2 Preliminary spectrum synthesis to selected near-IR hydrogen lines observed with VLT/CRIRES. See text for details. A colour version of this figure is available at dx.doi.org/10.1007/978-1-4020-9190-2_10

3 Near-IR Spectroscopy of BA-SGs with CRIRES High-S/N spectra with wide wavelength coverage in the J, H, K and L bands have been obtained for three Galactic BA-SGs (HD 92207, HD 111613, η Leo). Only wavelength ranges with small contamination by telluric lines are accessible for analysis here, as the telluric correction turned out to be laborious. Besides the expected H and He I lines numerous metal lines are also present in the spectra, many of which are in emission (in contrast to the visual spectra, which show absorption lines). Lines of C I, N I, O I, Mg II and Fe II are identified, numerous other features remain unidentified. The most important chemical species for the study of stellar and galactochemical evolution are therefore available for analysis. The Brα line is affected by the stellar wind and shows emission characteristics analogously to Hα in the optical. We have used the results (model atmospheres, level populations) from our previous analysis of the stars in the optical [14] to extend the spectrum synthesis to the near-IR for several metals with sophisticated model atoms: CNO & Mg [18–21]. This is in order to test in how far our methodology is suited to applications at these wavelengths. The results are encouraging. Good agreement is found for H, N, O and Mg (Figs. 2 & 3). Considerable non-LTE strengthening of the lines is found in all

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Fig. 3 As Fig. 2, but for near-IR metal lines. See text for details. A colour version of this figure is available at dx.doi.org/10.1007/978-1-4020-9190-2_10

cases. The C I model atom requires a revision in order to match results from the optical and the near-IR analysis. We conclude that the first steps in the quantitative near-IR spectroscopy of BASGs have been successful. However, the project will require dedicated efforts in the next decade to thoroughly prepare the science to be done with e.g. the E-ELT. Further near-IR spectroscopy with the VLT will play a crucial rôle in advancing the state-of-the-art.

References 1. 2. 3. 4. 5. 6. 7. 8. 9.

J.K. McCarthy, D.J. Lennon, K.A. Venn et al., Astrophys. J. 455, L135 (1995) K.A. Venn, J.K. McCarthy, D.J. Lennon et al., Astrophys. J. 541, 610 (2000) K.A. Venn, D.J. Lennon, A. Kaufer et al., Astrophys. J. 547, 765 (2001) A. Kaufer, K.A. Venn, E. Tolstoy et al., Astron. J. 127, 2723 (2004) R.P. Kudritzki, J. Puls, D.J. Lennon et al., Astron. Astrophys. 350, 970 (1999) R.P. Kudritzki, F. Bresolin, N. Przybilla, Astrophys. J. 582, L83 (2003) F. Bresolin, R.P. Kudritzki, R.H. Méndez, N. Przybilla, Astrophys. J. 548, L159 (2001) F. Bresolin, W. Gieren, R.P. Kudritzki et al., Astrophys. J. 567, 277 (2002) M.A. Urbaneja, A. Herrero, F. Bresolin et al., Astrophys. J. 622, 862 (2005)

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10. R.P. Kudritzki, Blue Supergiants—Starway to Virgo and Fornax, in From Extrasolar Planets to Cosmology: The VLT Opening Symposium, ed. by J. Bergeron, A. Renzini (Springer, Berlin, Heidelberg, New York, 2000), pp. 236–241 11. H.-U. Käufl, P. Ballester, P. Biereichel et al., SPIE 5492, 1218 (2004) 12. J.P. Aufdenberg, P.H. Hauschildt, E. Baron et al., Astrophys. J. 570, 344 (2002) 13. N. Przybilla, Quantitative spectroscopy of supergiants, Ph.D. Thesis, Univ. Munich (2002) 14. N. Przybilla, K. Butler, S.R. Becker, R.P. Kudritzki, Astron. Astrophys. 445, 1099 (2006) 15. N. Przybilla, K. Butler, Astrophys. J. 609, 1181 (2004) 16. N. Przybilla, Astron. Astrophys. 443, 293 (2005) 17. N. Przybilla, K. Butler, U. Heber, C.S. Jeffery, Astron. Astrophys. 443, L25 (2005) 18. N. Przybilla, K. Butler, R.P. Kudritzki, Astron. Astrophys. 379, 936 (2001) 19. N. Przybilla, K. Butler, Astron. Astrophys. 379, 955 (2001) 20. N. Przybilla, K. Butler, S.R. Becker et al., Astron. Astrophys. 359, 1085 (2000) 21. N. Przybilla, K. Butler, S.R. Becker, R.P. Kudritzki, Astron. Astrophys. 369, 1009 (2001)

Integral Field Spectroscopy of Protoplanetary Disks in Orion with VLT FLAMES Y.G. Tsamis, J.R. Walsh and D. Péquignot

Abstract We discuss integral field spectroscopy of proplyds in M42 using the FLAMES Argus unit and report the first detection of recombination lines of C II and O II from the archetypical LV2 object. These lines can provide important new diagnostics of the physical conditions in proplyds. We also draw attention to the future capabilities of the MUSE spectrograph in relation to similar studies.

1 Rationale This contribution focuses on optical integral field spectrophotometry (IFS) of protoplanetary disks (proplyds) in M42 taken with the Argus unit of VLT FLAMES. The proplyds in M42 are partially ionized, low-mass embedded young stellar objects immersed in the extreme UV radiation field of the Trapezium cluster [1]. They represent a unique nearby environment for the study of ongoing star formation in a region dominated by main sequence OB-type stars. At optical wavelengths proplyds usually present a photoionized skin facing the Trapezium, giving way to a neutral dusty envelope which is shielded from the ionizing photons and is often shaped into ‘cometary’ radiation-bounded tails. IFS mapping of the proplyds and their immediate surroundings can yield new insight on the influence of small-scale ‘inhomogeneities’ on integrated spectra of distant galactic and extragalactic H II regions. With this programme (078.C-0247A) we specifically aimed at recording the faint optical recombination line (ORL) spectra of carbon (C II) and oxygen (O II) ions and use them, for the first time, as abundance diagnostics of the photoionized surfaces of the proplyds and their outflows. Simultaneous coverage of the strong collisionally excited lines (CELs) of [O III] would allow us to obtain a separate estimate of the oxygen abundances across the field and thus study the ‘abundance discrepancy problem’, whereby heavy element abundances, relative to H, from ORLs are found to be higher than the corresponding CEL-based abundances for classic H II regions such as M42 and 30 Doradus, by factors of up to ∼ 2 [2]. The resolution of this problem is of high priority as it casts uncertainty on classic CEL-based methods of abundance determinations for local and distant nebulae and galaxies. It has Y.G. Tsamis () Department of Physics and Astronomy, University College London, Gower Street, London WC1E 6BT, UK e-mail: [email protected] A. Moorwood (ed.), Science with the VLT in the ELT Era, Astrophysics and Space Science Proceedings, doi: 10.1007/978-1-4020-9190-2_11, © Springer Science + Business Media B.V. 2009

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been proposed that temperature fluctuations [3], density fluctuations, or zones of hydrogen-deficient plasma [4] within the nebulae may contribute to the ‘ORL vs. CEL’ problem. Our observations were aimed at discriminating between the various possibilities.

2 Observations and Data Reduction Method Argus is a rectangular array of 22 × 14 microlenses fed by optical fibers: we used the small configuration which provides a sampling of 0.30 arcsec2 /microlens and projects 6.6 × 4.2 arcsec2 on the sky, yielding 297 positional spectra per field of view. The targets were three relatively bright proplyds, including the archetypical object LV2 (167–317; [5]), and were selected from HST WFPC2 Hα and [O III] images of the nebula. The spectra were taken in service mode, and in sub-arcsec seeing, with the LR1–5 and HR1, 3, 4, 6, 8, 14B grating set-ups of the Giraffe spectrograph at resolving powers of ∼ 10 000–46 000. For LV2 the total exposure time in the LR/HR modes was 4344 and 3173 sec respectively; similar times were allocated to the other two targets. The data reduction was done with the girBLDRS pipeline developed by the Geneva Observatory. Custom-made IRAF scripts allowed us to construct data cubes and spectral line maps, and a dedicated χ 2 minimization routine was used to automatically fit Gaussians to the emission lines.

3 First Results At the time of writing preliminary monochromatic maps have been obtained for LV2. We succeeded in detecting the C II λ4267 and O II λ4649 ORLs from the head of the proplyd, and imaged the proplyd head and outflow in the light of Hα, and the [O III] λλ4363, 4959, [Ar IV] λλ4711, 4740, and [S II] λλ6716, 6731 CELs. The [O III] line ratio yielded an electron temperature (Te ) map and the [Ar IV] and [S II] ratios yielded electron density (Ne ) maps. In Fig. 1 spectra of LV2 are shown extracted over the proplyd’s tip (∼ 9 spaxels). The C II λ4267 3d–4f line is well detected in the LR2 setting, as are the numerous O II V1 multiplet 3s–3p lines around 465.0 nm in the HR6 setting. All these lines are useful abundance diagnostics [6], and the intramultiplet relative intensities of O II V1 lines are a Ne diagnostic of their emitting regions too [7]. In Fig. 1b the O II lines appear single peaked meaning that they mostly arise at the LV2 rest frame velocity: in contrast, the [Fe III] λ4658 line exhibits two additional velocity components associated with the proplyd’s bipolar outflow; further detections of [Fe III] lines, which are good tracers of shocked gas, will allow us to measure the electron density in the outflow. Analysis of the [O III] λ5007 line substructure at a spectral resolution of 9 km s−1 pix−1 indicates that the lobes of the bipolar jet emanating from the embedded protostar have line of sight velocities of approximately −100 and +80 km s−1 (see Fig. 2). It is likely, however, that the

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Fig. 1 LV2 (167–317) Argus spectra: (a) LR2 spectrum showing the C II λ4267 recombination line; (b) HR6 spectrum showing the O II V1 multiplet recombination lines. Note the velocity splitting of the [Fe III] λ4658 line with the red- and blue-shifted components arising from the proplyd’s outflow Fig. 2 LV2 (167–317) Argus spectrum showing the [O III] λ5007 line taken with the HR8 setting (at 9 km s−1 pix−1 ); note the blue- and red-shifted components arising from the proplyd’s outflow and which are absent in the background nebula

velocity structure will differ amongst various ionic species; a detailed kinematical analysis of emission lines from the numerous Giraffe HR settings will clarify this. The velocity resolution of this dataset at Hα is a factor of 2.3 higher compared to a Gemini-S GMOS IFU analysis [8] (which focused only on the red part of the optical spectrum and did not go deep enough to detect any heavy element ORLs). In Fig. 3 we show a Te ([O III]) map of LV2 and its immediate vicinity based on the λ4363/λ4959 ratio with the corresponding O2+ /H+ CEL abundance map based on the λ4959 line. Notably, the electron temperature over the proplyd appears to be as high as 15 000 K whereas the background temperature is only ∼ 8500 K. This is mainly due to collisional suppression of the relatively low critical density λ4959 line over the dense proplyd (at Ne > 7 × 105 cm−3 ). As a direct result of such pseudoTe fluctuations, the O2+ /H+ CEL abundance (Fig. 3 bottom) in the close vicinity of LV2 appears to be about a factor of 3 lower than in the background nebula.

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Fig. 3 LV2 (167–317) and M42 background physical conditions: (Top) The electron temperature map based on the [O III] auroral to nebular line ratio; (Bottom) The corresponding forbidden-line doubly ionized oxygen abundance map

An intricate combination of pseudo-variations in CEL-based abundances coupled with enhanced ‘metallic’ ORL emission from dense proplyds, and similar types of

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condensations, could at last provide a factual explanation to the long-standing abundance discrepancy problem in H II regions. Detailed studies with VLT FLAMES can improve our understanding of the protoplanetary disks themselves and help elucidate the influence of such objects on integrated nebular spectra. In conclusion, we note that the planned second generation VLT instrumentation will include MUSE, an adaptive optics assisted IFU spectrograph, which in wide field mode (1×1 arcmin2 ) will be able to take integral field spectra of the whole central Orion region (and of distant giant H II regions), and to sample simultaneously the entire proplyd population with 0.2 × 0.2 spatial pixels. In narrow field AOassisted mode the spatial resolution of 0.025 × 0.025 will allow unprecedented views of individual proplyds and protostellar outflows. It is unfortunate, however, that the currently planned wavelength coverage of MUSE is only 465.0–930.0 nm; this will miss the brightest optical recombination lines from C II, C III, N II, N III, O II, Ne II species (< 465.0 nm) which have in recent years opened an exciting new window to the chemistry and astrophysics of nebulae.

References 1. J. Bally, C.R. O’Dell, M.J. McCaughrean, Astron. J. 119, 2919 (2000) 2. Y.G. Tsamis, M.J. Barlow, X.-W. Liu, I.J. Danziger, P.J. Storey, Mon. Not. R. Astron. Soc. 338, 687 (2003) 3. A. Peimbert, Astrophys. J. 584, 735 (2003) 4. Y.G. Tsamis, D. Péquignot, Mon. Not. R. Astron. Soc. 364, 687 (2005) 5. P. Laques, J.L. Vidal, Astron. Astrophys. 73, 97 (1979) 6. Y.G. Tsamis, J.R. Walsh, D. Péquignot, M.J. Barlow, X.-W. Liu, I.J. Danziger, Messenger 127, 53 (2007) 7. R.J. Bastin, P.J. Storey, in Planetary Nebulae in our Galaxy and Beyond, ed. by M.J. Barlow, R.H. Méndez. Proceedings IAU Symposium, vol. 234 (2006), p. 369 8. M.J. Vasconcelos, A.H. Cerqueira, H. Plana, A.C. Raga, C. Morisset, Astron. J. 130, 1707 (2005)

MAD@VLT: Deep into the Madding Crowd of Omega Centauri G. Bono, A. Calamida, C.E. Corsi, P.B. Stetson, E. Marchetti, P. Amico, P.G. Prada Moroni, I. Ferraro, G. Iannicola, M. Monelli, R. Buonanno, F. Caputo, M. Dall’Ora, S. Degl’Innocenti, S. D’Odorico, L.M. Freyhammer, D. Koester, M. Nonino, A.M. Piersimoni, L. Pulone and M. Romaniello

Abstract We present deep and accurate Near-Infrared (NIR) photometry of the Galactic Globular Cluster (GC) ω Cen. Data were collected using the MultiConjugate Adaptive Optics Demonstrator (MAD) on VLT (ESO). The unprecedented quality of the images provided the opportunity to perform accurate photometry in the central crowded regions. Preliminary results indicate that the spread in age among the different stellar populations in ω Cen is limited.

1 Introduction Quantitative constraints concerning the evolutionary properties of low-mass stars mainly rely on the comparison between predicted and observed Color-Magnitude Diagrams (CMDs) of GCs. The GCs present several key advantages when compared with field stars: (i) cluster stars typically present the same age and the same chemical composition; (ii) cluster stars are located at the same distance, since the depth effects are negligible, and present the same reddening; (iii) cluster stars in a CMD are distributed according to their evolutionary status, (consecutio), therefore, they are redundant systems; (iv) cluster cores host a zoo of compact objects: Cataclysmic Variables [6], Millisecond pulsars, Low-Mass X-ray Binaries [7] and low/intermediate-mass black hole [1]. The main drawback of GCs is that quite often more than half of the cluster stars are located in the innermost regions, and indeed the half mass radius of massive clusters is at most a few arcminutes. This is the reason why accurate and deep photometry of the innermost regions of GCs became chore only with the superb spatial resolution and image quality of the Hubble Space Telescope (HST) optical images. Accurate photometry in the crowded central regions of GCs is not a trivial effort even by using HST images. The GC ω Cen seems to be a peculiar system. It is the only one to show a well defined spread in the abundance of iron and α-elements, thus suggesting that it G. Bono () INAF-OAR, via Frascati 33, 00040 Monte Porzio, Italy e-mail: [email protected] A. Moorwood (ed.), Science with the VLT in the ELT Era, Astrophysics and Space Science Proceedings, doi: 10.1007/978-1-4020-9190-2_12, © Springer Science + Business Media B.V. 2009

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Fig. 1 Top—from left to right: the 3D shape of a group of five MS stars in a F 435W ACS image of 340 sec. Analytical Moffat PSF model of the stars plotted in the left panel. Top view of the same stars. The red plus marks the brightest star, while the blue plus the star that was identified in the MAD image. Residuals after the subtraction of the analytical PSF to the data. Bottom—same as the top panels, but in a Ks MAD image of 120 sec. Note the difference in spatial resolution of optical and NIR images and the smooth distribution of the residuals. A colour version of this figure is available at dx.doi.org/10.1007/978-1-4020-9190-2_12

might be a possible link between GCs and dwarf galaxies. It is also the most massive GC. This means that ω Cen is a perfect laboratory to investigate fast evolutionary phases [3]. The use of optical and NIR photometry of cluster stars presents several advantages when compared with either optical or NIR photometry. The opticalNIR colors can be adopted to pinpoint peculiar stellar populations that present either different ages or different chemical compositions [9].

2 Observations and Data Reduction Optical data were collected with the Advanced Camera for Surveys (ACS) on board the HST and the reduction strategy was already discussed by Castellani et al. [4] and by Calamida et al. [3]. MAD is a prototype instrument performing wide Field of View (FoV) real-time correction for atmospheric turbulence. The reader interested in more details concerning the MCAO techniques and MAD characteristics is referred to Gilmozzi and Spyromilio [8] and Marchetti et al. [10]. During the first on-sky demonstration run of MAD two 1 × 1 arcminute fields were observed in the region across the center of ω Cen. Five images of 5 × 24 s each were collected in Ks -band and three images of 5 × 24 s each in J -band, while four images of 10 × 24 s were collected in Ks -band and three images of 10 × 24 s in J -band. The seeing during the observations of the first night changed from 0.7 to 0.9 , while during the second it changed from 0.9 to 1.2 . The photometry was performed using DAOPHOTIV/ALLSTAR and ALLFRAME. We selected ≈ 100 isolated stars to estimate an analytical point-spread

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function (PSF) for each frame. We adopted a Moffat function for the PSF of the Ks band images with a quadratic positional change. For the J -band images we adopted for the PSF either a Moffat or a Lorentz function, and they were assumed linearly variable across the chip. The NIR and optical catalog includes ≈ 7.5 × 105 stars. The comparison between the optical and NIR catalogs indicates that on MAD images we identified on average more than 90% (Ks -band) and more than 75% (J -band) of the stars detected in the same cluster regions using ACS images. The photometric calibration of NIR data into the 2MASS system was performed using a sample of ∼ 5000 local standard stars [5]. We ended up with an optical-NIR catalog including ≈ 49 000 (Ks ) and ≈ 41 000 (J ) stars with at least one measurement in an optical and in a NIR band. The images plotted in the upper panels of Fig. 1 show a group of five Main Sequence (MS) stars in a F 435W -band ACS image. The exposure time of this image is 340 sec and it is located across the cluster center. The brightest object in the group is located a couple of magnitudes below the Turn-Off (TO) region (F 435W ∼ 19.9), while the faintest is almost two orders of magnitude fainter (F 435W ∼ 24.5). The lower panels show the same group of stars but in a Ks -band MAD image. The exposure time of this image is 120 sec. The stars quoted above present Ks magnitudes of ∼ 16.9 and ∼ 19.4 mag, respectively. The full-width-half-maximum (FWHM) of MAD images is typically better than 0.1 arcsec in the Ks -band and better than 0.25 arcsec in the J -band. This together with the good spatial resolution and image quality provided the opportunity to perform accurate ground-based PSF photometry in crowded cluster regions. The residuals of the fits (rightmost panels) are smaller than 1% both in the optical and in the NIR images. In passing, we note that the faintest star (blue plus) was firstly detected in the MAD images and then added in the fit of the ACS images. The NIR bands present several advantages when compared with the optical bands. (i) They are marginally affected by reddening uncertainties and by the possible occurrence of differential reddening [2, 11]. (ii) When moving from MSTO stars down to very-low-mass stars the range of NIR magnitudes covered by these structures is quite limited 17 ≤ Ks ≤ 21. On the other hand, the same structures in the F 435W and in the F 625W band cover 8 and 7 mag, respectively. The difference is due to the fact that less massive MS structures are also steadily cooler, therefore, their emissivity peaks in the NIR bands. Data plotted in Fig. 2 show the optical-NIR (Ks , F 435W − Ks ; J, F 435W − J ) CMDs of ω Cen based on MAD and ACS images. To our knowledge this is the deepest Ks , F 435W − Ks CMD ever collected for a GC. The fit with a cluster isochrone of 12 Gyr (red line) indicates that we detected MS stars with mass values M ≤ 0.3 M (Ks ≈ 21, left panel). The J -band photometry reaches similar limiting magnitudes, but the CMD is slightly shallower. In passing, we note that current photometry suggests that the age spread in ω Cen appears to be limited, and indeed the color age derivative at fixed metal content (Z = 0.0006) is (F 435W − Ks )/ age ∼ 0.06 mag/Gyr. In order to constrain the impact that MAD has on the photometry of crowded regions, Fig. 3 shows the optical-NIR CMD based on data of the same cluster regions collected with SOFI@NTT and with ISAAC@VLT. The exposure time of the

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Fig. 2 Left—Ks vs. F 435W − Ks CMD of ω Cen based on data collected with MAD@VLT and with ACS@HST. Current data extend from hot HB stars Ks ∼ 13.5, F 435W − Ks ∼ 1.0 down to the regime of very-low-mass stars Ks ∼ 21, F 435W − Ks ∼ 6.0. By assuming μ = 13.70, E(B − V ) = 0.11 and evolutionary prescriptions by Castellani et al. [4], in this region of the MS are located structures with a stellar mass M ≈ 0.3 M . Right—same as the left, but for the J vs. F 435W − J CMD of ω Cen. A colour version of this figure is available at dx.doi.org/10.1007/978-1-4020-9190-2_12

Fig. 3 Same as Fig. 2, but NIR data collected with SOFI@NTT and ISAAC@VLT. A colour version of this figure is available at dx.doi.org/10.1007/978-1-4020-9190-2_12

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NIR images is 240 (Ks ) and 180 (J ) sec for SOFI and 78 (Ks ) and 66 (J ) sec for ISAAC. The observing strategy we adopted to reduce these data is very similar to the approach discussed in Sect. 2. The NIR limiting magnitudes of this data set is ≈ 3–4 mag brighter than the MAD data set. The difference in the two data sets is mainly due to different seeing conditions: 0.5–0.7 arcsec (Ks , J ; SOFI), 0.4–0.6 arcsec (Ks , J ; ISAAC) versus 0.1–0.2 arcsec for MAD. An important role is also played by the spatial resolution ∼ 0.29 arcsec/px (SOFI), ∼ 0.15 arcsec/px (ISAAC) versus ∼ 0.03 arcsec/px (MAD). This has a twofold impact on the intrinsic accuracy of PSF photometry: (i) the improvement in the sampling implies a more accurate image deblending; (ii) the decrease in the pixel scale also implies smaller fluctuations in the sly background, and in turn, a more efficient identification of fainter sources. Preliminary results based on NIR data of ω Cen collected with MAD@VLT appear very promising in performing accurate and deep photometry in the innermost crowded regions of GCs. Up to now these regions have only been investigated using HST. Current Adaptive Optics systems have been developed for NIR bands. The use of these bands presents several advantages in constraining the evolutionary properties of very-low-mass stars and the cooling sequence of cluster white dwarfs (Calamida et al., 2008, in preparation).

References 1. 2. 3. 4. 5. 6. 7. 8. 9. 10. 11.

F.N. Bash et al., Astron. J. 135, 182 (2008) A. Calamida et al., Astrophys. J. 634, L69 (2005) A. Calamida et al., Astrophys. J. 673, L29 (2008) V. Castellani, A. Calamida, G. Bono et al., Astrophys. J. 663, 1021 (2007) M. Del Principe et al., Astrophys. J. 652, 362 (2006) P.D. Edmonds et al., Astrophys. J. 596, 1197 (2003) C.O. Heinke et al., Astrophys. J. 625, 796 (2005) R. Gilmozzi, J. Spyromilio, Messenger 127, 11 (2007) L.M. Freyhammer et al., Astrophys. J. 623, 860 (2005) E. Marchetti et al., Messenger 129, 8 (2007) J. Th, van Loon, Mon. Not. R. Astron. Soc. 382, 1353 (2007)

Chemical Evolution of the Galaxy and Supernova Yields after UVES G. Israelian and P. Bonifacio

Abstract The atmospheres of stars orbiting around black holes/neutron stars may contain a large amount of α and r-process elements as a result of contamination with the nucleosynthesis products in the matter ejected during any supernovae event associated to the formation of the compact object. The discovery of a strong overabundance of α-elements in the atmosphere of the companion in the massive black hole binary GRO J1655-40 has provided an unique way to establish the mass of the progenitor. Spectroscopy of other secondaries of selected LMXBs is crucial to establish an empirical relationship between the masses of black holes/neutron stars and those of their progenitors while the observations of radioactive elements can also provide a datation of any supernova event. Among the remarkable results provided by UVES at VLT is certainly the measurement of the isotopic ratio of lithium in several metal-poor stars. The presence of 6 Li in the early Galaxy is difficult to explain, one should be however aware of the uncertainty of this measurement, especially in presence of the line asymmetries induced by convective motions.

1 Chemical Abundances of Secondary Stars in Low Mass X-Ray Binaries Low mass X-Ray binaries (LMXB), are an ideal astrophysical site to investigate the link between the explosive end of massive stars and the formation of black holes and neutron stars. The chemical composition of the atmosphere of the secondary star can provide a unique opportunity to investigate the nucleosynthesis of elements in any supernova event associated to the precursor of the compact object. The matter ejected during the explosion will pollute the environment and very likely contaminate the atmosphere of the low mass companion. A detailed abundance analysis may then show evidence for any nucleosynthesis products from the supernova which in turn will provide constraints on the mass of the progenitor. This will be of key importance for our understanding on how the latest stages of stellar evolution lead to the formation of black holes and neutron stars. It was reported [9] the discovery of a strong overabundance of oxygen and other α-elements (Si, Mg and S) in the atmosphere of the star orbiting around the massive black hole in the binary system GRO J1655-40 (Nova Scorpii 1994). A highresolution spectrum of this star obtained with the Keck I 10 m telescope reveals the G. Israelian () Instituto de Astrofísica de Canarias, 38200 La Laguna, Tenerife, Spain e-mail: [email protected] A. Moorwood (ed.), Science with the VLT in the ELT Era, Astrophysics and Space Science Proceedings, doi: 10.1007/978-1-4020-9190-2_13, © Springer Science + Business Media B.V. 2009

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presence of strong absorption lines of O, Mg, Si and S. The analysis using spectral synthesis techniques and suitable model atmospheres indicated that these elements are 6 to 12 times more abundant in the star than in the Sun while the overall spectrum of the star resembles that of a normal subgiant star with solar abundances of iron group elements. We interpret the enhanced abundances of α-elements as a direct result of the nucleosynthesis in a Supernova explosion event associated with the progenitor of the black hole, providing direct evidence for a link between a supernova explosion and black hole formation. These results have been confirmed recently by González Hernández, Rebolo and Israelian [8] who used VLT/UVES to perform detailed abundance analysis and derive the stellar parameters. These authors found that the abundances of Al, Ca, Ti, Fe and Ni are consistent with solar values while O, Mg, Si and S and significantly enhanced in comparison with Galactic trends of these elements. They have concluded that the black hole in this system formed in hypernova explosion of a 30–35 solar mass progenitor with a mass cut in the range 2–3.5 solar masses. Recently González Hernández et al. [5–7] have carried out detailed spectroscopic analysis of A0620-00, Cen X-4 and XTE J1118+480 using the VLT/UVES and Keck I/HIRES. They computed synthetic spectra with the local thermodynamic equilibrium (LTE) code MOOG, adopting the atomic line data from the Vienna Atomic Line Database and using a grid of LTE model atmospheres (see for details [5]). Observed abundances have been compared with the expected abundances from yields of a variety of SN models of different metallicities, progenitor masses, explosion energies and geometries. The abundances in Nova Sco 1994 could be explained only if there was a SN explosion in the system. However, low Al abundance is not reproduced by current SN explosion models. A direct collapse cannot be ruled out in A0620 [5] while the abundances in this system are also consistent with a scenario where only the outer layers of the progenitor He core were ejected. Abundances and kinematics of Cen X-4 and XTE J1118+480 support agree with a SN explosion scenario. The progenitors of these systems were most likely born in the Galactic disc. Sadakane et al. [20] performed detailed spectroscopic studies for the secondary star in the black hole binary (micro-quasar) V4641 Sgr. These authors obtained abundances of 10 elements and found 0.8 dex over-abundances of N and Na. From line-by-line comparisons of eight other elements (C, O, Mg, Al, Ti, Si, Fe, and Cr) between V4641 Sgr and the two reference stars, they conclude that there is no apparent difference in the abundances of these elements between V4641 Sgr and their comparison targets. The analysis of these systems provides a relationship between masses of the compact objects and their progenitors (see Fig. 1). Let us remark that the fraction of the amount of captured matter by the secondary star is a free parameter in these models. This only allows to estimate a lower limit of the masses of the He cores. Secondary stars in LMXBs are faint and their high resolution spectroscopic analysis is only possible with 8–10 m class telescopes. The superior wavelength accuracy of ESPRESSO should allow isotopic ratio measurements in these targets.

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Fig. 1 Relationship between the masses of compact object in LMXBs and the progenitor stars. The crosses indicate the error bars. A colour version of this figure is available at dx.doi.org/10.1007/978-1-4020-9190-2_13

2 Lithium Isotopic Ratio The remarkable finding by Asplund et al. [1] that warm main sequence/turnoff stars display a measurable amount of 6 Li has attracted considerable attention over the last two years. On the one hand the presence of 6 Li, which is more fragile than 7 Li, places a strong constraint on any model which implies Li depletion. One has to do a careful tuning to deplete Li, but still preserve some 6 Li. On the other hand the presence of 6 Li at metallicities as low as 1/1000 solar pose serious problems to justify its creation. The lighter Li isotope is not formed in standard big bang nucleosynthesis. Up to now it has been assumed that its main production channel was cosmic ray spallation, like for Be and the two B isotopes. However at very low metallicities the ordinary spallation, which requires C, N and O nuclei is not effective and one has to rely essentially on the α–α fusion. The formation of 6 Li at the level reported [1] requires an energy of at least 1014 erg g−1 and it is difficult to find sources which can inject this amount throughout the volume of the Galaxy [19]. Many models have been suggested to explain these puzling observations and they range from new physics, implying a non-standard big bang nucleosynthesis [2, 10– 13, 16], to production in stellar flares [22], passing through production by cosmic rays generated by structure formation [17, 18, 21]. In spite of the impressive amount of theoretical work to explain the 6 Li observations one should keep in mind that the measurement is very difficult. The wavelength shift between the two isotopes is of 7.15 km s−1 (on a velocity scale), this may be compared to the thermal broadening in the atmosphere of a typical G or F dwarf, which is of the order of 6 km s−1 . It is thus obvious that in the spectra of stars 6 Li manifests itself as a slight extra absorption on the red wing of the Li feature. The situation is further complicated by the fact that the Li I 6708 Å resonance transition is in fact a doublet, whose two components are blended due to thermal broadening,

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Fig. 2 The Li I doublet in HD 74 000, fitted with a synthetic 3D NLTE profile (observed and theoretical data and from Cayrel et al. [3]). The best fit is shown as a green solid line, the individual 7 Li components are shown as blue dashed lines and the 6 Li components as red dashed lines. The best fit is obtained with no 6 Li. The asymmetry of the line due to convection adequately accounts for the extra absorption on the red wing of the feature. A colour version of this figure is available at dx.doi.org/10.1007/978-1-4020-9190-2_13

and has also hyperfine structure. The spectrum of the halo star HD 74 000 is shown in Fig. 2, from which it is clear that in a stellar spectrum the blend of the 7 Li and 6 Li doublets is seen as an unresolved feature. The spectrum shown in Fig. 2 has been taken with the HARPS spectrograph [14] on the ESO 3.6 m telescope with a resolution R = 120 000. It is well known that the profiles of spectral lines emerging from G and F stars are asymmetric [4] due to the Doppler effect of convective motions. Cayrel et al. [3] investigated the impact of these asymmetries on the measurement of the Li isotopic ratios. Both making use of an empirical profile, derived from unblended Fe I lines which arise from the same atmospheric layers as the Li I doublet, and of theoretical profiles computed from 3D hydrodynamical simulations with NLTE line transfer, they convincingly demonstrated that the convective asymmetry, if ignored, can mimic the presence of 6 Li at the level of a few per cent. This implies that it is mandatory to take into account the convective line asymmetries in order to measure the 6 Li content. This result, by itself, does not imply that all the existing measurements of 6 Li should be discarded, however a revision is warranted. At very minimum one should expect that all measurements will be revised downwards; some or even all the detections may have to be withdrawn. From the observational side high S/N ratio and high resolution is clearly necessary. However in the work of Cayrel et al. [3] the superbe wavelength stability and radial velocity accuracy of HARPS proved to be a great advantage. The spectrum shown in Fig. 2 is the result of the co-addition of twenty one hour exposures, an extremely stable wavelength scale allows to co-add the spectra to reach a very high S/N ratio, without introducing spurious line asymmetries. Such spurious line asymmetries, if present, would make the measurement of 6 Li impossible. Another important conclusion of Cayrel et al. [3] is that in the line profile fitting the isotopic ratio is degenerate with the velocity shift and it would be desirable to avoid having

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use the velocity shift as a free parameter. However this requires a radial velocity with an absolute accuracy of the order of 20 m s−1 . Molaro et al. [15] have investigated the radial velocity accuracy of UVES and HARPS using the spectra of asteroids. They concluded that while HARPS is accurate to the level of 1 m s−1 , UVES shows shifts up to 50 m s−1 , likely due to non-uniform slit illumination. This suggests that it would be extremely valuable to have a spectrograph with a performance like HARPS (or better!) at the VLT, since UVES is not designed for this. The current requirements for ESPRESSO at the VLT (see Pasquini, these proceedings) meet in fact this need. The superior wavelength accuracy of ESPRESSO should translate into higher S/N in the coadded spectra ad more reliable line profiles with respect to UVES. It is likely that ESPRESSO will have a major impact on the topic of isotopic ratios.

References 1. 2. 3. 4. 5. 6. 7. 8. 9. 10. 11. 12. 13. 14. 15. 16. 17. 18. 19. 20. 21. 22.

M. Asplund et al., Astrophys. J. 644, 229 (2006) C. Bird, K. Koopmans, M. Pospelov, 703, arXiv:hep-ph/0703096 (2007) R. Cayrel et al., Astron. Astrophys. 473, L37 (2007) D.F. Gray, The Observation and Analysis of Stellar Photospheres, 3rd edn. (Cambridge University Press, Cambridge, 2005) J.I. González Hernández, R. Rebolo, G. Israelian, J. Casares, A. Maeder, G. Meynet, Astrophys. J. 609, 988 (2004) J.I. González Hernández, R. Rebolo, G. Israelian, J. Casares, K. Maeda, P. Bonifacio, P. Molaro, Astrophys. J. 630, 495 (2005) J.I. González Hernández, R. Rebolo, G. Israelian, E.T. Harlaftis, A.V. Filippenko, R. Chornock, Astrophys. J. 644, L49 (2006) J.I. González Hernández, R. Rebolo, G. Israelian, Astron. Astrophys. 478, 203 (2008) G. Israelian, R. Rebolo, G. Basri, J. Casares, E.L. Martín, Nature 401, 142 (1999) K. Jedamzik, Phys. Rev. D 74, 103509 (2006) K. Jedamzik, 710, arXiv:0710.5153 (2007) K. Jedamzik, 707, arXiv:0707.2070 (2007) K. Jedamzik et al., J. Cosmol. Astropart. Phys. 7, 7 (2006) M. Mayor et al., Messenger 114, 20 (2003) P. Molaro et al., Astron. Astrophys., in press, arXiv:0712.3345 (2007) M. Pospelov, 712, arXiv:0712.0647 (2007) T. Prodanovi´c, B.D. Fields, Astrophys. J. 645, L125 (2006) T. Prodanovi´c, B.D. Fields, Phys. Rev. D 76, 083003 (2007) N. Prantzos, Astron. Astrophys. 448, 665 (2006) K. Sadakane et al., Publ. Astron. Soc. Jpn. 58, 595 (2006) T.K. Suzuki, S. Inoue, Astrophys. J. 573, 168 (2002) V. Tatischeff, J.-P. Thibaud, Astron. Astrophys. 469, 265 (2007)

Part II

VLTI Science Highlights

VLTI Science Highlights Guy Perrin

Abstract The 2000 decade is remarkable for aperture synthesis in the optical and near-infrared wavelength regimes. Major interferometers such as VLTI have started science operation and have boosted the rate of astrophysics-oriented publications based on interferometric results. Also, the advent of long-baseline interferometers making use of large pupils has opened the way to faint science and first results on extragalactic objects has made it a reality. As the interest for the technique from the astronomical community keeps growing more ambitious instrumental projects are contemplated which will definitely put optical/near-infrared interferometry into mainstream astrophysics. It is shown here how VTLI plays a leading role in all these aspects.

1 Introduction: The Growing Impact of Interferometry in General and of VLTI in Particular Interferometry1 was mostly used for astronomical research by instrument builders until the beginning of the 2000 decade. As a consequence few papers were published by such a small community. With the building of the Keck and VLT interferometers, larger communities got access to interferometers and the number of astrophysical papers has started to grow after the first fringes in 2001 as is visible on Fig. 1. Other interferometers will benefit from the support of well-organized communities as is already the case for CHARA and hopefully for the MRO interferometer. The share of VLTI publications is of about 50% of the total amount. This number keeps increasing and 30 new published results are to be expected in 2008 from the European facility. One reason of the success is the wealth and diversity of instruments giving access to near up to mid-infrared wavelengths. The production of the three VLTI instruments (VINCI, MIDI, AMBER) is listed in Table 1. VINCI has not been offered since 2004. MIDI is in operation since 2004 and the first AMBER papers have been published in 2007. The counts were stopped at the end of September 2007. Most VLTI science targets are stars. However, the large pupils have allowed observations of active galactic nuclei. In either domains, breakthroughs have been possible thanks to VLTI. The list of stellar physics fields addressed with the VLTI is impressive: fundamental parameters, stellar surfaces, rotation, pulsation, dust, spots, disks, binaries, G. Perrin () Observatoire de Paris, LESIA, CNRS, UMR 8109, UPMC, Université Paris Diderot, 92190 Meudon, France e-mail: [email protected] 1 In the following, interferometry stands for optical/near-infrared interferometry. A. Moorwood (ed.), Science with the VLT in the ELT Era, Astrophysics and Space Science Proceedings, doi: 10.1007/978-1-4020-9190-2_14, © Springer Science + Business Media B.V. 2009

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Fig. 1 Number of interferometry-based science refereed publications over the last decade. The red curve is the annual total number of science papers. The blue curve is the annual number of papers for VLTI only. The dashed red curve is a fit to the annual number of papers. Source: http://olbin.jpl.nasa.gov/. A colour version of this figure is available at dx.doi.org/10.1007/978-1-4020-9190-2_14 Table 1 Number of science refereed publication per VLTI instrument and number of citations. Counts were stopped before October 2007 Instrument

Refereed science papers

Citations (before 10/2007)

VINCI

33

527

MIDI

36

321

AMBER

8

50

winds. . . . A few examples are presented in this paper. Others could have been chosen. This clear trend of the growing impact of interferometry (VLTI) in astrophysics will not stop and prospects are discussed to conclude this quick overview.

2 Stellar Physics The first of fundamental parameters to be derived from interferometric observations of a star is an apparent diameter or a typical size, whatever the source shape. This parameter can be obtained with a uniform disk model. The apparent diameter

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is a good approximation to the physical diameter for stars on the main sequence measured at near-infrared wavelengths where limb darkening is small. This simple measurement, combined with the bolometric flux, yields the effective temperature of a star or, equivalently in the case of a pure black body, its surface brightness. A surface brightness scale for dwarf and subgiant stars has been established with VINCI [1]. It allows to predict the diameter of these stars from their B–L index with an accuracy of a percent. In the same vein, given a parallax, the mass-radius relation of low-mass stars and brown dwarves was also obtained with VINCI [2] and showed good agreement with theoretical isochrones, an important tool to predict the characteristics of giant exoplanets (Fig. 2). In the special case of Procyon A for which asteroseismic, spectroscopic and interferometric data were available, an accurate mass could be determined (1.43 ± 0.03 M ) thus bringing a very strong constraint on its evolutionary path making this peculiar star a benchmark for evolutionary models [3]. Beyond fundamental parameters, surfaces and shapes of stars can be characterized. Even at infrared wavelengths, limb-darkening albeit small can be measured and yields clues on the structure of the atmosphere and of its temperature distribution [4]. Another degree of accuracy can be reached by studying the structures due to convection [5]. The most spectacular are rapidly rotating stars whose surface can be modeled with an elongated ellipsoïd as is the case for Achernar whose oblateness of 1.56 (Fig. 2) places the star near its stability limit [6]. The size of stars may also change with time in a spectacular way for Mira stars [7]. The more subtle diameter variations of cepheid stars allow to calibrate the distance scale and bridge the gap between cosmological distance scales and parallaxes [8].

Fig. 2 Left: the mass-radius relation for low-mass stars and brown dwarves [2]. Right: the oblateness of Achernar [6]. A colour version of this figure is available at dx.doi.org/10.1007/978-1-4020-9190-2_14

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3 Around Stars: From Gas and Dust to Planets The close environment of stars is very rich and the most internal regions require interferometry. One of the first AMBER results has given an accurate view of the intricate star, disk and wind of the young stellar object MWC 297 which had not been disentangled so far [9]. The spectrally dispersed visibilities measured in the K band show a resolved wind in Br γ and a more compact disk + star system in the continuum. This observation explains the skewness of the line in spectroscopy, the receding part of the wind being partly occulted by the disk. The study of the dusty and molecular environment of evolved stars is directly benefiting from the dispersed visibilities of MIDI between 8 and 13 µm. The shell of the C-rich mira star V Oph varies with phase [10]. Close to the star photosphere, the C2 H2 shell has a smallest radius and column density at minimum light. Further away, the optical depth of the amorphous carbon and SiC dust shell also reaches a minimum at minimum light. The coordinate use of VLTI and VLBA data of the O-rich mira star S Ori gives a full picture of the pulsating star and of its molecular and dusty environment up to the SiO maser ring [11]. The phase-dependent picture suggests that the mass loss increases near visual minimum and that the dust shell expands after maximum. MIDI measurements also shed light on the symbiotic star HM Sge [12]. The wind of a mira star is accreted by a hot white dwarf. A rather dense dusty shell around the mira has been found and its nature—a mix of amorphous carbon and silicate dust—has been understood. VLTI is also very useful to understand the process of planet formation. The study of protoplanetary disks by spectral energy distribution fitting is quite difficult as the uniqueness of the models is not guaranteed without spatial information. Here the ability to combine spectroscopy and interferometry is a key. The first spectacular result was the understanding of the changing nature of dust in protoplanetary disks as a function of distance to the central star [13]. Comparison of MIDI spectral visibilities to spectra disclosed the segregation of grains (Fig. 3). Grains in the central astronomical unit are larger and cristalline whereas grains in the outer regions of disks are smaller and amorphous. This result showing evidence of the processing of dust was interpreted as the signature of the first step towards planet formation. The disk around TW Hya was supposed to be the home of a massive planet to explain a 3 a.u. inner cavity. This hypothesis was ruled out thanks to MIDI data which instead showed a thick disk reaching down to 0.5 a.u.

4 Extreme Objects One of the strengths of VLTI is the access to spectral data from the J up to the N band and, more over, in a relatively short time. This is very important for very sudden and unexpected events as novae or supernovae. As a first example of such an astronomical event, RS Oph was observed with AMBER just 5.5 days after it exploded as a nova [14]. The complex structure of the source was revealed: an elongated central

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Fig. 3 The processing of dust in protoplanetary disks revealed by MIDI: cristalline dust in the inner a.u. and amorphous dust in the outer parts [13]. A colour version of this figure is available at dx.doi.org/10.1007/978-1-4020-9190-2_14

compact object detected in the continuum at K is embedded in relatively slow equatorial ejecta whose fast wind and shock front are identified thanks to observations in the Br γ and He I lines. Jets perpendicular to the equatorial plane have been detected at radio wavelengths two weeks later and images taken with HST five months later confirmed the bipolar shape detected earlier by AMBER and whose starting point has been established thanks to interferometry. Another example of extreme and complex object is the blue luminous variable star η Car. Studied in the past with HST and ground-based large telescopes, the mystery of the central source was still complete. It has been the subject of several studies with all three VLTI instruments [15–18]. VINCI measured the central source oblateness consistent with the large scale bipolar shape of the wind. MIDI data led to the identification of the central source enshrouded in a thick homonculus whose dusty structure was also better understood. The model of the bipolar wind produced by the fast rotating star has been well constrained by AMBER data.

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Fig. 4 The optical depth across the N band of the 2 pc scale structure around the core of NGC 1068 compared to silicate dust models [21]. A colour version of this figure is available at dx.doi.org/10.1007/978-1-4020-9190-2_14

5 Active Galactic Nuclei The direct investigation of the core of Active Galactic Nuclei is the most spectacular breakthrough of interferometry. These sources were too faint to be observable by small telescope interferometers. With the advent of the interferometric connection of 8-m class telescopes, AGNs have become a new field for interferometers. Although no images have been taken yet, the simple information on the size of the central region has allowed to challenge the unified theory of AGNs. NGC 1068, the archetype of Seyfert 2 galaxies, was the first source to be observed by VLTI. VINCI found an unresolved compact source possibly embedded into a more extended 2 pc component [19]. Observations at mid-infrared wavelengths with MIDI [20] found the source completely resolved and identified the 2 pc scale component with the putative torus of the unified scheme. The optical depth of the cool (226 K) extended component shows the signature of silicate dust (Fig. 4), one more evidence to support the unified scheme of AGNs [21]. Because of the complex structure of the object, short spatial frequencies are required to capture the object reality. VISIR observations allowed to connect the MIDI high angular resolution information with the source more extended regions, a necessity to better model the central core [22]. NGC 1068 is the first chapter of a long success story as eight more AGNs have been observed with MIDI (see W. Jaffe’s paper in this book).

6 Still Much More to Come. . . The first five years of science with VLTI have been very rich. As more astronomers use the instruments and as more sources become observable thanks to the technical

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progress achieved, VLTI science reaches new frontiers. AGN science was a first encouraging limit to be broken. The very next step towards still more exciting science is imaging which is of great help to better understand a priori unknown and complex sources, most of them becoming complex when observed with enough spatial resolution. The next generation of VLTI instruments will all make imaging a regular tool in interferometry. Fringe tracking to fight atmospheric turbulence at the interferometer scale is on the verge to dramatically increase the sensitivity of VLTI by many magnitudes. The frontier of faint object science will then be broken as well. A taste of what will become possible is already given by the first studies of the brighter objects around the galactic center [23]. The high accuracy astrometric capability of VLTI must be added to this. New mines of exoplanets are to be dug and black hole gazing will extraordinarily widen the scope of interferometry. The tens of papers published so far are just an appetizer of the whole meal of which only a few bites have been taken. The contribution of VLTI to astronomy will be overwhelming in ten years from now.

References 1. 2. 3. 4. 5. 6. 7. 8. 9. 10. 11. 12. 13. 14. 15. 16. 17. 18. 19. 20. 21. 22. 23.

P. Kervella et al., Astron. Astrophys. 426, 297 (2004) D. Segransan et al., Astron. Astrophys. 397, L5 (2003) P. Kervella et al., Astron. Astrophys. 413, 251 (2004) M. Wittkowski et al., Astron. Astrophys. 413, 711 (2004) J. Aufdenberg et al., Astrophys. J. 633, 424 (2005) A. Domiciano et al., Astron. Astrophys. 407, L47 (2003) H. Woodruff et al., Astron. Astrophys. 421, 703 (2004) P. Kervella et al., Astron. Astrophys. 416, 941 (2004) F. Malbet et al., Astron. Astrophys. 464, 43 (2007) K. Ohnaka et al., Astron. Astrophys. 466, 1099 (2007) M. Wittkowski et al., Astron. Astrophys. 470, 191 (2007) S. Sacuto et al., Astron. Astrophys. 465, 469 (2007) R. Van Boekel et al., Nature 432, 479 (2004) O. Chesneau et al., Astron. Astrophys. 464, 119 (2007) R. Van Boekel et al., Astron. Astrophys. 410, L37 (2003) O. Chesneau et al., Astron. Astrophys. 435, 1043 (2005) P. Kervella et al., Astron. Astrophys. 464, 1045 (2007) G. Weigelt et al., Astron. Astrophys. 464, 87 (2007) M. Wittkowski et al., Astron. Astrophys. 418, L39 (2004) W. Jaffe et al., Nature 429, 47 (2004) A. Poncelet et al., Astron. Astrophys. 450, 483 (2006) A. Poncelet et al., Astron. Astrophys. 472, 823 (2007) J.-U. Pott et al., Astron. Astrophys., in press

MIDI Sees Active Galactic Nuclei W. Jaffe, D. Raban, K. Meisenheimer, K. Tristram, Ch. Leinert and H. Röttgering

1 Introduction One of the primary motivations for building the VLTI was to observe the inner, dusty regions of Active Galactic Nuclei (AGNs). Although the VLTI is limited to relatively bright objects, there are a dozen or so Seyfert galaxies, and at least one quasar, that are bright enough to see. These are the only extragalactic targets so far observed with optical/infrared interferometers. The VLTI mid-Infrared instrument, MIDI, is a spectroscopic instrument working in the N-band (8–14 µ) which corresponds to the peak emission from blackbodies at ∼ 300 K. With a few exceptions, for example some cases of synchrotron emission, MIDI has mostly been used to observe warm dust at near this temperature. For AGNs the primary questions concerning dust are related to the unified theories that attempt to reduce the bewildering number of AGN types by attributing many of their characteristics to the obscuration of the nuclear regions by a dust structure, perhaps toroidal in form, with a characteristic size of a parsec or so. This scale size is determined by the typical UV luminosity of an AGN core, ∼ 1044 erg s−1 , and the need to keep the dust cooler than its sublimation temperature, ∼ 1500 K. For the nearest extragalactic AGNs at distances of ∼ 10 Mpc direct measurement of the dust morphology then requires a resolution of order 10 milliarcsec (mas). Single telescope observations of AGNs have shown mid-IR SEDs that prove the existence of warm dust there, but attempts to constrain the nature of the torus from these SEDs have proved difficult. Thus the primary motivation of the interferometric measurements was to image the dust structures and determine their shape, size, temperature and orientation with respect to the axes of the galaxy, the radio structures and the larger scale ionization structures. We also hoped to learn something about their chemistry from the shape of the silicate emission feature at 10 µ and possibly about their degree of inhomogeneity or “clumpiness”. The crudest taxonomic classification divides AGNs into “Type 1” where the torus is seen face-on, and “Type 2” where it is seen edge-on. In the first case we expect to see the AGN core as a bright unresolved point source surrounded by a near-circular cooler emission disk. In the second case we would only see the cooler disk and we would expect it to be flattened. W. Jaffe () Leiden Observatory, Nielsbohrweg 2, Leiden, Netherlands e-mail: [email protected] A. Moorwood (ed.), Science with the VLT in the ELT Era, Astrophysics and Space Science Proceedings, doi: 10.1007/978-1-4020-9190-2_15, © Springer Science + Business Media B.V. 2009

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2 Nearby Seyfert Galaxies The MIDI GTO observing list includes about a dozen objects, mostly Seyfert types 1 and 2 but including one radio galaxy and one quasar. They were chosen to have strong pointlike emission as seen by TIMMI 2. These are listed in Table 1. A single interferometric measurement does not produce a complete image but only a single component of the Fourier transform of an image at a single point in the “UV-plane”. If many such UV-points are observed, something like a map can be reconstructed but observing time constraints limited such complete mapping to two sources: NGC 1068 and the Circinus galaxy, both nearby Sy 2 galaxies. The results of these two observations are contained in the theses of David Raban [1] and Konrad Tristram [2]. A set of typical MIDI interferometric spectrum measured at different UV-positions is shown in Fig. 1. The analysis of these data show that in both galaxies there is indeed a flattened dust disk/torus of parsec size around the AGN core. In both cases we find a well defined hot central disk of size ∼ 1 pc with T ∼ 600 K and a larger, possibly less flattened, cooler extension with T ∼ 250 K out to tens of parsecs. We interpret the hot component as the central throat of the torus, directly exposed to radiation from the AGN core. The structure of the larger component is difficult to determine because we have not measured enough short UV points. This deficit we hope to correct with Auxiliary Telescope observations and VISIR data. Unexpectedly we found that the hot dust emission coincides in size and orientation with the positions of H2 O maser emission spots mapped by VLBI radio techniques [3]. Thus these steam masers come from exactly the regions where the warmest dust is found, apparently because there are warm molecules and the mid-IR pumping photons here. Radio continuum VLBI maps of NGC 1068 show that the dust/steam region surrounds a much hotter region characterized by free-free emission at X-ray temperatures. While the dust/maser disk in Circinus is perpendicular to the larger ionization cones in this galaxy, in agreement with the Table 1 AGN targets observed with MIDI Name

Flux (10 µ) Jy

NGC 1068

13

Sy 2

Yes

5

Sy 2

Yes

Centaurus A

0.6

RG

Yes

NGC 3783

0.5

Sy 1

Yes

Mrk 1239

0.4

Sy 1

Yes

MCG-05-23-016

0.3

Sy 2

Yes

NGC 7469

0.4

Sy 1

Yes

NGC 1365

0.6

Sy 2

Yes

IC 4329A

0.6

Sy 1

?

Circinus

Type

Detect?

NGC 253

?

Sy 2/SB

No

NGC 7582

?

Sy 2

No

QSO

Yes

3C 273

0.3

MIDI Sees Active Galactic Nuclei

91

Fig. 1 MIDI correlated flux spectra for NGC 1068, showing the broad silicate absorption feature, and the artefacts due to atmospheric O3 at 9.6 µ. A colour version of this figure is available at dx.doi.org/10.1007/978-1-4020-9190-2_15

simplest picture of the confinement of UV-light by the torus, this is not the case in NGC 1068, where the disk is tipped 45◦ from the radio/ionization cone axis. We are currently trying to understand how this misalignment was created, how long it can last, and what exactly the relation is between the disk axis and the UV-light propagation. Another, more subtle, indication of interesting physics in these disks is that the spectrum of the 9.7 µ silicate feature seen in absorption against the hot cores differs significantly from that of “standard” amorphous interstellar dust. These seems to indicate a dust composition rich in Aluminum, perhaps because such dust particles have higher sublimation temperatures. We also have indications that the cooler dust at larger distances is in fact inhomogeneous or clumpy. This was predicted from the SEDs in some models [4] and by hydrodynamic modeling of the motions of the dust and gas [5].

3 Other Targets Most of the other sources in Table 1 are Sy 1 or Sy 2 galaxies at larger distances than the two just discussed. For these we have typically only 1 or 2 UV-points per galaxy,

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and cannot say much about the size or morphology of the emission. Most seem to show a “cool” spectrum without obvious emission or absorption features. This is perhaps unexpected for the Sy 1 galaxies where much of the visible emission might come from the hot central core itself. The spectra in fact resemble the predictions of emission from an AGN surrounded by a disk which is everywhere optically thin to its own emission in the mid-IR [6] and hence the same in all directions. The two non-Seyfert targets are particularly interesting, although our UVcoverage here is also very limited. Centaurus A, a nearby radio galaxy, shows an correlated spectrum that is consistent with most of the emission arising from synchrotron processes at the radio core, modified by absorption by foreground cold dust [7]. However, the emission is in fact slightly resolved perpendicular to the radio jet on scales of 10 mas (∼ 0.2 pc), so apparently there is circumnuclear dust, warmed either by a weak hot central accretion disk or by the short wavelength synchrotron emission itself. The shape of the IR synchrotron spectrum allows us to estimate the magnetic field in the radio core as 26 µT and the energy loss/reacceleration timescale as ∼ 4 days, consistent with the variability of this source. The low mid-IR luminosity of the disk, relative to the large central mass and jet luminosity of Cen A, indicate that it is a very inefficient accretor, with Lbol ∼ 10−4 LEdd . The other non-Seyfert, 3C273, is a quasar at a redshift of 0.158, or about 650 Mpc, certainly the most distant object ever detected with optical/IR interferometers. From a single interferometric spectrum taken on 7 February 2007 we found a correlated flux of about 0.24 Jy, somewhat higher than a total flux measured three weeks later on VISIR. This apparent contradiction probably means that we caught this object during a short outburst. In this case it must arise from synchrotron emission in a very small radius (a few light days) and must be unresolved, even by the VLTI.

4 Conclusions 1. Most AGNs showing Jansky (or even sub-Jansky) level pointlike emission from single 4-m telescopes are detectable by the VLTI. 2. The best studied Sy 2 galaxies show dust disks that agree in size and orientation with the water masers, show hot central “throats”, signs of clumpiness, and perhaps Aluminum rich chemistry. 3. The disk of NGC 1068 is tipped relative to the direction expected from the ionization cones. 4. The more distant AGNs of both types show relatively flat spectra that resemble optically thin models. 5. The core of Cen A is dominated by synchrotron emission, but shows also an underluminous thermal disk.

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References 1. D. Raban, in preparation 2. K.R.W. Tristram, K. Meisenheimer, W. Jaffe, M. Schartmann, H.-W. Rix, Ch. Leinert, S. Morel et al., Astron. Astrophys. 474, 937 (2007) 3. J.F. Gallimore, S.A. Baum, C.P. O’Dean, Astrophys. J. 613, 794 (2004) 4. R. Mason, T.R. Geballe, C. Packham, N.A. Levenson, M. Elitzur, R.S. Fisher, E. Perlman, Astrophys. J. 640, 612 (2006) 5. M. Schartmann, K. Meisenheimer, M. Camenzind, H. Klahr, S. Wolf, Th. Henning, Astron. Nachr. 328, 671 (2007) 6. G.L. Granato, L. Danese, A. Franceschini, Astrophys. J. 486, 147 (1997) 7. K. Meisenheimer, K.R.W. Tristram, W. Jaffe, F. Israel, N. Neumayer, D. Raban, H. Röttgering, W.D. Cotton, Th. Henning, Ch. Leinert, B. Lopez, G. Perrin, A. Prieto, Astron. Astrophys. 471, 453 (2007)

The Use of the VLTI for Studying the Asymmetric Mass Loss of Evolved Stars Olivier Chesneau

The first generation interferometric recombiners of the VLTI have allowed to extend the range of targets that can be studied by optical interferometry. In the context of evolved stars, the VLTI is a breakthrough for its capabilities to hunt and study circumstellar disks, and also to observe transient events like Novae. Some synergy has been found with some VLT instruments like NACO, and also with the Plateau de Bure interferometer, despite an unfavorable location. These first results are promising for the second generation VLTI instruments, and an improved synergy with the new generation VLT instruments and ALMA is foreseen. The Paranal site would represent in this context a key facility for high spatial resolution studies of timevariable, spatially complex compact dusty sources.

1 Planetary Nebulae After having been extensively studied in the visible, the Planetary Nebulae (Pne), and among them the youngest and more dusty ones are now studied in the near and mid-IR. PNe exhibit amazing shapes, a large fraction of which show more or less marked deviations from the spherical symmetry, up to the most extreme bipolar ones. These asymmetries illustrate how complex the shaping mechanisms can be, but clear and conclusive observations supporting one or the other models proposed so far are missing and needed. High spatial resolution observations are promised to play a key in this field since most of the proposed shaping mechanism operate in the vicinity of the dying star, in the frontiers between its extended photospheres and its dusty environment. The ability of the VLTI to investigate the dust at high spatial resolution (1–10 mas) is well suited for the study of many aspects of the late evolution of stars and in particular, to deal with the asymmetry in PNe back to the AGB. O. Chesneau () Observatoire de la Côte d’Azur, Avenue Nicolas Copernic, 06130 Grasse, France e-mail: [email protected] A. Moorwood (ed.), Science with the VLT in the ELT Era, Astrophysics and Space Science Proceedings, doi: 10.1007/978-1-4020-9190-2_16, © Springer Science + Business Media B.V. 2009

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1.1 Complex Sources The discovery by De Marco et al. [7] of a dark lane in the STIS/HST spectra of the PN CPD-56◦ 8032, was interpreted as a disk and prompted us to investigate whether the MIDI/VLTI interferometer could help characterizing this compact dusty environment, and may be applied to other asymmetrical PNs. CPD-56◦ 8032 and Hen2113 were chosen for this first attempt. These targets, whose Wolf-Rayet central stars and nebulae share many characteristics, exhibit a core dominated by dust emission, and a complex dust chemistry witnessed by the simultaneous appearance of oxygenrich (silicate) emission and carbon-rich features in the form of strong PAHs emission. The chosen baseline, (UT2–UT3 ∼ 46 m), was in the lowest range allowed by the fixed UTs, but was already too long for these ‘extended’ environments. The torus in Hen2-113 (described in Lagadec et al. [8]) has a typical diameter of 0.6 and the torus/disk of CPD-56◦ 8032 appears more compact, 0.15 in diameter, but still very extended for such a long baseline. A global picture of CPD-56◦ 8032 emerged from the ISO (14 × 22 ) and MIDI (0.6 × 0.6 ) spectra, the HST images and the MIDI visibility curves. The outer nebula shows multiple lobes excavated in a dense environment. The ISO spectrum is dominated at long wavelength by the signature of crystalline silicates, while in the N band, the emission is compact and almost unresolved for a 8 m telescope, except in the PAHs bands. From the visibilities, a sharp, inclined ring (diameter of ∼ 200 AU) could be inferred. The inner rim is at about 500 K, but faces the 225 km s−1 wind from the central Wolf-Rayet star. This dusty structure may be caught is a stage of fast dissipation. OH 231.8 + 4.2 is a well studied pre-planetary nebula. The central source is a binary system formed by an M9-10III Mira variable and a hot companion. In the near to mid-IR, the core is deeply embedded into a highly optically thick dust envelope. This core is unresolved at 0.1 scale in K and L bands with NACO (see [11]). The first MIDI measurements, using baselines roughly perpendicular to the bipolar lobes were able to provide an extension of the source (assuming a Gaussian model) from 30 mas at 8 µm to 50 mas at 13 µm, i.e. 40–70 AU at 1.5 kpc. New measurements recently obtained in the direction of the bipolar axis show that the core cannot be ascribed by a spherical shell, although it is not sure that the term ‘disk’ is pertinent for a structure with such a large opening angle. The system is probably caught in a phase of intense mass-transfer, in which the dense dusty wind of the cool primary is highly perturbed by the presence of the hot companion orbiting inside. The geometry of the dust is probably very complex. CPD-56◦ 8032 and OH 231.8 + 4.2 are very interesting targets that would deserve a more extensive interferometric mapping. However, MIDI is a simple 2-beam recombiner, lacking closure phase information and such an effort should be undertaken with MATISSE, the second-generation 4-beam successor of MIDI [10].

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1.2 The Bipolar, Edge-on Sources The first targets observed with MIDI/VLTI were too complex to be investigated in depth at low cost via our partial interferometric observations. Moreover, the intermediate inclination of these targets affects their aspect in a manner difficult to infer. Our team decided to focus on compact systems whose inclination is close to 90◦ (edge-on configuration).1 We decided to focus on the famous bipolar nebulae Mz 3 (the Ant) and M2–M9 (the Butterfly) exhibiting the tighter waists, among the extreme cases of asymmetrical shaping. The Ant Nebula is one of the most striking planetary nebulae known. The morphology of the Ant Nebula is a bright core, three nested pairs of bipolar lobes and a ring-like outflow. The core is not unresolved by NACO from K to L bands. Mz 3 was observed using UT2–UT3 and UT3–UT4 baselines, and the observations could rapidly be interpreted in terms of a flat 10 × 20 mas structure. There was indeed no doubt that a disk was detected and this prompted us to use directly a radiative model to interpret these observations [3]. A simple passive disk could fit almost perfectly the data (one of the best match to date from a disk model and interferometric observations), giving confidence about our estimation of its mass content and its geometry. This model will be used to constrain the mass-loss history of the system, but information of the progenitor(s) is missing, hidden from your view with this high inclination. Some data have also been recorded on M2–M9, and the preliminary results are already promising. They show that the putative ‘disk’ inside the nebula is more extended than the one in Mz 3. This is in line with the larger mass (factor ∼ 4) inferred by Smith and Gehrz and with the SED suggesting dust at somewhat cooler average temperature than Mz 3. Moreover, the first visibilities reveal some spectral features indicative of crystalline silicates, generally associated to long-lived disks, i.e. a disk probably older than the bipolar structures seen in the HST images. Such features are much more difficult to infer in the ISO spectra. The analysis of these results is at the moment too early to draw firm conclusions, but the ‘disk’ around M2–M9 appears already as more massive and probably much older than the one of Mz 3, despite the fact that the kinetical ages of the nebulae are quite similar.

2 Novae It is tempting to associate the origin of the narrow waist of the extreme bipolar planetary nebulae by the influence of a hidden companion that perturb deeply mass-loss of the primary star, and many asymmetrical nebulae could witness the presence companions in symbiotic systems. In this context, it is not excluded that the Ant or the 1 It is interesting to note that other groups chose systems with pole-on configuration (i ∼ 0◦ ). Assuming circular symmetry, a few baselines can efficiently constrain the density law and structure of the disk (see [14]).

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Butterfly my be nascent cataclysmic systems, in which the conditions of a nova outburst are fulfilled. A classical nova eruption results from a thermonuclear runaway on the surface of a White Dwarf (WD) that is accreting material from a companion star in a close binary system. All aspects of the outbursts, from the temporal development of the fireball, followed by a dust formation phase or the appearance of the coronal phase can be studied with the VLTI, both in near and mid-IR. The detailed geometry of the first phases of novae in outburst remains virtually unexplored. The recent outburst from the recurrent2 nova RS Oph showed how complex such an ejection can be, as seen in the radio-interferometers and the HST images of the bipolar nebula formed rapidly [1, 12]. AMBER observations were secured only 5.5 days after discovery [5] providing an excellent K band dataset, including two bright emission lines, Brγ and HeI 2.06 µm. The triplet of baselines used was not sufficient for determining the geometry of the outburst, but could provide good constraints on the size of the K band continuum (dominated by the mostly optically thick emission from the nova), of the Brγ line forming region (located in the nova wind), and of HeI 2.06 µm extended forming regions (located close to the shock propagating inside the slow and dense wind of the red giant). An important information could also be extracted on the two components kinematics inside the Brγ line, an equatorially enhanced ‘slow’ (∼ 1800 km s−1 ) ejection, and the jet-like E–W emission at ‘high’ speed (∼ 3000 km s−1 ). Associated to the motions measured in the plane of the sky measured from the HST images taken 150 days later, a good picture of the true velocity field of the outburst emerges. The RS Oph observations were based on a single snapshot, followed by a few AMBER measurements 60 days later, but the source could not be really be monitored. Recently, a great opportunity to observe one of the brightest classical novae the classical was not missed by the VLTI. Soon after its discovery, the slow nova V1280 Sco showed an intense event of dust formation, faded abruptly in visual, and became very bright in the near and mid-IR. This nova was monitored during 4 months (two DDT accepted 278.D-5053, 279.D5014), providing the first spatially resolved observations of a dust forming nova. In absence of clear evidence of asymmetry, each stage is currently modeled using the spherical radiative transfer code dusty, providing of wealth of information on the physical conditions that leaded to such an efficient, and long duration dust formation event.

3 Conclusion The studies presented here are illustrative of the numerous applications of the VLTI for the study of evolved stars. Many articles, referring more on the evolved stars photosphere and their winds could also be cited, such as the recent extensive study 2 Contrary to classical novae for which the recurrence time can be of the order of 10 000 yrs or more, the recurrent novae are characterized by several outburst registered in an historical period of time. The recurrence timescale of RS Oph is one of the shortest, of the order of 20 yrs.

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of the mass-loss of some Mira stars in connection with their stellar pulsations: S Ori by Wittkowski et al. [16] and V Oph by Ohnaka et al. [13]. Another aspect not treated here, is the detection, in the dust forming site, of big dust clumps seen in the most extreme envelopes, such as the ones reported in R CrBs and recently observed around RY Sgr by MIDI (Leão et al. [9]). Our language simplifications are often not reflecting the spatial complexity of the dusty sources observed. Out of the targets cited here, only Mz 3 surely harbors what is usually denoted as a ‘flat passive disk’. The flared, dissipating structures seen around the Wolf-Rayet stars in CPD-568032 and Hen2–113 are better ascribed with a dusty torus than with a stratified passive disk. The situations gets more complicated with OH 231.8 and dusty symbiotic systems like HM Sge (Sacuto et al. [15]) for which a dense, spherical Mira wind is perturbed by a hot WD. Spherical and disk simple geometries are no longer pertinent in that context, and imaging capabilities are necessary to clarify the complex interaction between the stars and their environment. The VLTI is already in the Phase B study of the second generation recombiners that should be able to recombine the 4 ATs/4 UTS simultaneously in the mid-IR (MATISSE, Lopez et al.). Also very promising is the convergence in terms of spatial resolution between the optical and millimetric interferometry techniques. The Plateau de Bure’s interferometer new long baselines provide a spatial resolution in the millimetric domain comparable with the resolution of a single-dish UT at 10 µm (∼ 300 mas) and ALMA should extend this complementarity to the scale of the dusty disks currently discovered with the VLTI (i.e. ∼ 10 mas). High dynamics images from the ELT could also provide a unique view of the disk/nebula interface, allowing to precise the conditions leading to the launch and channeling of material out the core of the system.

References 1. 2. 3. 4. 5. 6. 7. 8. 9. 10. 11. 12. 13. 14. 15. 16.

M.F. Bode, D.J. Harman, T.J. O’Brien et al., Astrophys. J. Lett. 665, L63 (2007) J.-F. Desmurs, J. Alcolea, V. Bujarrabal et al., Astron. Astrophys. 468, 189 (2007) O. Chesneau, F. Lykou, B. Balick et al., Astron. Astrophys. Lett. (2007), accepted O. Chesneau, N. Nardetto, F. Millour et al. O. Chesneau, A. Collioud, O. De Marco et al., Astron. Astrophys. 455, 1009 (2006) O. Chesneau, T. Verhoelst, B. Lopez et al., Astron. Astrophys. 435, 563 (2005) O. De Marco, M.J. Barlow, M. Cohen, Astrophys. J. Lett. 574, L83 (2002) E. Lagadec, O. Chesneau, M. Matsuura et al., Astron. Astrophys. 448, 203 (2006) I.C. Leão, P. de Laverny, O. Chesneau et al., Astron. Astrophys. 466, L1 (2007) B. Lopez et al., in Science with the VLT in the ELT Era, ed. by Alan Moorwood. Astrophysics and Space Science Proceedings (Springer, Netherlands, 2009), p. 353 M. Matsuura, O. Chesneau, A.A. Zijlstra et al., Astrophys. J. Lett. 646, L123 (2006) T.J. O’Brien, M.F. Bode, M.W. Porcas et al., Nature 442, 279 (2006) K. Ohnaka, T. Driebe, G. Weigelt, M. Wittkowski, Astron. Astrophys. 466, 1099 (2007) Th. Ratzka et al., in Science with the VLT in the ELT Era, ed. by Alan Moorwood. Astrophysics and Space Science Proceedings (Springer, Netherlands, 2009), p. 101 S. Sacuto, O. Chesneau, M. Vannier, P. Cruzalèbes, Astron. Astrophys. 465, 469 (2007) M. Wittkowski, D.A. Boboltz, K. Ohnaka et al., Astron. Astrophys. 470, 191 (2007)

Mid-infrared Interferometric Observations of Young Circumstellar Discs Th. Ratzka, Ch. Leinert, R. van Boekel and A.A. Schegerer

MIDI was designed to open the mid-infrared wavelength regime to the VLTI. It allows to resolve and study the warm regions of dusty discs around young stellar objects in great detail due to the perfect match between the spatial resolution provided by the VLTI and the size of the regions dominating in the N-band. The recent years thus consequently lead to new insights into the circumstellar environments of a wide variety of pre-main-sequence stars. With the aim to illustrate the capabilities of MIDI and the VLTI in this field, we will give typical examples of what has been done since the first investigation of the mid-infrared sizes of discs around Herbig Ae/Be stars.

1 TW Hya With its distance of about 50 pc and its estimated age of about 10 Myr TW Hya is a key object for investigating the transition phase from discs to planets. A detailed analysis of the spectral energy distribution (SED) led [1] to conclude that the inner disc region within 4 AU of the central star is almost devoid of dust, with only an optically thin layer of small (sub)micron-sized dust grains remaining. With Keck interferometry measurements, [2] could spatially resolve the dust continuum emission at a wavelength of 2 µm. They concluded that the inner disc consists of optically thin, submicron-sized dust extending to within 0.06 AU where it may be magnetospherically truncated. In May 2005 we were able to interferometrically measure this comparatively faint young stellar object in the mid-infrared (MIR) [3]. The actual projected baseline length of about 50 m provided a spatial resolution or, more precisely, fringe spacing of about 40 mas at 10 µm, i.e. 2 AU at the distance of TW Hya. The TW Hya disc was simulated by using a model for passive irradiated flared discs [4]. This model describes the disc with two layers: a disc interior layer and a disc surface layer. The flared shape of the disc ensures that the stellar light can irradiate and heat the surface layer to the typical temperature of optically thin dust. Th. Ratzka () Astrophysical Institute of Potsdam, An der Sternwarte 16, 14482 Potsdam, Germany e-mail: [email protected] A. Moorwood (ed.), Science with the VLT in the ELT Era, Astrophysics and Space Science Proceedings, doi: 10.1007/978-1-4020-9190-2_17, © Springer Science + Business Media B.V. 2009

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Fig. 1 Left: Comparison of the observed SED (triangles) and the model fit (Mstar = 0.6 M , Lstar = 0.19 L , Tstar = 4000 K, Rdisc = 140 AU). The total SED (solid thick) consists of the contribution from the star (dashed–dotted), the optically thin inner disc (solid thin), the midplane of the optically thick outer disc (dashed), and its atmosphere (dotted). The additional component in the NIR is the contribution due to scattering. Right: Comparison of observed visibility and the best simultaneous fit (solid) to both SED and visibility. The visibility found by reproducing the model described by [1] is also shown (dashed–dotted). The grey regions (9.3–9.9 µm, > 12.9 µm) are affected by the atmosphere

However, similar to the model of [1] we included a sort of parametrised ‘roundedoff’ inner rim. This was done by reducing the pressure scale height over a certain radius range. Following the result of [2] also an optically thin distribution of dust inside the rim has been added. Simultaneous fits of both the SED and the visibility (see Fig. 1) confirmed that the circumstellar disc of TW Hya can be described best by a rather standard outer part with an abrupt transition to an inner part with reduced emission. However, based on our visibility measurements, we put this transition region much closer to the star than it has been done in earlier models, namely between 0.5 and 0.8 AU. The model presented by [1] fails to reproduce both the shape and the low values of the measured visibility, although it includes the required ‘cliff-like’ rim at an even larger distance from the central star than in our model. The reason for this finding becomes apparent when analysing the contributions of the various disc components to the SED. In the model developed by [1] the ‘wall’ is so far away from the central heating source that it cannot contribute significantly to the N-band flux at all. In their model the mid-infrared visibility traces only the optically thin inner disc. Assuming that the inner disc, as already discussed by [2], does not extend inwards fully to the dust sublimation and corotation limit, we could also reproduce the Keck near-infrared (NIR) results.

2 RY Tau RY Tau is a well-known T Tauri star that belongs to the Taurus-Auriga association at a distance of about 140 pc and has shown irregular photometric variability in the

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visible and NIR wavelength range. RY Tau was observed in November 2004 with projected baselines of 48 m and 78 m [5]. In contrast to the above described analytical study we used here the code MC3D which is based on the Monte Carlo method [6, 7]. The vertical density distribution in the disc is calculated self-consistently assuming hydrostatic equilibrium, i.e. the balance of gravitational and thermal pressure [5]. Furthermore, we assume a two-layer dust model, i.e. not-evolved, interstellar dust of astronomical silicate and carbon in the surface regions while in the disc interior dust grains with a maximum size of 1 mm are additionally considered. The latter approach preliminarily simulates dust settling in the disc. An active disc model is our favourite approach consisting of a passively heated (dust) disc where accretion effects are added. Moreover, an circumstellar envelope can be considered, too. The best fit parameters for our active model without an envelope are given in the caption of Fig. 2. In our study it is shown that accretion effects as well as an optically thin envelope can generate additional NIR and MIR radiation and have similar effects on the SED and visibilities. Only complementary observations in the UV range where the accretion rate can be independently measured, will provide additional constraints to disentangle accretion effects and envelope contributions. We have to mention that the measurements (SED and MIR visibilities) could be even reproduced without considering any accretion effects. Since in our model the vertical density distribution is calculated assuming hydrostatic equilibrium, the potential formation of a ‘puffed-up’ inner wall, which also affects the NIR and MIR range of the SED and visibilities, is included in a natural way. However, we do not see an excessively puffed-up inner wall with our approach. We find by fitting the correlated spectra for RY Tau a decreasing contribution of non-evolved, i.e. small (0.1 µm) amorphous, dust grains and an increasing crystallinity with increasing baseline length, i.e. decreasing contribution from the colder circumstellar material (Fig. 2). The crystalline material is thus concentrated mainly in the inner parts of the disc (point C). The abundance decreases strongly with decreasing resolution (point B) and converges for the single-dish observation (point A) to a lower limit that approximately corresponds to the abundance of crystalline dust in interstellar matter [8]. The result for RY Tau shows that the formation of crystalline dust grains is not only favoured in the innermost disc region of Herbig Ae/Be stars [9], but also in T Tauri stars.

3 UX Ori A subset of the intermediate mass Herbig Ae stars show irregular variability at optical wavelengths, thought to be due to changing column densities of intervening dust causing variable extinction [10]. These stars are called UXOR variables, after their prototype UX Ori. While continuously variable on a level of a few tenths of mag (e.g. [11–13]), they occasionally undergo much deeper ‘Algol-type’ minima. During such events, which occur a few times per year on average and last for several days, the stars dim by 2−3 magnitudes. The decreasing brightness is accompanied

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Fig. 2 Model of the SED (upper left) and the visibilities (right column) of RY Tau. We consider a passively heated disc (Rin = 0.3 AU, Rout = 270 AU, Mdisc = 4 · 10−3 M , Mstar = 1.69 M , Lstar = 10.0 L , Tstar = 5560 K) and accretion with a rate of 9.1·10−8 M yr−1 . For the inclination only an upper limit of 70◦ could be derived. The dashed line in the first panel represents the intrinsic, stellar flux. The visibilities were calculated from the corresponding model images for 8.5 µm, 9.5 µm, 10.6 µm, 11.5 µm, and 12.5 µm. Triangles and squares represent the upper and lower limit of the visibilities for different position angles but the same inclination of the model. Lower left: Relative mass contribution (RMC) of crystalline (squares) and 0.1 µm-sized amorphous (triangles) silicate grains plotted versus the reached spatial resolution of our observations

by a strong increase in the degree of linear polarisation. In order to explain the photometric and polarimetric behaviour during brightness minima, it has been proposed that UXOR systems harbour discs that are seen nearly edge-on. While UX Ori has been studied extensively at mm wavelengths using both single-dish telescopes and interferometers, its disc has not yet been spatially resolved, though a stringent upper limit on the disc size of 100 AU could be set [14]. The spectral slope at mm wavelengths was found to be very shallow: Fν ∝ ν α , α = 2.1 ± 0.2 [14]. Two disc configurations were proposed to explain this shallow slope: (1) a compact flared disc (∼ 30 AU) which is optically thick all the way to the mm regime and contains only small dust grains; and (2) a larger disc containing an additional reservoir of cm-sized particles (‘pebbles’) in the disc midplane. Appropriate selection of telescope combinations and timing of the observations with MIDI in December 2004 and October 2007 resulted in nearly orthogonal projected baselines, roughly along the anticipated major and minor disc axes. The interferometric visibilities on both baselines are significantly below unity, i.e. we spatially resolve the disc’s thermal infrared emission. While the baselines are of similar

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length (∼ 4 : 3 ratio) the visibility levels differ strongly, indicating the emission strongly deviates from circular symmetry and providing supportive evidence for a highly inclined disc. We have simultaneously modelled our visibility data and the SED of UX Ori, using a 2D radiative transfer model of a circumstellar disc whose density structure is determined by two parameters: the scale height of the disc at the inner edge H0 and the ‘flaring index’ γ (the ratio of the disc scale height and distance to the star is given by H /R = (H0 /Rin )(R/Rin )γ ). Two additional fit parameters are the disc inclination and the position angle of the major axis. We explore both disc configurations that were proposed to explain the mm observations, and find that we can reproduce the 10 µm visibilities and the SED with both approaches. The strongest constraint that our measurements provide is the high disc inclination of 70◦ to 75◦ . This agrees well with existing ideas that were based on spatially unresolved observations. With only two visibility measurements the position angle of the minor axis is not strongly constrained. The range of values found (133◦ –142◦ E of N) is in reasonable agreement with the value of 127◦ derived from polarimetric measurements taken during a deep photometric minimum in the mid 80s [11].

References 1. 2. 3. 4. 5. 6. 7. 8. 9. 10. 11. 12. 13. 14.

N. Calvet, P. D’Alessio, L. Hartmann et al., Astrophys. J. 568, 1008 (2002) J.A. Eisner, E.I. Chiang, L.A. Hillenbrand, Astrophys. J. 637, L133 (2006) T. Ratzka, C. Leinert, T. Henning et al., Astron. Astrophys. 471, 173 (2007) C.P. Dullemond, C. Dominik, A. Natta, Astrophys. J. 560, 957 (2001) A. Schegerer, S. Wolf, T. Ratzka et al., Astron. Astrophys., in press S. Wolf, T. Henning, B. Stecklum, Astron. Astrophys. 349, 839 (1999) I. Pascucci, S. Wolf, J. Steinacker et al., Astron. Astrophys. 417, 793 (2004) H.P. Gail, in: Astromineralogy, ed. by T.K. Henning. LNP, vol. 609 (Springer, Berlin, 2003) pp. 55–120 R. van Boekel, M. Min, C. Leinert et al., Nature 432, 479 (2004) W. Wenzel, in: Non-Periodic Phenomena in Variable Stars, ed. by L. Detre. IAU Colloq. 4 on Variable Stars (Budapest, Academic Press, 1969) pp. 61–73 N.V. Voshchinnikov, V.P. Grinin, N.N. Kiselev et al., Astrophysics 28, 182 (1988) V.P. Grinin, P.S. The, D. de Winter et al., Astron. Astrophys. 292, 165 (1994) V.P. Grinin, O.V. Kozlova, A. Natta et al., Astron. Astrophys. 379, 482 (2001) A. Natta, T. Prusti, R. Neri et al., Astron. Astrophys. 350, 541 (1999)

VLTI-AMBER Observations of η Carinae with High Spatial Resolution and Spectral Resolutions of λ/λ = 1500 and 12 000∗ G. Weigelt, S. Kraus, T. Driebe, K.-H. Hofmann, F. Millour, R. Petrov, D. Schertl, O. Chesneau, K. Davidson, A. Domiciano de Souza, T. Gull, J.D. Hillier, F. Malbet, F. Rantakyrö, A. Richichi, M. Schöller and M. Wittkowski

We report spectro-interferometric observations of the Luminous Blue Variable (LBV) η Car using the AMBER instrument of ESO’s Very Large Telescope Interferometer (VLTI) [1]. The observations around the He I 2.059 µm and the Brγ 2.166 µm emission lines allow us to investigate the wavelength dependence of the visibility, differential phase, and closure phase of η Car’s stellar wind region. If we fit visibility profiles derived from Hillier et al. models [2] to the AMBER visibilities, we obtain 50% encircled-energy diameters of 4.3, 6.5, and 9.6 mas in the 2.17 µm continuum, the He I, and the Brγ emission lines, respectively. For the interpretation of the non-zero differential and closure phases measured within the Brγ line, we developed a simple wind model. Our observations support theoretical models of winds from fast-rotating, luminous hot stars.

1 Introduction η Car is one of the most luminous and most massive (M ∼ 100 M ) LBVs, exhibiting an extremely high mass loss rate [3]. Spectroscopic studies of its surrounding Homunculus nebula showed that η Car’s wind is latitude-dependent [4]. η Car’s aspherical wind can be explained by models of fast rotating, luminous hot stars [5, 6], which predict an increased wind speed and density along the polar axis. Van Boekel et al. [7] resolved η Car’s wind using the VLTI/VINCI instrument. Hillier et al. [2, 8] developed a radiative transfer model which is able to predict the wavelength dependence of η Car’s wind zone. There is indirect evidence, such as η Car’s 5.5-yr spectroscopic periodicity, which suggests that η Car is a binary [9]. G. Weigelt () Max-Planck-Institut für Radioastronomie, Auf dem Hügel 69, 53121 Bonn, Germany e-mail: [email protected] A. Moorwood (ed.), Science with the VLT in the ELT Era, Astrophysics and Space Science Proceedings, doi: 10.1007/978-1-4020-9190-2_18, © Springer Science + Business Media B.V. 2009

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2 Observations The AMBER instrument of ESO’s Very Large Telescope Interferometer is a nearinfrared beam-combiner [10–12] which is able to measure visibilities, differential phases, and closure phases [10, 13] with spectral resolutions of 35, 1500, or 12 000. Our observations were performed on 2004 December 26, 2005 February 25, and 2005 February 26 with three VLTI 8.2 m telescopes [1]. The data were reduced with the amdlib software package (http://amber.obs.ujf-grenoble.fr; see [14]).

3 Results 3.1 Visibilities, Differential Phases, and Closure Phases Figure 1 shows medium-resolution spectra, visibilities, differential phases, and closure phases extracted from the AMBER interferograms. From these measurements and the comparison of the observations with various types of models, we obtained the results discussed in the following sections (see [1] for more details).

3.2 Comparison of the Observed Visibilities with Model Predictions Figure 2 shows a comparison of our AMBER visibilities with the Hillier et al. 2001 model visibilities [2]. Most of the Hillier model predictions are approximately in agreement with our observations. However, in the He I line we find significant deviations between the measured wavelength-dependent visibilities and the prediction of the 1-D Hillier model, possibly indicating that this line contains major contributions from the wind-wind collision region of η Car’s primary star and its hypothetical companion [9, 15, 16].

3.3 Diameter of η Car’s Wind Zone in the Continuum, He I, and Brγ Emission Line The comparison of the AMBER continuum visibilities with the Hillier model predictions [2] shows that there is a good agreement. Fitting the Hillier center-to-limb intensity profiles to the observations, we obtain a 50% encircled-energy diameter of 4.3 mas for the continuum-emitting region around 2.17 µm. Furthermore, we found an elongation towards position angle 120 ± 15◦ with a projected axis ratio of 1.18 ± 0.1, in agreement with previous VLTI/VINCI observations [7]. These results

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Fig. 1 VLTI-AMBER observations of η Car around the Brγ (left) and He I line (right; medium spectral resolution R = λ/λ = 1500): (top) spectra and visibilities for the three baselines described in the figures, (middle) wavelength-differential phases, and (bottom) closure phase. From [1]. A colour version of this figure is available at dx.doi.org/10.1007/978-1-4020-9190-2_18

support theoretical predictions of an enhanced mass loss in polar direction [5, 6]. Fits of Hillier model visibilities to the observed emission line visibilities yield 50% encircled-energy diameters of 6.5 and 9.6 mas in the He I and the Brγ emission lines, respectively.

3.4 Differential Phase and Closure Phase Within the Brγ and He I emission lines we measured non-zero differential and closure phases (see Fig. 1) which suggest a complex, asymmetric object structure. We developed an aspherical, latitude-dependent stellar wind model which is able to explain the observations (see [1]).

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Fig. 2 Comparison of the AMBER η Car spectra and visibilities with the predictions of the Hillier et al. model [2]: spectra (upper row) and visibilities (bottom three rows) of the AMBER measurements (red) and model predictions [2] (green). The data from 2004 (middle and right) was taken with medium spectral resolution (R = 1500), whereas the 2005 Brγ measurement (left) was obtained with high spectral resolution (R = 12 000). The differences in the He I line can possibly be explained by the emission from a wind collision zone [8, 15]. From [1]. A colour version of this figure is available at dx.doi.org/10.1007/978-1-4020-9190-2_18

4 Conclusion • The comparison of our AMBER Brγ and continuum visibilities with the radiative transfer model from Hillier et al. [2] shows that there is a good agreement between the model predictions and the observations. The differences between the model and observations in the He I line can possibly be explained by the emission from a binary wind-wind collision zone. • From fits of Hillier model visibility profiles [2] to the AMBER visibilities, we derived 50% encircled-energy diameters of 4.2, 6.5, and 9.6 mas in the 2.17 µm continuum, the He I, and the Brγ emission lines, respectively. • In the continuum around the Brγ line, we found an asymmetry towards position angle PA = 120 ± 15◦ with a projected axis ratio of 1.18 ± 0.10. This result supports theoretical studies which predict an enhanced mass loss in polar direction for massive stars rotating close to their critical rotation rate (e.g., [5, 6]). • The measured non-zero differential phases and closure phases within the Brγ and the He I emission lines suggest an asymmetric, wavelength-dependent ob-

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ject structure. We developed an aspherical, latitude-dependent stellar wind model which can explain all available observations. Acknowledgements Based on observations collected at the European Southern Observatory, Paranal, Chile, within the AMBER programmes 074.A-9025 and 074.A-9024. This research has made use of the SIMBAD database at CDS, Strasbourg (France), and the Smithsonian/NASA Astrophysics Data System (ADS). The project has also benefited from funding from the French Centre National de la Recherche Scientifique (CNRS) through the Institut National des Sciences de l’Univers (INSU) and its Programmes Nationaux (ASHRA, PNPS).

References 1. 2. 3. 4. 5. 6. 7. 8. 9. 10. 11.

12. 13. 14. 15. 16.

G. Weigelt, S. Kraus, T. Driebe et al., Astron. Astrophys. 464, 87 (2007) D.J. Hillier, K. Davidson, K. Ishibashi, T. Gull, Astrophys. J. 553, 837 (2001) K. Davidson, R.M. Humphreys, Annu. Rev. Astron. Astrophys. 35, 1 (1997) N. Smith, K. Davidson, T.R. Gull et al., Astrophys. J. 586, 432 (2003) S.P. Owocki, S.R. Cranmer, K.G. Gayley, Astrophys. J. 472, L115 (1996) S.P. Owocki, S.R. Cranmer, K.G. Gayley, Astrophys. Space Sci. 260, 149 (1998) R. van Boekel, P. Kervella, M. Schöller et al., Astron. Astrophys. 410, L37 (2003) D.J. Hillier, T. Gull, K. Nielsen et al., Astrophys. J. 642, 1098 (2006) A. Damineli, Astrophys. J. 460, L49 (1996) R.G. Petrov, F. Malbet, G. Weigelt et al., SPIE 4838, 924 (2003) R.G. Petrov, F. Millour, O. Chesneau et al., in: The Power of Optical/IR Interferometry Recent Scientific Results and 2nd Generation VLTI Instruments, ed. by A. Richichi, F. Delplancke, F. Paresce, A. Chelli (Springer, Berlin, 1973) pp. 143–153 R.G. Petrov, F. Malbet, G. Weigelt et al., Astron. Astrophys. 464, 1 (2007) F. Millour, M. Vannier, R.G. Petrov et al., EAS Publ. Series 22, 379 (2006) E. Tatulli, F. Millour, A. Chelli et al., Astron. Astrophys. 464, 29 (2007) K.E. Nielsen, M.F.T. Gull et al., Astrophys. J. 660, 669 (2007) J.M. Pittard, M.F. Corcoran, Astron. Astrophys. 383, 636 (2002)

Resolving the Inner Active Accretion Disk Around the Herbig Be Star MWC 147 with VLTI/MIDI + AMBER Spectro-interferometry S. Kraus, Th. Preibisch and K. Ohnaka

Abstract We studied the geometry of the inner (AU-scale) circumstellar environment around the Herbig Be star MWC 147. Combining, for the first time, near(NIR, K band) and mid-infrared (MIR, N band) spectro-interferometry on a Herbig star, our VLTI/MIDI and AMBER data constrain not only the geometry of the brightness distribution, but also the radial temperature distribution in the disk. For our detailed modeling of the interferometric data and the spectral energy distribution (SED), we employ 2-D radiation transfer simulations, showing that passive irradiated Keplerian dust disks can easily fit the SED, but predict much lower visibilities than observed. Models of a Keplerian disk with emission from an optically thick inner gaseous accretion disk (inside the dust sublimation zone), however, yield a good fit of the SED and simultaneously reproduce the observed NIR and MIR visibilities. We conclude that the NIR continuum emission from MWC 147 is dominated by accretion luminosity emerging from an optically thick inner gaseous disk, while the MIR emission also contains strong contributions from the outer dust disk.

1 Introduction and Observations Herbig Ae/Be (HAeBe) stars are intermediate-mass objects, which are still accreting material, probably via a circumstellar disk composed of gas and dust. Understanding the structure of these disks and the processes through which they interact with the central star is critical for our understanding of the formation process of stars. Since, until recently, the spatial scales of the inner circumstellar environment (a few AU) were not accessible to infrared imaging observations, conclusions drawn on the 3-D geometry of the circumstellar material were mostly based on the modeling of the SED (e.g. Hillenbrand et al. [1]). However, these fits are known to be ambiguous (e.g. Men’shchikov and Henning [2]) and have to be complemented with spatial information, as provided by infrared interferometry. The first systematic studies of HAeBes using the technique of infrared longbaseline interferometry have revealed that for most HAeBes the NIR size correlates with the stellar luminosity L following a simple R ∝ L1/2 law, suggesting that the NIR continuum emission mainly traces hot dust at the inner sublimation radius. S. Kraus () Max-Planck-Institut für Radioastronomie, Auf dem Hügel 69, 53121 Bonn, Germany e-mail: [email protected] A. Moorwood (ed.), Science with the VLT in the ELT Era, Astrophysics and Space Science Proceedings, doi: 10.1007/978-1-4020-9190-2_19, © Springer Science + Business Media B.V. 2009

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However, for the more luminous Herbig Be stars, the NIR-emitting structure is more compact than predicted by the size-luminosity relation (Monnier et al. [3]). To further investigate the origin of the “undersized” Herbig Be star disks, we observed the Herbig Be star MWC 147 with the VLTI. First infrared interferometric observations of this star by Akeson et al. [4] have indicated that the NIR-emitting region is surprisingly compact (∼ 0.7 AU, assuming a uniform disk profile). In the course of three ESO open time programs and using the 8.2 m UT telescopes, we have obtained seven VLTI/MIDI measurements and one VLTI/AMBER measurement. The MIDI observations cover baseline lengths between 39 and 102 m. From the AMBER data, one wavelength-dependent visibility, corresponding to a 101 m baseline, could be extracted. For our modeling of MWC 147, we adopt the stellar parameters by Hernández et al. [5], namely a spectral type of B6, a distance of 800 pc, a bolometric luminosity of 1550 L , a mass of 6.6 M , and a stellar radius of 6.63 R .

2 The Power of Joint NIR/MIR Spectro-interferometry The VLTI instruments AMBER and MIDI combine the high spatial resolution achievable with infrared interferometry with spectroscopic capabilities, measuring the fringe visibility as a function of wavelength. As circumstellar disks exhibit a temperature gradient, different spectral channels trace different spatial regions. Therefore, spectro-interferometric observations, which cover a sufficiently large wavelength range, can constrain not only the disk geometry, but also the radial temperature profile of the disk. For the investigation on YSO disks, the NIR and MIR wavelength regimes are particularly well suited since the NIR wavelength regime (∼ 2 µm) is most sensitive to the thermal emission from dust located at the dust sublimation radius (T ≈ 1500 K, a few AU from the star), while MIR wavelengths (∼ 10 µm) trace dust with a temperature of several hundred Kelvin, located a few 10 AU from the star (see Fig. 1).

Fig. 1 Illustration of the inner environment of HAeBe stars. Due to the different temperatures, spectro-interferometry can disentangle multiple emission components. A colour version of this figure is available at dx.doi.org/10.1007/978-1-4020-9190-2_19

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3 Radiative Transfer Modeling As a first step of analysis we compared the interferometric data to commonly used analytic disk models with a simple temperature power-law (T (r) ∝ r −α , with α = 3/4 or 1/2). Using the assumption that each disk annulus radiates as a blackbody, we can compute the wavelength-dependence of the disk size corresponding to these analytic models and find that they cannot reproduce the measured NIR and MIRsizes simultaneously (see Fig. 2). Therefore, we applied a more sophisticated modeling approach using the mcsim_mpi 2-D radiative transfer code (Ohnaka et al. [6]). For each model, we first check the agreement with the SED of MWC 147 and then fit the spectrointerferometric visibilities (see Kraus et al. [7] for details). The dust density distribution of the accretion disk in our models resembles a flared, Keplerian-rotating disk with a radial density distribution ρ(r) ∝ r −3/2 , which extends from the dust sublimation radius to 100 AU. In order to reproduce the shape of the SED, we found that, in addition to the disk, an extended envelope is required, for which we use the radial density distribution ρ(r) ∝ r −1/2 . Figure 3 shows model images, the SED, as well as the visibilities corresponding to our best-fit model of a passive, irradiated accretion disk. Although irradiated disks are able to reproduce the SED of MWC 147, they predict visibilities which are much smaller than the measured visibilities (χr2 ≈ 26) and are therefore in strong conflict with our interferometric measurements.

Fig. 2 Comparing the wavelength-dependent disk size of MWC 147 (as derived from one of our VLTI/AMBER + MIDI measurements) with the wavelength-dependent disk size predicted by standard temperature power-law disk models, we find that these models cannot reproduce our measurements. In this figure, the disk model was scaled to match the measured MIR size. A colour version of this figure is available at dx.doi.org/10.1007/978-1-4020-9190-2_19

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Fig. 3 Model images (a–d), SED (e), and NIR/MIR visibilities f–g corresponding to our best-fit radiative transfer image of an irradiated dust disk geometry. The poor agreement between measured and model visibilities (χr2 ≈ 26) indicates that passive disk models can be ruled out. A colour version of this figure is available at dx.doi.org/10.1007/978-1-4020-9190-2_19

In passive circumstellar disks, the infrared emission is generally assumed to originate almost entirely from dust; the emissivity of the inner, dust-free gaseous part of the disk, at radii smaller than the dust sublimation radius, is negligible. In an actively accreting disk, on the other hand, viscous dissipation of energy in the inner dust-free gaseous part of the accretion disk can heat the gas to high temperatures and give rise to significant amounts of infrared emission from optically thick gas. The inner edge of this gas accretion disk is expected to be located a few stellar radii above the stellar surface, where the hot gas is thought to be channeled towards the star via magnetospheric accretion columns. While the magnetospheric accretion columns are too small to be resolved in our interferometric data (3R correspond to 0.09 AU or 0.12 mas), infrared emission from hot gas between the dust sublimation radius and the stellar surface should be clearly distinguishable from the thermal emission of the dusty disk due to the different temperatures of these components and the resulting characteristic slope in the NIR- and MIR-visibilities. As MWC 147 is a quite strong accretor (M˙ acc ≈ 10−5 M yr−1 ; Hillenbrand et al. [1]), significant infrared emission from the inner gaseous accretion disk is expected. Muzerolle et al. [8] found that, even for smaller accretion rates, the gaseous inner accretion disk is several times thinner than the puffed-up inner dust disk wall and is optically thick (both in radial as well as in the vertical direction). In order to add the thermal emission from the inner gaseous disk to our radiative transfer models, we assume the radial temperature power-law by Pringle [9]. Including the accretion luminosity from an inner gaseous disk to the model improves the agreement between model predictions and observed visibilities strongly. With a flared

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Fig. 4 Observables corresponding to our best-fit radiative transfer model with optically-thick inner gas disk (similar to Fig. 3), yielding good agreement (χr2 ≈ 1.0). A colour version of this figure is available at dx.doi.org/10.1007/978-1-4020-9190-2_19

disk geometry and an accretion rate of M˙ acc = 9 × 10−6 M yr−1 , both the SED and the interferometric visibilities are reasonably well reproduced (χr2 = 1.0, see Fig. 4).

4 Conclusions and Outlook Our VLTI interferometric observations of MWC 147 constrain, for the first time, the inner circumstellar environment around a Herbig Be star over the wavelength range from 2 to 13 µm. We find evidence that the NIR emission of MWC 147 is dominated by the emission from optically-thick gas located inside of the dust sublimation radius, while the MIR also contains contributions from the outer, irradiated dust disk. Our study demonstrates the power of infrared spectro-interferometry to probe the inner structure of the disks around young stars and to disentangle multiple emission components. Future investigations on YSO accretions disks will benefit substantially from the proposed 2-nd generation VLTI instruments, such as MATISSE, increasing not only the number of recorded baselines, but also expanding the spectral coverage to the L and M band.

References 1. L.A. Hillenbrand, S.E. Strom, F.J. Vrba, J. Keene, Astrophys. J. 397, 613 (1992)

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2. A.B. Men’shchikov, T. Henning, Astron. Astrophys. 318, 879 (1997) 3. J.D. Monnier, R. Millan-Gabet, Astrophys. J. 579, 694 (2002) 4. R.L. Akeson, D.R. Ciardi, G.T. van Belle, M.J. Creech-Eakman, E.A. Lada, Astrophys. J. 543, 313 (2000) 5. J. Hernández, N. Calvet, C. Briceño, L. Hartmann, P. Berlind, Astron. J. 127, 1682 (2004) 6. K. Ohnaka et al., Astron. Astrophys. 445, 1015 (2006) 7. S. Kraus, T. Preibisch, K. Ohnaka, arXiv:0711.4988 (2007) 8. J. Muzerolle, P. D’Alessio, N. Calvet, L. Hartmann, Astrophys. J. 617, 406 (2004) 9. J.E. Pringle, Annu. Rev. Astron. Astrophys. 19, 137 (1981)

Multi-epoch VLTI/MIDI Observations of the Carbon-rich Mira Star V Oph K. Ohnaka, T. Driebe, G. Weigelt and M. Wittkowski

1 Introduction The driving mechanism of mass outflows in Mira variables has not yet been fully understood. Recent progress in optical and infrared interferometric techniques has been contributing to studies of the region between the top of the photosphere and the innermost region of the circumstellar dust shell, exactly where mass outflows are expected to be initiated. Infrared interferometric observations of oxygen-rich Mira stars have spatially resolved dense, warm (∼ 1000–2000 K) molecular layers consisting of H2 O, SiO, and CO and extending to ∼ 2–3 R (e.g., Mennesson et al. [1]; Perrin et al. [6]; Ohnaka et al. [4]). The optically thick emission from the warm molecular layers affect the apparent size of Mira stars: the object appears larger than the star itself at wavelengths where the opacities of the above molecular species are higher. In particular, H2 O has strong spectral features in the mid-infrared, which causes the angular size to increase from the near-infrared to the mid-infrared as observed toward oxygen-rich Miras. However, these studies have been limited to oxygen-rich or S-type (C/O ≈ 1) Mira variables up to now, and the physical properties of the outer atmosphere of carbon stars and their temporal variations have not yet been well probed. Particularly, mid-infrared interferometric observations of carbon stars are still very scarce, as are multi-epoch measurements to follow temporal variations of the photosphere and the circumstellar environment. The MID-infrared Interferometric instrument (MIDI) at the ESO’s Very Large Telescope Interferometer (VLTI) is well suited for a study of the circumstellar environment close to the star. It enables us to obtain visibilities with a spectral resolution of 30 or 230 from 8 to 13 µm, where the absorption features due to C2 H2 and HCN as well as the SiC and (featureless) amorphous carbon dust emission are observed for carbon stars. We present the first multi-epoch N band spectro-interferometric observations of the carbon-rich Mira star V Oph using MIDI. K. Ohnaka () Max-Planck-Institut für Radioastronomie, Auf dem Hügel 69, 53121 Bonn, Germany e-mail: [email protected] A. Moorwood (ed.), Science with the VLT in the ELT Era, Astrophysics and Space Science Proceedings, doi: 10.1007/978-1-4020-9190-2_20, © Springer Science + Business Media B.V. 2009

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Fig. 1 a: N -band uniform-disk diameters of V Oph observed with MIDI. The filled circles, diamonds, and triangles represent the observational results obtained at phases 0.18, 0.49, and 0.65, respectively. The photospheric size estimated from the observed bolometric flux and effective temperature of V Oph is shown with the dashed–dotted line. The solid, dotted, and dashed lines represent the best-fit models for phases 0.18, 0.49, and 0.65, respectively (see Sect. 3). b: N -band spectra observed at three phases. The hatched region marks the wavelength range severely affected by the telluric ozone absorption. A colour version of this figure is available at dx.doi.org/10.1007/978-1-4020-9190-2_20

2 Observations V Oph was observed with MIDI on six nights between April and September 2005 at baseline lengths of 42–124 m using four 8.2 m Unit Telescopes (UT2–UT4, UT1– UT4, and UT2–UT3 baseline configurations). The six data sets were binned into three different phases of 0.18, 0.49, and 0.65 with two data sets assigned to each phase. A prism with a spectral resolution of λ/λ ≈ 30 at 10 µm was used to obtain spectrally dispersed fringes between 8 and 13 µm. The details of our MIDI observations and the data reduction are described in Ohnaka et al. [5]. At each phase, we fit the observed visibilities at a given wavelength with a uniform disk to obtain characteristic angular sizes, and the resulting uniform disk diameters in the N -band are shown in Fig. 1a. The figure illustrates a temporal variation of the angular size of V Oph: the star appears smaller at minimum light (phase 0.49) than at post-maximum (phase 0.18) or post-minimum (phase 0.65). Also, the N -band angular sizes obtained at three epochs are significantly larger than the photospheric size of 4–6 mas which is estimated from the bolometric flux and the effective temperature. Only the 9–10 µm uniform-disk diameters obtained at phase 0.49 are rather close to the photospheric size. This means that we detect the same trend as found in oxygen-rich Mira stars: the mid-infrared angular sizes are much larger than the photospheric angular size. Figure 1b shows the absolutely calibrated N -band spectra of V Oph obtained at the three phases. The figure reveals that the spectra are characterized by absorption between 8 and 9 µm and a broad emission feature centered at ∼11.3 µm. The feature seen between 8 and 9 µm is actually the tail of the broad absorption feature centered at ∼ 7.5 µm due to the C2 H2 ν4 + ν5 band and the HCN 2ν2 band, although this

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HCN band extends only up to ∼ 8.3 µm. The broad emission feature at ∼ 11.3 µm is attributed to SiC dust. In addition to these conspicuous features, a broad absorption feature centered at ∼ 14 µm due to the C2 H2 ν5 and HCN ν2 bands extends down to ∼ 11 µm, where it becomes blended with the SiC emission feature. Comparison between the spectra at phases 0.18 and 0.49 reveals that the 8 µm absorption feature becomes somewhat stronger at minimum light than at post-maximum.

3 Modeling with C2 H2 Layers and a Dust Shell In order to interpret the observed N -band visibilities and spectra, we construct a model consisting of optically thick C2 H2 gas close to the star and a more extended dust shell. The C2 H2 gas is approximated with two layers (hot and cool C2 H2 layers), while the dust shell consists of amorphous carbon and SiC grains. Since the extended dust shell is resolved out at the baseline lengths used for our MIDI observations, we fix the inner boundary radius of the dust shell to 2.5 R , which corresponds to the condensation temperature of amorphous carbon (∼ 1600 K). The radii, temperatures, and C2 H2 column densities of the hot and cool C2 H2 layers as well as the dust shell’s optical depth and the fraction of SiC dust are varied to search for the parameter set which best reproduces the observed visibilities and spectra. The details of our modeling is described in Ohnaka et al. [5]. Comparison between the best-fit models and the observations for three phases is shown in Fig. 1a, while the parameters of these models are plotted in Fig. 2. It should be stressed here that we compare the model visibilities to the observations, not uniform-disk diameters, because the uniform-disk diameter fit is merely for obtaining some representative angular size of the object. Detailed comparison between the model and the observational data is presented in Ohnaka et al. [5]. The radius of the cool C2 H2 layer shows noticeable variations: the layer is more extended at post-maximum (phase 0.18) and post-minimum (phase 0.65) than at minimum light (phase 0.49). Furthermore, the column density of the cool C2 H2 layer appears to be the smallest at minimum light, although the uncertainties of the derived column densities are rather large. The temperature of the cool layer is found to exhibit no remarkable temporal variation. On the other hand, the hot C2 H2 layer shows a temporal variation somewhat different from that of the cool layer. The hot layer becomes more extended, denser, and a little cooler at phase 0.65 compared to the preceding two phases. The dust shell also shows noticeable temporal variations: its optical depth is the smallest at minimum light, while the fraction of SiC dust (∼ 20%) does not show significant changes. The observed temporal variation of the C2 H2 layers and the dust shell can be qualitatively understood by dynamical model atmospheres of carbon-rich Mira stars with dust formation included (Nowotny et al. [2, 3]). In these models, the largeamplitude stellar pulsation causes the formation of shock fronts propagating outward, and polyatomic molecules such as C2 H2 and HCN are likely to form behind shock fronts where the density is high. Therefore, it is possible that the observed

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Fig. 2 Temporal variations of the radii, column densities, and temperatures of the hot and cool C2 H2 layers (a, b, and c), and the optical depth of amorphous carbon dust in the dust shell (d) derived from our modeling of the MIDI data on V Oph

temporal variation of the radii of the C2 H2 layers results from changes of the positions of C2 H2 formation behind shock fronts. This means that the cool layer extending to ∼ 1.8 R at phase 0.18 moves farther outward and may no longer contribute to the N -band flux at phase 0.49. And the small radius of the cool layer obtained at phase 0.49 suggests that a new C2 H2 formation front may just be emerging at a smaller radius and that this front may build up and move outward from phase 0.49 to 0.65. The phase dependence of the grain optical depths we derived for V Oph can also be interpreted in a similar manner. Dust can form behind shock fronts where the density is sufficiently high and the temperature is below condensation temperature. If such a hot dust layer is responsible for the higher grain optical depths we derived for phase 0.18 (post-maximum light), the dust layer is expected to expand from phase 0.18 to 0.49 (minimum light) and become diluted. At minimum light, a new (inner) dust layer can be forming, but still with a low degree of condensation, which can qualitatively explain the smallest grain optical depths at minimum light. When this new dust layer has fully grown, it can provide higher grain optical depths at post-minimum phases. However, we note that the mass loss rate of V Oph (∼ 10−8 M yr−1 ) is by two orders of magnitude lower than that of the models of Nowotny et al. [2, 3]. It is necessary to compare with dynamical models computed

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with parameters corresponding to V Oph for a better understanding of the physical processes responsible for molecule and dust formation close to the star.

4 Conclusion Our multi-epoch MIDI observations and modeling of the carbon-rich Mira star V Oph have revealed that the observed N -band visibilities and spectra can be reproduced by optically thick C2 H2 layers and a dust shell consisting of amorphous carbon and SiC. The temporal variation of the properties of the C2 H2 layers and the dust shell can be qualitatively explained by dynamical model atmospheres. In order to obtain a more comprehensive picture of how the outer atmosphere and the dust shell respond to stellar pulsation, it is indispensable to carry out near-infrared interferometric observations simultaneously with MIDI observations. Such coordinated observations are feasible with VLTI, using MIDI and the near-infrared beam combiner, AMBER (Petrov et al. [7]), which is capable of measuring visibilities in the J , H , and K bands with three different spectral resolutions of 75, 1500, and 12 000.

References 1. B. Mennesson, G. Perrin, G. Chagnon et al., Astrophys. J. 579, 446 (2002) 2. W. Nowotny, B. Aringer, S. Höfner, R. Gautschy-Loidl, W. Windsteig, Astron. Astrophys. 437, 273 (2005a) 3. W. Nowotny, T. Lebzelter, J. Hron, S. Höfner, Astron. Astrophys. 437, 285 (2005b) 4. K. Ohnaka, J. Bergeat, T. Driebe et al., Astron. Astrophys. 429, 1057 (2005) 5. K. Ohnaka, T. Driebe, G. Weigelt, M. Wittkowski, Astron. Astrophys. 466, 1099 (2007) 6. G. Perrin, S.T. Ridgway, B. Mennesson et al., Astron. Astrophys. 426, 279 (2004) 7. R.G. Petrov, F. Malbet, G. Weigelt et al., Astron. Astrophys. 464, 1 (2007)

A Mid-infrared Interferometric Study of the Circumstellar Environment of Dusty OH/IR Stars with VLTI/MIDI T. Driebe, K. Ohnaka, K. Murakawa, K.-H. Hofmann, D. Schertl, G. Weigelt, T. Verhoelst, O. Chesneau, A. Domiciano de Souza, D. Riechers, M. Schöller and M. Wittkowski

OH/IR stars are evolved stars exhibiting large infrared excess as well as prominent OH maser emission. The vast majority of this heterogeneous object class are highlyevolved low- and intermediate-mass stars ascending the upper part of the asymptotic giant branch (AGB). Thus, these stars are in the final phase of their AGB evolution which is characterized by intensive mass loss with mass-loss rates that can reach up to 10−4 M . This high mass loss leads to the development of an extended, usually optically thick dusty circumstellar envelope which can be well studied with infraredinterferometric observations.

1 Observations and Modeling For a small sample of dust-enshrouded OH/IR stars, we carried out N -band observations with MIDI, the mid-infrared interferometric instrument of ESO’s VLTI at Table 1 Overview of the MIDI measurements of our small sample of OH/IR stars. N is the total number of MIDI observations, and Bp and P.A. give the range of projected baselines and position angles IRAS

RA

DEC

Obs. time

number

[2000]

[2000]

[2006]

Bp

P.A.

[m]

[◦ ]

13328-6244

13:36:19

−62:59:19

Feb.–Jun.

13581-5930

14:01:40

−59:44:34

Apr.–Jun.

21

13.4–31.9

42–120

17

13.6–32.0

16292-5004

16:33:00

−50:10:56

May.–Jun.

20–112

12

13.5–32.0

14–102

16444-4527

16:48:05

−45:32:55

May.–Aug.

4

11.0–16.0

63–117

16460-4022

16:49:34

−40:27:39

Apr.–Aug.

10

10.2–16.0

42–115

18172-2305

18:20:19

−23:03:51

Jun.–Aug.

9

14.0–32.0

53–90

N

T. Driebe () Max-Planck-Institut für Radioastronomie, Bonn, Germany e-mail: [email protected] A. Moorwood (ed.), Science with the VLT in the ELT Era, Astrophysics and Space Science Proceedings, doi: 10.1007/978-1-4020-9190-2_21, © Springer Science + Business Media B.V. 2009

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Fig. 1 Comparison of the MIDI observations of IRAS 13581-5930 with a 1-D dust radiative transfer model obtained with DUSTY [1]; upper panels: model (blue squares) and observations (red bullets with error bars) for two of our MIDI visibility measurements using the E0-G0-16m and D0-G0-32m baselines, respectively; lower left: N -band model spectrum (solid red line), calibrated MIDI spectrum (blue bullets with error bars), and IRAS LRS spectrum (green squares); lower right: radial model intensity profiles for different wavelengths across the N band. The labels give some basic model parameter values. rin and Tin are the radius and temperature at the inner dust shell boundary, and τV is the optical depth at 0.55 µm. A colour version of this figure is available at dx.doi.org/10.1007/978-1-4020-9190-2_21

Cerro Paranal. All measurements were obtained using the 1.8 m-class auxiliary telescopes (ATs). A summary of the MIDI observations is given in Table 1. For the data interpretation, we started to carry out radiative transfer calculations using the 1-D code DUSTY [1] and the 2-D Monte Carlo code presented in [2] to simultaneously model the wavelength dependence of the visibility across the N band and the midinfrared spectrum of our sample stars as they have been observed with VLTI/MIDI. An example of our MIDI observations and DUSTY modeling is shown in Fig. 1.

References 1. Z. Ivezi´c, M. Elitzur, Astrophys. J. 445, 415 (1995) 2. K. Ohnaka, T. Driebe, K.-H. Hofmann et al., Astron. Astrophys. 445, 1015 (2006)

The Closest Dusty Cloud Ever Detected Around a R CrB Variable Star Using the VLTI/MIDI Instrument I.C. Leão, P. de Laverny, O. Chesneau, D. Mékarnia and J.R. De Medeiros

1 Introduction In recent VLTI/MIDI observations (see Leão et al. [2] for more details), we collected N-band dispersed (7.5–13.5 µm, λ/λ = 230) visibilities of RY Sgr using two telescope pairs: UT1–UT4 and UT3–UT4. We obtained 5 different baselines with 7 runs (see Fig. 1) under rather good atmospheric conditions. The data represent a snapshot obtained within June 2005.

2 Results Monochromatic model: In a first-order analysis, we neglected any variations with wavelength of the source geometry and fitted the curves with simple monochromatic geometric models, obtaining a global morphology in the N-band. Proposed model: central star + dust cloud + extended circumstellar envelope. Results for the best fit to the visibility curves (see Fig. 1): • • • •

Separation between the cloud and the central star: 16 ± 1 mas. PA of the cloud: 75◦ ± 10◦ (modulo 180◦ ). FWHM of the Gaussian CSE: 18 ± 3 mas. Flux contributions to the total N-band flux of the whole system: 10% ± 2% for the central star and 8% ± 2% for the cloud.

Chromatic model: In a more detailed analysis, we considered possible spectral variations of the model parameters. Very good fits were found with parameter values close to those given in the global analysis, validating the global morphology. In addition, the CSE FWHM grows slightly from 17 to 19 mas (within uncertainty of ±3 mas) toward larger wavelengths. No significant variation of the stellar and cloud fluxes were observed. I.C. Leão () OCA, Dpt. Cassiopée, CNRS – UMR6202, BP4229, 06304 Nice cedex 4, France e-mail: [email protected] A. Moorwood (ed.), Science with the VLT in the ELT Era, Astrophysics and Space Science Proceedings, doi: 10.1007/978-1-4020-9190-2_22, © Springer Science + Business Media B.V. 2009

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Fig. 1 Dispersed visibilities as a function of spatial frequency and PA. Solid curves are the observed visibilities. Non-solid curves are the best theoretical fits. The total error bars on the visibilities are shown at each curve extremity to illustrate their mean variations as a function of wavelength. The chart at the top left corner illustrates the u–v plan with the 7 observed baselines, labeled (a) to (g). A grayscale background representing the best theoretical fit is also shown in the chart. The geometric model of the image is shown in the middle panel at left

We explored the inner 60 mas (∼ 110 AU) of the RY Sgr environment with a dynamic range better than 20. The FWHM of the dusty envelope is ∼ 120 R∗ , (∼ 35 AU). The cloud separation is ∼ 100 R∗ (∼ 30 AU) from the center. This is the closest dusty cloud ever detected around a R CrB-type variable since the first direct detection with NACO by de Laverny and Mékarnia [1].

The Closest Dusty Cloud Around a R CrB Using VLTI/MIDI

References 1. P. de Laverny, D. Mékarnia, Astron. Astrophys. 428, L13 (2004) 2. I.C. Leão, P. de Laverny, O. Chesneau et al., Astron. Astrophys. 466, L1 (2007)

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Part III

Future VLT and VLTI Science Priorities

A Twenty Year Science Vision for European Astronomy Guy Monnet and Tim de Zeeuw

1 Introduction Three years ago, a group of European funding agencies, plus ESA and ESO, put forward the Astronet program to ensure close coordination of the astronomical developments within the Continent. One prime objective of this ERANET contract financed by the European Commission is to produce a long-term (∼ 20-year) strategic plan, similar in scope and content to the recurrent US decadal surveys. This is being pursued through a two-step process, first articulating a ’Science Vision’ with a set of prioritised science goals as well as an analysis on the generic facilities needed to reach them, followed by an Infrastructure Roadmap outlining the ways and means to implement the Vision. The first step is now completed, with the Science Vision document published at the end of September 2007; this paper presents the process that has been followed and the main conclusions reached.

2 Astronomy The scientific domain considered is astronomy at large, i.e., the study of every physical process beyond planet Earth, through the whole electromagnetic spectrum from γ -rays to radio, plus neutrinos and gravitational waves, and for all modes of observation (e.g. imaging, spectroscopy, time series). Here, “study” should be taken in the widest sense, encompassing such crucial aspects of present day astronomy as combining very different types of observation and putting forward massive exploration searches and surveys, including in situ measurements throughout the Solar System. Astronomy is a physical science driven by observation; testing models versus observations is thus a vital element for which theory and numerical simulations are essential components. That also implies a strong need to maintain close links with sister research fields in physics, chemistry, computer science, laboratory experiments, geophysics, and biology. Astronomy benefits from and drives advances in technology. As a result, it is now possible e.g. to (a) study objects over 95% of the age of the Universe, (b) detect and G. Monnet () ESO, Karl-Schwarzschild-Str. 2, 85748 Garching, Germany e-mail: [email protected] A. Moorwood (ed.), Science with the VLT in the ELT Era, Astrophysics and Space Science Proceedings, doi: 10.1007/978-1-4020-9190-2_23, © Springer Science + Business Media B.V. 2009

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Fig. 1 Historical gains in spatial resolution and sensitivity. A colour version of this figure is available at dx.doi.org/10.1007/978-1-4020-9190-2_23

characterise planets around other stars, (c) use not only photons but also particles and possibly soon gravitational waves to study celestial sources and (d) explore in situ virtually any object in the Solar System. On the exploitation side, it is possible to simulate complicated astrophysical processes and analyse large data streams. These observing feats are made possible by the wide range of general purpose or dedicated observatories, on the ground or in space, presently in operation or under development. As an illustration of the increasing power of these observing facilities, Fig. 1 shows the gains in angular resolution in the optical/IR domain and in sensitivity at radio wavelengths, through the past century and up to 2025. Astronomy is important for society and culture. Examples can be as diverse as its enabling role for navigation devices and mobile phones, the relevance of asteroid impacts on Earth and the way astronomy puts in context our own existence with the discovery of many other worlds and their potential for life. Equally important, as-

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tronomy helps attract young people to the physical sciences and is a major provider of exciting discoveries for amateur astronomers and the layman alike.

3 Predicting the Future: Ambitious Science Goals There is general agreement on four key questions of interest to a broad science community as well as to the general public, and for which significant breakthroughs should be achieved in the coming two decades from ground and space telescopic observations, combined with modelling and theoretical developments: A. B. C. D.

Do we understand the extremes of the Universe? How do galaxies form and evolve? What is the origin and evolution of stars and planets? How do we (and the Solar System) fit in the global picture?

Resolving these questions which are among the most fundamental in all of science requires setting up ambitious European plans through 2025, with several billions of Euros for new investments and operation of present and future facilities. The European Commission will fund only a modest fraction and the bulk of the support would have to come from national funding agencies. In turn, the funding agencies do need as a prerequisite a comprehensive plan, quite similar to the US decadal surveys, covering all of astronomy including links with neighbouring fields. The main rationale for Astronet is indeed to develop such a plan together with the entire European astronomical community.

4 Developing the Science Vision The four key questions listed above were tackled by a Science Vision Working Group (SVWG) and four supporting panels, bringing together about 50 scientists, appointed by the funding agencies, with a good distribution of expertise, gender and nationalities. The SVWG consisted of the panel chairs and co-chairs, plus eight members at large. Each of the panels concentrated on one of the four key questions. They looked broadly at the combination of observations, simulations, laboratory experiments, interpretation and theory required for major breakthroughs and identified corresponding generic observing facilities crucial for these endeavours. The panels and the SVWG systematically used the large amount of already available information e.g. in the various national strategic plans, ESA’s Cosmic Vision, the ASPERA road mapping exercise [1] (http://www.aspera-eu.org/images/stories/files/Roadmap.pdf) and the three ESA-ESO topical studies [2–4] (http://www.stecf.org/coordination/ esa_eso/). This work led to specific recommendations incorporated in a draft version of the Science Vision document which was released to the community late 2006 and subjected to a web-based forum discussion. The amended draft was then presented

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Fig. 2 Science Vision book cover. A colour version of this figure is available at dx.doi.org/10.1007/978-1-4020-9190-2_23

and discussed during a Symposium in Poitiers on January 23–25, 2007. Its 228 participants from 31 countries enrolled themselves in the four much enlarged panels for in-depth discussions. Much constructive input was generated, leading to further sharpening of the science requirements and improved balance across the fields. The Science Vision book [5] was finalised by the panels, the SVWG and the editors, released end of September 2007 and widely distributed (Fig. 2). It can be accessed from http://www.eso.org/public/outreach/press-rel/pr-2007/pr-44-07.html. The intended audience is quite varied, from the astronomical community at large to funding agencies and ministries. An early condensed version, outlining the main science priorities in each thematic area, together with the present or new (generic) facilities requested to tackle them, was input to the Infrastructure Roadmap exercise on mid-May 2007, as it began to embark on the next—and last—leg towards building a comprehensive strategic plan for European astronomy.

5 Recommendations Each of the main book chapters describes the background of one of the four fundamental questions, spells out the key scientific goals and describes the most promis-

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ing approaches. A clear distinction is made throughout the presentation between ‘Essential’ and ‘Complementary’ facilities with respect to a particular science goal, as well as to their ‘Current’ or ‘Future’ status. In addition, a number of crossdisciplinary requirements were analysed, viz. theory and simulations, computing resources, astronomical data management and laboratory astrophysics. A highly condensed summary of the results and recommendations, with for each of the four themes the main goals identified and the essential facilities needed, is presented below. Current facilities in operation or construction are shown in standard fonts; possible future generic ones in italics.

5.1 Panel A: Do We Understand the Extremes of the Universe? • Evolution of dark energy with cosmological epoch [Planck, SKA, X-ray Survey, Wide-Field Space Telescope]. • Consistent picture of dark matter and dark energy [Planck, Dark Matter Detectors, SKA, X-ray Survey, Wide-Field Space Telescope, ELT, CTA]. • Search for relic gravitational waves from CMB polarisation [CMB Polarisation Satellite]. • Direct detection of gravitational waves in strong gravity regions [Ground Gravitational Detectors, LISA]. • Direct studies near supermassive black-hole horizons [SWIFT, XMM, INTEGRAL, 8–10 m Telescopes, VLTI, HESS, X-ray Observatory, Sub-mm VLBI]. • Understanding the astrophysics of compact objects [SWIFT, XMM, INTEGRAL, X-ray Survey, ELT, X-ray Observatory]. • Understanding the origin of high-energy neutrinos and cosmic rays [HESS, AUGER, CTA].

5.2 Panel B: How Do Galaxies Form and Evolve? • Mapping the early growth of matter density fluctuations [LOFAR, Planck, ELT, SKA]. • Detect the Universe first lights [JWST, X-ray spectroscopy]. • Constrain dark energy equation of state through evolution of galaxy clusters [XMM]. • Inventory of the metal content over cosmic time [8–10 m Telescopes, Large UV Space Telescope]. • Metallicity of the warm-hot local IGM [HST, X-ray spectroscopy]. • Evolution of the Hubble sequence and central black-holes growth [8–10 m Telescopes, ALMA, ELT, SKA, Large cooled IR Telescope]. • Dynamics and evolution census of stars in our Galaxy and the nearest ones [GAIA].

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5.3 Panel C: What Is the Origin and Evolution of Stars and Planets? • Determine the full sequence of star formation [8–10 m Telescopes, HST, Herschel, JWST, ALMA, Large mm Single Dish, ELT, Large UV Space Telescope, X-ray Observatory, SKA, Sub-mm Space Interferometer]. • Unveil stellar structure and evolution [8–10 m Telescopes, HST, GAIA, ELT, Large UV Space Telescope, X-ray Observatory, High Precision Photometric Network]. • Understand the life-cycle of the interstellar medium [8–10 m Telescopes, Herschel, JWST, ALMA, Large mm Single Dish, ELT, Large UV Space Telescope, X-ray Observatory, SKA]. • Determine how planetary systems form and evolve [Herschel, JWST, ALMA, ELT, SKA, Sub-mm Space Interferometer]. • Explore the diversity of planetary systems [8–10 m Telescopes, GAIA, ELT, High Precision Photometric Network]. • Make a census of Earth-like planets in the habitable zone and search for biomarkers [ELT, Sub-mm Space Interferometer, High Precision Photometric Network].

5.4 Panel D: How Do We (and the Solar System) Fit in the Global Picture? • Understand physical processes in Solar System plasmas [SOHO, STEREO, Hinode, ACE, CLUSTER, SDO, Ulysses, THEMIS, SST, GREGOR, High-latitude Solar Orbiter, Radio Imaging, Large Solar Telescope, EUV-X Satellite, Flying Formation, Ground-based Radar Network]. • Develop a unified picture of the Sun and the heliosphere [SOHO, STEREO, Hinode, ACE, CLUSTER, SDO, Ulysses, High-latitude Solar Orbiter, Radio Imaging, Flying Formation]. • Understand what drives solar variability [SOHO, STEREO, Hinode, ACE, CLUSTER, SDO, Ulysses, High-latitude Solar Orbiter]. • Role of magnetic fields and turbulence in the formation & evolution of the Solar System [Genesis, Rosetta, In Situ Probes]. • Determine the dynamical history & composition of comets and asteroids [Genesis, Rosetta, In Situ Probes]. • Develop full models of planet internal structures & atmospheres [Mars probes, Venus Express, Cassini, Bepi-Colombo, Space Missions, Mars Sample Return Mission, JWST, ELT]. • Search for liquid water on Mars & Europa and Study of Titan atmosphere [Mars probes, Venus Express, Cassini, Bepi-Colombo, Space Missions, Mars Sample Return Mission JWST, ELT].

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6 Conclusions Answering these key scientific questions requires an optimal use of existing and still competitive facilities, plus those being constructed or yet to be decided within the next decades. The latter comprises a next generation of ground-based telescopes plus new specific space observatories and missions, in particular in the context of ESA’s cosmic vision. Studies of individual objects, large scale dedicated surveys and investigations of the time-domain will be essential. All observations will have to be supported by theoretical programs, extensive modelling and numerical simulations, and in many domains by laboratory experiments. While decidedly endowed with a European flavor, the Science Vision was developed with a keen knowledge of the largely parallel efforts that are and will occur from the North-American continent and elsewhere. With soon an Infrastructure Roadmap, hopefully leading to an effective and timely implementation of the plans outlined in the Science Vision, Europe will have the opportunity to build an optimum mix of cooperation and competition with the other main players worldwide, reaping a large fraction of the rich scientific harvest ahead. Finally, considering the more specific topic at this conference, it is gratifying to see the strong need for continued exploitation and upgrade of 8–10 m class telescopes, including the VLT/VLTI. The Science Vision is also a great help to focus the current E-ELT design phase and provides compelling scientific arguments for its realisation. Acknowledgements Developing a common Science Vision for European Astronomy was made possible by the help and trust of the Astronet partners and the Commission, and above all from the essential contribution by the community at many different levels.

References 1. ApPEC Status and Perspective of Astroparticlephysics in Europe (ASPERA, 2007) 2. M. Perryman et al., ESA-ESO Working Groups Report No 1: Extra-Solar Planets (ESA, ESO, 2005) 3. T.L. Wilson et al., ESA-ESO Working Groups Report No 2: The Herschel ALMA–Synergies (ESA, ESO, 2006) 4. J. Peacock et al., ESA-ESO Working Groups Report No 3: Fundamental Cosmology (ESA, ESO, 2006) 5. P.T. de Zeeuw, F.J. Molster (eds.), A Science Vision for European Astronomy (ASTRONET, 2007)

Baryonic Acoustic Oscillations Gavin Dalton

Since the discovery and confirmation of the requirement for a substantial Λ-like Dark Energy component to the Universe there has been substantial progress on the development of various observational approaches to understanding the nature of the Dark Energy. In this paper I will review some of the recent refinements to the observational evidence, and discuss opportunities for the VLT to make a substantial contribution to this field over the next ten years.

1 Introduction The evidence for the existence of a substantial energy component to the Universe comes from a number of different cosmological observations, with the earliest indications coming from large-scale structure (e.g. [7]). The best current constraints on Dark Energy come from a combination of the observed luminosities of type-1a supernovae [10], the high angular resolution measurements of the cosmic microwave background [5], and large-scale structure measurements at low redshift from the 2dFGRS and SDSS surveys [6]. This evidence has led to many attempts to develop models for Dark Energy, and these models usually attempt to constrain the form of the equation of state of the Dark Energy as w(z) = P /ρ. Key questions for the investigation of the Dark Energy phenomenon at the next level include whether Dark Energy is simply a constant energy term, i.e. the Cosmological Constant, (w(z) = −1), or whether Dark Energy represents some more novel physics such as a departure from classical General Relativity. At present, however, it seems clear that there is no single compelling theory, and that progress in this area must be driven by new observations, for which there are only two distinct approaches: We can either study Dark Energy by means of geometrical tests of the Universal expansion (Supernovae or Baryonic Acoustic Oscillations, BAOs), or we can study the growth of structure in the Universe (by weak lensing, cluster abundances, or the integrated Sachs-Wolfe effect) and relate these observations to models of large-scale structure evolution from large numerical simulations. The recent report from the US Dark Energy Task Force (DETF, [4]) produced G. Dalton () Astrophysics, University of Oxford, Denys Wilkinson Building, Keble Road, Oxford OX1 3RH, UK e-mail: [email protected] A. Moorwood (ed.), Science with the VLT in the ELT Era, Astrophysics and Space Science Proceedings, doi: 10.1007/978-1-4020-9190-2_24, © Springer Science + Business Media B.V. 2009

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a figure of merit estimator for the various approaches, with the broad conclusions that the geometric approaches are concrete and well-understood, whilst the growthbased approaches represent a bigger challenge in terms of model-dependence and risk, but provide the possibility of more substantial progress gains in the longer term. The clear recommendation of the DETF report was for a multi-strand approach to addressing the Dark Enery issue.

2 Baryonic Acoustic Oscillations The phenomenon of BAOs occurs because the harmonics of a sound wave corresponding to the physical size of the Universe at the time of recombination (i.e. the same physical scale as the primary peak in the CMB power spectrum) remain imprinted at a low level on the present day matter power spectrum. The signal is highest at large scales, and successively damped at smaller scales because of the relatively low amplitude of the power spectrum. At the smallest scales the signal is erased completely due to the onset of non-linear evolution within gravitationally bound structures. The best current results come from a sample of 520 000 galaxies drawn from the SDSS DR5 [2], shown in Fig. 1 together with the 3 year WMAP data.

Fig. 1 (Top) 3 year WMAP data shown with the best fit WMAP model overplotted (Ωm = 0.24, [3]). (Bottom) The power spectrum of galaxy clustering from 520 000 galaxies from SDSS DR5, with the smooth component removed to highlight the oscillations [2]

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Fig. 2 Residual power spectra for the combined 2dFGRS and SDSS sample (a), the SDSS LRG sample (b) and the full combined dataset (c) (from [1])

3 Oscillations as a Function of Redshift Recent progress with the SDSS main survey and its Large Red Galaxy (LRG) component has now allowed a secondary test to be made where the oscillations can now be used as a measure of relative distance [1]. Figure 2 shows the model-subtracted power spectra for the combined SDSS and 2dFGRS low refshift samples (608 000 galaxies with a median redshift of z ∼ 0.2), together with that for the SDSS LRG sample (79 000 galaxies with a median redshift of z ∼ 0.35). Fitting the position of the first acoustic peak in both samples gives the ratio of the volume averaged distances: (1 + z)2 czDA (z)2 1/3 Dv (z) = , (1) H (z) with the result D0.35 /D0.2 = 1.812 ± 0.060, compared with a prediction of 1.67 for flat CDM, which is intriguing, but does not yet represent compelling evidence for a discrepancy. Including the constraints from supernovae measurements gives 1.66 ± 0.01!

4 Upcoming Surveys The results presented in Sect. 3 probably represent the limit of what can be done with current surveys. SDSS is now preparing for an extension of the LRG sample,

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Fig. 3 Simulated P (k) for the proposed FAST-SOUND survey. A colour version of this figure is available at dx.doi.org/10.1007/978-1-4020-9190-2_24

and there will be photometric redshift surveys from PanSTARRS, VST/VISTA and the Dark Energy Survey. The next major redshift survey that may provide new information will be the FAST-SOUND survey to be conducted with Subaru’s FMOS instrument [8, 9]. FMOS will allow a unique large-scale structure probe of the key interval in redshift, 0.7 < z < 1.8 using H-α emission line redshifts for several hundred thousand galaxies (Fig. 3), commencing in 2009. Beyond FMOS, a second generation fibre system on Subaru will deliver visible band spectra for up to 4000 simultaneous targets over a 1.5 degree field of view, with the aim of producing a redshift sample of up to 5 × 106 galaxies out to z ∼ 1.1, specifically targetted at BAO measurements, and yielding an order of magnitude improvement over the current SDSS results [11]. One of the key aspects of a survey of this size will be the ability to divide the power spectrum measurements into angular and line-of-sight directions, giving independent measures of the evolution of dL (z) and H (z), respectively, and providing discrimination between curvature effects and a true dynamical Dark Energy component (Fig. 4).

5 Dark Energy with the VLT? The delivery of a WFMOS-like programme over the next 5–10 years appears to be a key element of any coordinated strategy to tackly the Dark Energy problem. At the time of this meeting, the funding status of WFMOS as a Subaru instrument is far from clear, and so the question must be asked as to whether the VLT might provide a better opening for such a facility. Preliminary estimates suggest that a 2dF-like conversion of one UT to a lightweight top end ring could provide adequate mass margin for such an instrument and corrector without compromising other operational aspects of the VLT/VLTi if careful thought is given to scheduling issues. Possible aspects of such an instrument will be discussed elsewhere in these proceedings.

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Fig. 4 Example of the power of WFMOS-scale surveys arising from the ability to divide measurements into angular and line-of-sight power spectra. A colour version of this figure is available at dx.doi.org/10.1007/978-1-4020-9190-2_24

References 1. 2. 3. 4. 5. 6. 7. 8. 9. 10. 11.

W.J. Percival, S.M. Cole, D.J. Eisenstein et al., Mon. Not. R. Astron. Soc. 381, 1053 (2007) W.J. Percival, R.C. Nichol, D.J. Eisenstein et al., Astrophys. J. 657, 51 (2007) D.N. Spergel, R. Bean, O. Dore et al., Astrophys. J. Suppl. Ser. 170, 377 (2007) A. Albrecht, G. Bernstein, R. Cahn et al., astro-ph/0609591, 2006 D.N. Spergel, L. Verde, H.V. Peiris et al., Astrophys. J. Suppl. Ser. 148, 175 (2003) W.J. Percival, W.J. Sutherland, J.A. Peacock et al., Mon. Not. R. Astron. Soc. 337, 1068 (2002) G.B. Dalton, R.A.C. Croft, G. Efstathiou et al., Mon. Not. R. Astron. Soc. 271, 47 (1994) G.B. Dalton, I.J. Lewis, D.G. Bonfield et al., Proc. SPIE 6269, 136 (2006) K. Glazebrook, C.A. Blake, G.B. Dalton, in Science with FMOS, ed. by T. Maihara, K. Ohta (Kyoto University Press, Kyoto, 2004) P. Astier, J. Guy, N. Regnault et al., Astron. Astrophys. 447, 31 (2006) D. Parkinson, C.A. Blake, B. Basset et al., Mon. Not. R. Astron. Soc. 377, 185 (2007)

Galaxy Formation and Evolution J. Bergeron

1 Established Properties of High Redshift Galaxies Major steps in our knowledge of high redshift galaxies and quasars were made possible by the synergy between the Hubble Space Telescope (HST) and 8–10 m class optical-near infrared ground-based telescopes. The galaxy samples were built from very deep, pencil-beam surveys with the HST and deep, wider area surveys undertaken with ground-based facilities. The most efficient selection technique of high-redshift star-forming galaxies is based on the presence of the Lyman-break within a given broad-band filter. Spectroscopic confirmation is then achieved with 8–10 m class telescopes. The first comparison of the galaxy Luminosity Function (LF) between two cosmic epochs was made for Lyman-Break Galaxies (LBGs) at z ∼ 3 and z ∼ 4 [1]. At these bright magnitudes, there is little evidence for an evolution of the LF shape and star formation density between these two epochs, whereas this is no longer the case at higher z (mostly photometric samples at z 5). There is a deficiency of luminous galaxies at z ∼ 6, with a brightening of M by ∼ 0.6 magnitude between z ∼ 6 and z ∼ 3, but little evolution of the luminosity density and Star Formation (SF) rate density [2]. Sizes of high z galaxies have been measured from HST deep images [2, 3]. They are small: the mean half-light radius is equal to rhl = 2.3((1 + z)/3)−1.05 kpc (assuming Ω , Ωm , h = 0.7, 0.3, 0.7). This corresponds to only a small decrease in apparent size between z = 3 and z = 6, from rhl = 220 mas to 175 mas. At z ∼ 2, the brightest galaxies are several times larger. The surface density, n, of high redshift LBGs is derived from deep HST surveys. Bright LBGs at z ∼ 6, i-dropouts with z884,AB ≤ 25.4, are rare with n 0.015 arcmin−2 . The surface density of idropouts increases rapidly with decreasing luminosities with n (z884,AB ≤ 28.1) 1.4 arcmin−2 [4]. At z ∼ 7–8, one expects about a few tenths of galaxies per arcmin2 at these faint AB magnitudes. Another powerful technique to discover high-redshift galaxies is the detection of Lyα emission from narrow-band imaging. These surveys uncover a star-forming, young galaxy population, while the Lyman-break technique selects both young and older age galaxies. There are several hundreds candidate Lyα Emitters (LAEs) at z > 5.5. They are confirmed by subsequent spectroscopy. Their observed equivalent widths are large, wobs (Lyα) ≥ 120 Å. The number density of LAEs at z ∼ 5.7 is J. Bergeron () Institut d’Astrophysique de Paris, 98bis Boulevard Arago, 75014 Paris, France e-mail: [email protected] A. Moorwood (ed.), Science with the VLT in the ELT Era, Astrophysics and Space Science Proceedings, doi: 10.1007/978-1-4020-9190-2_25, © Springer Science + Business Media B.V. 2009

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n 60 deg−2 , number similar to that of bright LBGs at z ∼ 6 [5]. Their SF rate density is SFRD = 10−3.2 M yr−1 Mpc−3 , i.e. about 10% that of z ∼ 6 faint LBGs with L > 0.3Lz=3 [6]. It should be noted that the highest z object with spectroscopic confirmation, discovered so far, is a LAE at z = 6.96 [7]. Searches for LAEs during the reionization era, at z = 8.8 with ISAAC and SINFONI, have been unsuccessful [8]. Using the z = 5.7 and 6.5 LAE samples, the predicted number down to a line flux f = 1.3 × 10−17 erg s−1 cm−2 is 3–10 deg−2 which implies search areas far larger than previously explored [8]. A major step forward in our knowledge of the internal physics of galaxies is now possible with integral-field spectroscopy. In-depth physical diagnostics from 3D velocity fields of z ∼ 2 bright galaxies with SINFONI yield dynamical mass, merging rate and metallicity [9]. It also led to the discovery of a large rotating disc at z = 2.38 [10]. Another topic explored recently is the history of mass assembly characterized by the morphological evolution of galaxies. Imagers with Adaptive Optic (AO) systems such as NACO has enabled the estimate of the distribution of galaxy morphological types at z ∼ 1, confirming that the fraction of irregular galaxies increases with redshift [11]. Finally, extensive multi-wavelength imaging surveys from the near-infrared to mid-infrared with the Spitzer telescope, and the submillimeter and radio ranges with ground-based telescopes have revealed the existence of a population of rare, bright, massive galaxies at z ∼ 1–3. Their redshifts have been derived either from spectroscopy with Spitzer or optical, near-infrared follow-up with 8–10 m telescopes (e.g. [12, 13]).

2 Galaxy Evolution: Open Questions To better understand the physics of galaxy formation and evolution, we need to assess the effects of different observable properties for identifying the key physical processes that drive the evolutionary trends. This requires high-quality observations of large samples covering a wide redshift range. The star formation history should be known as a function of galaxy mass, type and environment. In addition to redshift determination, galaxy spectra of high signal-to-noise ratio are mandatory to detect stellar and interstellar gas absorption lines, thus enabling an estimate of the chemical composition of galaxies as a function of their masses. Imaging with HST revealed small galaxy sizes in their UV rest-frame; this should be extended in the optical rest-frame to probe the older stellar population. Morphological studies of galaxies must be done for samples large enough to ascertain the importance of interactions and mergers at early epochs. This is also linked to a better knowledge of galaxy clustering versus galaxy luminosity, type and stellar mass. Furthermore, a multi-wavelength approach is essential to understand the connection between the formation of black holes and their host galaxies, i.e. the role of feedback versus accretion. Finally, spatially resolved imaging and spectroscopy is needed to understand the internal physics of distant galaxies by measuring rotation and inflow in high redshift discs and providing the galaxy masses.

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On large spatial scales, one needs to push the search for clusters of galaxies beyond z ∼ 1.5 in order to constrain the basic cosmological parameters, in particular the dark energy equation of state. This is complementary to surveys with APEX and X-ray satellites, while ALMA will provide key follow-up. To understand the growth of large-scale structures, search for super-clusters at z > 0.5 must also be undertaken. At high redshift, z ∼ 5–7, although there is not yet full consensus, the evolution of the LBG LF appears to arise mostly from an increase in the characteristic luminosity M over cosmic time, with little evolution of the LF faint end slope and a modest decrease in the SF rate density [14]. This must be confirmed by additional extremely deep photometric surveys in the optical and near-infrared, since these results suggest that the lower luminosity galaxies play an important role in reionizing the Universe. Another open issue is whether the stronger clustering of bright (compared to faint) LBGs at z ∼ 4–5 [15] extends to higher z. For LAEs, there is no strong evidence for large-scale clustering at z ∼ 5.7 over ∼ 2 deg2 [6]. There is a deficit of brighter LAEs at z = 6.5 compared to z = 5.7. Although very model dependent, this may put some constraints on the neutral fraction of the Universe close to the end of the reionization epoch. Finally, exquisite spatial resolution is needed to determine the importance of merging at these early times.

3 Galaxy Evolution: Implied Studies and Tools Establishing the build-up of the present-day Hubble sequence of galaxies and constraining the later stages of galaxy evolution will require large area imaging surveys over several 102 deg2 . These surveys must cover the optical, including the U band, and the near-infrared to get accurate estimates of photometric redshifts. They should be fairly deep down to AB magnitudes of 25. They will also enable the discovery of rare, very bright and/or massive objects. There is a strong synergy with ALMA for obscured galaxies. Requirement: wide-field imagers covering all bands between U (mandatory) and H-K.

To determine the SF rate, stellar mass, mean stellar age and galaxy clustering at intermediate redshifts z ∼ 0.5–3, thus including the redshift desert, one must catch galaxies in all their evolutionary phases. One could then verify if scaling laws observed in the nearby Universe still hold at higher redshift. Spectroscopy at low resolution must be performed for samples of at least ∼ 105 galaxies at z > 0.5. This will also enable the detection of super-clusters of galaxies beyond z ∼ 0.5. Requirement: strong dedication of VIMOS to large spectroscopic surveys plus next generation of wide-field, multi-object spectrographs (optical and near-infrared).

The internal kinematics of z ∼ 1–3 galaxies provide galaxy total masses as well as signatures of galactic winds and star formation feedback. This requires 3D intermediate resolution spectroscopic surveys (multiple IFU) of fairly large samples,

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103 targets at z 2. For the brighter objects, one could then also derive the distribution of metals, dust and stars across the galaxies. These investigations have started with SINFONI (small field of view, thus single target). Requirement: second generation VLT instruments—KMOS, MUSE.

At higher redshifts, depth is crucial while sample size is also an issue. Deep imaging surveys in the optical and near-infrared are needed to constrain the SF rate at an epoch close to the end of reionization. Bright LBGs at z ∼ 6 with magnitudes zAB ≤ 25.4, and LAEs at z ∼ 5.7 with Lyα emission flux f ≥ 1.8 × 10−17 erg cm−2 s−1 are rare objects; the number density for both populations is n ∼ 50 deg−2 . That of faint LBGs at z ∼ 6 with zAB ≤ 28.1 reaches n ∼ 1 arcmin−2 . This implies surveys covering several deg2 down to (v, r, i)AB = 28, zAB = 26, KAB = 24. Requirement: wide-field imagers in the optical AND near-infrared.

Large-scale clustering studies require spectroscopic identification. At z > 5.5, this can only be achieved for LAEs and the brightest LBGs with 8–10 m telescopes. As these as rare objects, typically one per 70–100 arcmin2 , this will usually be single target spectroscopy. Requirement: second generation VLT instruments—X-shooter, MUSE.

Spatial resolution and morphological analysis at the 100 mas scale is achievable with current instruments as e.g. NACO on the VLT in the K-band, and the HSTACS in the i-band. Galaxies at z ∼ 6 are clearly resolved at this scale, but are too small (mean half-light diameter of 350 mas) for detailed morphological information, including the presence of merging blobs, with these imagers. Future facilities will address these questions, e.g. JWST will have a resolution of 35 mas at 1 µ. Requirement: VLT instruments—NACO, HAWK-I with AO; JWST; E-ELT.

4 High Redshift QSOs: Established Properties and Open Questions As mentioned above, there is a strong link between Active Galactic Nuclei (AGN) activity and galaxy formation. This is demonstrated by the correlation between black hole and bulge masses or that between black hole mass and stellar velocity dispersion. The sequence of activity AGN-starburst in the framework of galaxy evolution is still an open question. Nevertheless, feedback from the black hole during its accretion phase appears to play a major role in quenching star formation. The Sloan Digital Sky Survey (SDSS) has been essential for the discovery of an unprecedented large number of QSOs up to z = 6.42. The SDSS z ∼ 6 QSOs are a very rare population: at a survey magnitude limit of zAB = 20.2, or M1450 = −26.4, their number density is n 2.0 × 10−3 deg−2 [16]. The number density of QSOs with mAB < 20.2 decreases by ∼ 103 from 3 < z < 5 to z ∼ 6 and, at z > 3, their integrated luminosity function roughly varies as ρ ∝ 10−0.5z [17]. Even

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now, the faint end of the QSO LF has only be probed up to z ∼ 3.5. As at lower redshift, the QSO LF at z ∼ 3 is best fit by a double power law; the knee of the LF occurs at M1450 −26 and its faint end slope has a value of α −1.5 [18]. There is evidence that the LF bright end slope is shallower at z ∼ 4, β ∼ −2.5, than at z < 2.2, β ∼ −3.3 [17]. Does this flattening extend to even higher redshifts? This does not seem to be the case: from a SDSS deeper, smaller area survey down to zAB = 21, the value at z ∼ 6 is β ∼ −3.1, although still uncertain due to the small sample size [19, 20]. The number of faint QSOs at intermediate redshifts can also be derived from LBG surveys, since the Lyman-break selection technique uncover both the LBG and QSO populations. Spectroscopic identification differentiates these two classes of objects. At 2 z 4, the fraction of QSOs among the LBG samples (down to RAB = 25.5 at z ∼ 3) is constant and equal to 3% [1, 21, 22]. These QSO samples have indeed been considered in building the QSO LF at z ∼ 3 [18]. If the QSO fraction was unevolving with redshift, photometric LBG samples at z ∼ 5–6 could be used to estimate the knee and the faint end of the QSO LF. Alternatively, fainter QSO surveys could probe the LF knee. It is expected that the luminosity of the LF knee is decreasing with increasing redshift, as it is the case for the characteristic luminosity M of galaxies. The on-going z ∼ 6 faint QSO survey down to zAB = 23.9 [22] should at least provide an upper limit of the luminosity of the LF knee. At z ∼ 2–3, The UV luminosity density should be dominated by the ionizing radiation from QSOs and AGN. This is no longer the case at z ∼ 6, the galaxies being then the main source of ionizing radiation, implying a softening of the hard UV background flux with increasing redshift. A main issue is the number density of faint AGN at high redshift. In addition to the census of the sources contributing to the UV ionizing background, the latter can be constrained by the evolution of the ionization state of the Intergalactic Medium (IGM). This evolution can be studied from an analysis of high redshift QSO spectra. Two important parameters can be measured: the Gunn-Peterson optical depth in Lyα, Lyβ, Lyγ transitions and the distribution of the length of dark absorption gaps within the Lyman forest. This approach has been applied to the mall sample of 19 bright SDSS QSOs at z ∼ 6 [23]. The outcome is a fast evolution of the IGM physical state from z = 5.7 to z = 6.4: the ionization rate decreases by a factor over 4, the mean HI fraction increases by ∼ 10 and the mass averaged HI fraction is ∼ 1–4% at z ∼ 6.2.

5 High Redshift QSOs: Implied Studies and Tools To study the evolution of black holes over time and evaluate the contribution of the QSOs to the UV background, at least the bright end and the knee of the QSO LF should be better known at z > 4. This would also help distinguishing between a luminosity and/or a density evolution of the LF. The required surveys must cover large areas of several 102 deg2 down to (i, z)AB = 25, KAB = 23. For spatially unresolved objects at z > 5.5, there is a most severe contamination by brown dwarfs

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if the Lyman break, color selection technique is restricted to optical bands; observations must be extended to the near-infrared. As these are extremely rare objects, single target spectroscopy is needed for redshift confirmation. Requirement: wide-field imagers in the optical AND near-infrared, and second generation VLT instruments—X-shooter.

The physical state and metal enrichment level of the IGM towards the end of the reionization era requires QSO spectroscopy. The sample of high redshift QSOs should be large enough, 102 down to (i, z)AB = 23, to detect spatial inhomogeneities in the IGM ionization level and metallicity. Medium-low spectral resolution in the optical will give information on the IGM ionization level, while mediumhigh spectral resolution in the near-infrared is needed for metal line absorption studies. Requirement: second generation VLT instruments—X-shooter.

Kinematic studies of the stellar component of QSO host galaxies will be performed with 3D intermediate resolution spectroscopy in the near-infrared, and the mapping of Lyα emission from gaseous halos will require integral-field spectroscopy in the optical. Requirement: second generation VLT instruments—KMOS, MUSE.

The instrumentation requirements for morphological studies of QSO host galaxies are the same as those for normal galaxies, although with more extreme AO requirements. Only observations of QSO hosts at moderate redshift will be feasible with 8–10 m class ground-based telescopes. Requirement: VLT instruments—NACO, HAWK-I with AO; JWST; E-ELT.

6 Conclusions Third generation VLT instrumentation. As described in Sects. 3 and 5, the most needed new instruments are (1) Wide-field imagers (∼ 1 deg2 ) covering the optical, including the U-band, and the near-infrared, (2) a next generation Wide-field, high-multiplex spectrograph, together with an upgrade of VIMOS. For optical spectroscopy, trade-offs between covering the U- versus z-band must be investigated. Scientific priorities. From the scientific goals outlined in the previous sections, one can identify a few broad-scope projects, each of a few 103 hr: – establish the galaxy Hubble sequence together with the galaxy LF and SFR as a function of morphological type, mass and environment, – understand the internal physics of z ∼ 1–3 galaxies: dynamics, feedback and chemical composition, – explore the reionization epoch: search for faint QSOs and LAEs at z 6, and first statistical studies of the metal forest beyond z ∼ 5.5, – prepare ELT spectroscopic follow-up: large samples of faint, high-redshift candidate galaxies and QSOs.

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For the success of these large projects, the role of ESO in providing advanced data reduction tools is essential.

References 1. C.C. Steidel, K.L. Adelberger, M. Giavalisco, M. Dickinson, M. Pettini, Astrophys. J. 519, 1 (1999) 2. R. J Bouwens, G.D. Illingworth, J.P. Blakeslee, M. Franx, Astrophys. J. 653, 53 (2006) 3. R. J Bouwens, G.D. Illingworth, J.P. Blakeslee, T.J. Broadhurst, M. Franx, Astrophys. J. 611, L1 (2004) 4. K. Shimasaku, M. Ouchi, H. Furusawa et al., Publ. Astron. Soc. Japan 57, 447 (2005) 5. N. Kashikawa, K. Shimasaku, M.A. Malkan et al., Astrophys. J. 648, 7 (2006) 6. T. Murayama, Y. Taniguchi, N.Z. Scoville et al., Astrophys. J. Suppl. 172, 523 (2007) 7. M. Iye, K. Ota, N. Kashikawa et al., Nature 443, 1861 (2006) 8. J.-G. Cuby, P. Hibon, C. Lidman et al., Astron. Astrophys. 461, 911 (2007) 9. N.M. Förster Schreiber, R. Genzel, M.D. Lehnert et al., Astrophys. J. 645, 1062 (2006) 10. R. Genzel, L.J. Tacconi, F. Eisenhauer et al., Nature 442, 786 (2006) 11. M. Huertas-Company, D. Rouan, G. Soucail et al., Astron. Astrophys. 468, 937 (2007) 12. E. Daddi, M. Dickinson, R. Chary et al., Astrophys. J. 631, L13 (2005) 13. K.I. Caputi, H. Dole, G. Lagache et al., Astrophys. J. 637, 740 (2006) 14. R. J Bouwens, G.D. Illingworth, M. Franx, H. Ford, Astrophys. J. 670, 928 (2007) 15. N. Kashikawa, M. Yoshida, K. Shimasaku et al., Astrophys. J. 637, 631 (2006) 16. X. Fan, J.F. Hennawi, G.T. Richards et al., Astron. J. 128, 515 (2004) 17. G.T. Richards, M.A. Strauss, X. Fan et al., Astron. J. 131, 2766 (2006) 18. A. Bongiorno, G. Zamorani, I. Gavignaud et al., Astron. Astrophys. 472, 443 (2007) 19. L. Jiang, X. Fan, J. Annis et al., submitted to Astron. J. (2007), see arXiv:0708.2578 20. C.J. Willott, P. Delorme, A. Omont et al., Astron. J. 134, 2435 (2007) 21. C.C. Steidel, A.E. Shapley, M. Pettini et al., Astrophys. J. 604, 534 (2004) 22. K. Nandra, E.S. Laird, C.C. Steidel et al., Mon. Not. R. Astron. Soc. 360, L39 (2005) 23. X. Fan, M.A. Strauss, R.H. Becker et al., Astron. J. 132, 117 (2006)

Exoplanets: The Road to Earth Twins S. Udry, F. Pepe, C. Lovis, M. Mayor, the HARPS and ESPRESSO/CODEX Teams

Abstract Recent HARPS discoveries have demonstrated that very quiet stars exist, with intrinsic radial-velocity variations below 1 m s−1 . These results allow us in particular to characterize an emerging new population of very light planets down to a few Earth masses, as for example the components of a new 3-planet system, with masses between 3.6 and 8 M⊕ . Neptune-mass planets seem to be numerous. Their properties are especially important to constrain planet-formation models. Our HARPS experience also allows us to discuss the limitations of the radial-velocity method and the associated optimistic perspectives for the future detection of Earthlike planets in the Habitable Zone of solar-type stars, especially in the context of the foreseen development of ultra-stable spectrographs for the VLT (ESPRESSO) or the E-ELT (CODEX).

1 A Decade of Giant Planet Detections The discovery 13 years ago of an extra-solar planet orbiting the solar-type star 51 Peg [8] has encouraged the launch of numerous new search programs, leading now to a steadily increasing number of exoplanet detections. More than 270 other planetary companions have been found to orbit dwarfs of spectral types from F to M and more massive evolved stars. From this sample, we have learned that giant planets are common and that the planetary formation process may produce an unexpectedly large variety of configurations covering a wide range of planetary masses, orbital shapes, and planet-star separations (see e.g. [18]). The very large majority of the exoplanets have been found through the induced Doppler spectroscopic variations of the primary star (the so-called radial-velocity (RV) technique). Most of the candidates are giant gaseous planets similar in nature to Jupiter. With the development of a new generation of very stable spectrographs led by the HARPS spectrograph on the ESO 3.6-m telescope at La Silla, the past few years have known a new step forward in planet discoveries with the detections of lighter (5–20 M⊕ ), mainly “solid” planets (Table 1). The interest for very low mass planets (Neptune masses or super-Earths) follow several motivations: (i) the statistical properties of planetary systems provide constraints to the complex physical scenarios of planet formation, (ii) simulations of planetary formation also furnish information on the planet internal structure closely S. Udry () Geneva Observatory, Geneva University, 51 ch des Maillettes, 1290 Versoix, Switzerland e-mail: [email protected] A. Moorwood (ed.), Science with the VLT in the ELT Era, Astrophysics and Space Science Proceedings, doi: 10.1007/978-1-4020-9190-2_26, © Springer Science + Business Media B.V. 2009

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Table 1 Known very low mass planets (m2 sin i < 23 M⊕ ). Npl is the number of planets in the system. O–C are the velocity residuals around the orbital solutions Planet

P (d)

μ Ara

c 9.64

10.5

3.1

1.4

4

HARPS

[12, 15]

55 Cnc

e 2.81

14.2

6.7

5.4

5

HET

[9]

GJ 436

b 2.64

22.6

5.3

2?

HIRES

[4]

HD 190360

c 17.1

18.1

4.6

3.5

2

HIRES

[19]

GJ 876

d 1.94

5.9

6.3

4.6

3

HIRES

[14]

HD 4308

b 15.6

14.1

4.1

1.3

1

HARPS

[16]

HD 69830

b 8.67

10.2

3.5

0.6

3

HARPS

[6]

c 31.6

11.8

2.7

0.6

3

HARPS

[6]

d 197.

18.1

2.2

0.2

3

HARPS

[6]

b 5.4

15.7

12.5

1.2

3

HARPS

[2]

c 12.9

5.1

3.0

1.2

3

HARPS

[17]

GJ 581

m2 sin i (M⊕ )

K m/s

18.

O–C m/s

Npl

Instr.

Ref.

d 84.0

8.2

2.5

1.2

3

HARPS

[17]

HD 219828

b 3.83

19.8

7.0

1.7

2

HARPS

[10]

GJ 674

b 4.69

11.0

8.7

0.8

1

HARPS

[3]

New HARPS

b 4.31

3.6

1.8

1.2

3

HARPS

Mayor et al., in prep.

c 9.63

7.0

2.65

1.2

3

HARPS

Mayor et al., in prep.

d 20.5

8.0

2.3

1.2

3

HARPS

Mayor et al., in prep.

linked to the planet radius, and (iii) in a more distant future, space missions as DARWIN (ESA) or TPF (NASA) will search for life on terrestrial-type planets. Before the detailed design of such ambitious missions, we need a first insight on the frequency of terrestrial planets, and on the properties of their orbits. Planets in the Habitable Zone of our closest neighbors will be especially valuable. Here, we will concentrate on the mass distribution of exoplanets focusing mainly on its lowest-mass end, and on the prerequisites to the detection of Earth-type planets, in term of RV precision required improvements.

2 Planetary Mass Distribution The distribution of known planet masses is illustrated on Fig. 1 (left). The low-mass edge of the distribution is poorly defined because of observational incompleteness (smaller RV variations for the lower mass planets). However, in this planet-mass distribution, low-mass planets already start to draw a new population at very low masses. The emerging bimodal aspect of the distribution strongly suggests that the decrease of the distribution for masses less than about one mass of Jupiter is not only the result of the detection bias but should be real, and provides an interesting

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Fig. 1 Left. Planetary mass distribution from giant planets to super-Earths. The double-hatched histogram represents HARPS detections. Right. Histogram of RV rms for stars in the high-precision HARPS subprogram. Part of the “higher”-rms tail results from stellar activity or from still undetected planetary systems. A colour version of this figure is available at dx.doi.org/10.1007/978-1-4020-9190-2_26

constraint for planetary formation scenarios. Because of their small masses and central locations in the systems, these low mass planets are probably mainly composed of rocky/icy material. The discovery of very low mass planets close to the detection threshold of RV surveys, over a short period of time, suggests that this kind of objects are rather common. Moreover, at larger separations (2–3 AU), the microlensing technique is finding similar mass objects (the lightest with a mass of 5.5 M⊕ , [1]) showing that smaller mass planets can be found over a large range of separations. This is in complete agreement with the latest Monte Carlo simulations of planet formation, that furthermore predict a large population of still lighter Earth-like planets [11].

3 HARPS High-Precision Program The global histogram of the observed RV dispersions for the stars in the HARPS high-precision GTO program (G- and K-dwarf targets) presents a mode at 1.4 m s−1 (Fig. 1, right). The velocities have been obtained over several seasons from 2003 to 2007. Part of the observed rms in the tail of the distribution, at the level of 2 to 3 m s−1 , can be explained by stars with still some chromospheric activity. Also, several multi-planetary systems have been detected with a global RV rms at the level of 3–4 m s−1 (before fitting the orbital solution). A good example of the latter case is provided by the “Trio of Neptunes” (HD 69830; [6]). The rough 4 m s−1 rms of the observed velocities drop to residuals of 0.6 m s−1 around the 3-planet Keplerian solution, and even down to 0.2 m s−1 around the 6-months 3rd planet (removing the contributions of the 2 shorter-period

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Fig. 2 Phase-folded orbital solutions (left) and top view (right) of the 3 planets orbiting the next HARPS planetary system to be announced (Mayor et al., in prep.; Table 1). Planet masses are between 3.6 and 8 M⊕ . A colour version of this figure is available at dx.doi.org/10.1007/978-1-4020-9190-2_26

planets and considering run-averaged velocities). For the global HARPS highprecision sample, we can thus reasonably suspect that part of the observed RV scatter is really the result of undetected, low-amplitude, multi-planetary systems. Actually, a close analysis of these data provides hints for the existence of many more very low-mass planetary candidates. Among them, an exciting 3-planet system to be published soon, with masses between 3.6 and 8 M⊕ , is presented for illustration in Fig. 2 (orbital parameters are given in Table 1).

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4 Limitations for Precise Radial-Velocity Measurements The planetary mass estimated from the Doppler measurements is directly proportional to the amplitude of the reflex motion of the primary star. The measure of very precise RVs requires that all steps along the light path, from the star to the detector, are well understood. The main aspects are discussed e.g. in [13], we just recall here a few important points. Assuming that the target spectrum is not contaminated by external sources (e.g. stellar neighbors), the interstellar medium, or the Earth atmosphere, the main limitations are basically divided in 3 categories: (1) The stellar noise groups error sources which are produced at the emission of the light, i.e. by the observed source itself: stellar pulsations, surface granulation and activity-related jitter. The cross-correlation technique is so efficient that, for most of the stars in the HARPS high-precision sample, the photon noise is at the level of 0.5 m s−1 after an exposure time smaller than the typical periods of stellar acoustic modes. Long integrations (15 minutes) are sufficient to damp these RV variations below 0.2 m s−1 (rms). At longer variation time scales, Kjeldsen et al. [5] suggest that granulation can induce RV variability larger than (or comparable to) 1 m s−1 for solar-type stars. To damp the granulation noise several measurements spanning a few hours could be required. Test observations with HARPS are ongoing to better characterize this point. Finally, any anisotropies of the stellar atmosphere will also induce RV variations with time scales comparable to the stellar rotation period. The amplitude of this RV jitter is correlated with stellar chromospheric activity. The reemission in the core of the calcium lines is an efficient indicator to select a sample of “non-active” stars. (2) Instrumental errors are those related to the detection process. They include the whole light path starting at the telescope and ending on the detector. Among these instrumental effects we can distinguish two main contributions: errors affecting the measurements (stability or repeatability) supposing a perfect wavelength reference, and calibration errors (i.e. errors on the wavelength scale). Lovis and Pepe [7] have considerably improved the precision of the wavelength of thorium lines as well as the number of lines to be used for the calibration of the spectrograph. The performances demonstrated by HARPS have excited the imagination of astronomers. New ultra-stable spectrographs are now studied for large telescopes, with the aim of reaching long-term RV precisions down to 10 cm s−1 (ESPRESSO/VLT; Pasquini et al., this volume) or even at the level of a few cm s−1 (CODEX/E-ELT). (3) Photon noise finally sets the fundamental limit for the attainable precision as the latter scales with the signal-to-noise of the spectra. To get to ultra-high precisions, a huge number of photons is needed. [email protected] ESO telescope belongs to the most efficient RV spectrographs; it reaches a precision of about 1 m s−1 in less than one minute, on a late G dwarf of mv = 7.5. Reaching 1 cm−1 would then require an exposure time of 10 000 minutes. This makes evident that larger telescopes are urgently needed. To achieve a 10 cm−1 precision on a mv = 8 star, a 10-m class telescope is required. The situation becomes even more dramatic if we consider that typical transit candidates delivered by space-based surveys (COROT, Kepler), carry

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magnitudes typically above mv = 12. If 10 cm−1 precision had to be achieved in a single exposure, a 50-m telescope would then be ideally needed. It is quite difficult to correctly estimate the real amplitude of the different sources of noise intrinsic to the star or due to the instrument. The final precision will depend on which of these factors actually dominates the error budget. However, we can set an upper limit for the quadratic sum of these noises (for non-active stars) at less than 1 m s−1 . This limit is estimated from the O–C around the orbital solutions of planets detected with HARPS (Table 1) or from the lowest RV dispersion of HARPS measurements (Fig. 1, right).

5 Searching for Earth-Type Planets in the Habitable Zone Is it possible to detect terrestrial planets in the Habitable Zone (HZ) of neighboring stars? For planets orbiting M dwarfs this is already feasible. The detection of two low mass planets (5.1 M⊕ and 8 M⊕ ) at both edges of the HZ of the M4V star GJ 581 [17] is a good example of that possibility. For Earth-type planets orbiting solar-type stars, the situation is obviously more challenging. About an order of magnitude in the precision of the measured RVs has to be gained. We have seen that perspectives for such an improvement are good. In particular, the ESPRESSO@VLT project presently studied to be implemented at the incoherent combined focus of the VLT’s, and later CODEX@E-ELT, will provide the required efficient and ultra-stable instrumentation to reach this challenging goal. The unambiguous identification of the signature of an Earth-like planet from Doppler measurements will require a large number of observations (to beat activityrelated effects), and the set up of an adequate observing strategy (to diminish the influence of acoustic modes and granulation). A search for Earth-like planets around a sample of a couple of tens of stars will then be possible but will be expensive in term of telescope time.

6 Concluding Remarks The future detection of very low mass planets and possibly Earth twins requires precision at the cm−1 -level. Opposite to what was thought 3 to 6 years ago, HARPS has been able to demonstrate that sub-m s−1 precision can be reached. Although stellar noise may be one of the limiting factors, there seem to be many stars that allow us to reach even 20 to 30 cm−1 rms on short-term. How quiet the most stable stars are, and over which time scales, needs to be investigated in more detail. On the instrumental side, many progresses have been made. In particular, from the HARPS experience, we have not identified any show stopper until date, and we think that a final precision of better than 5–10 cm−1 on the most stable stars is within reach. The development of extremely stable spectrographs on large telescopes will provide us with high instrumental precision and high efficiency. Despite the different

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Fig. 3 Masses of detected exoplanets as a function of the year of discovery. The planets orbiting a neutron star (1992) or detected by microlensing (∼ 2005) are represented by squares, transiting planets and the most numerous exoplanets discovered by Doppler spectroscopy by circles having radii proportional to their orbital eccentricities. The lower envelope illustrates the continuous progresses of the spectrograph sensitivity. Since a few years we have entered the era of super-Earth detections, rushing towards Earth masses. A colour version of this figure is available at dx.doi.org/10.1007/978-1-4020-9190-2_26

sources of noise, we are confident that Doppler spectroscopy will be able to detect rocky planets in the Habitable Zone of solar-type stars well before the launch of DARWIN- or TPF-type space satellites, and thus provide the prerequisite targets for these enthusiastic missions.

References 1. 2. 3. 4. 5. 6. 7. 8. 9. 10. 11.

J.-P. Beaulieu, D.P. Bennett, P. Fouqué et al., Nature 439, 437 (2006) X. Bonfils, T. Forveille, X. Delfosse et al., Astron. Astrophys. 443, L15 (2005) X. Bonfils, M. Mayor, X. Delfosse et al., Astron. Astrophys. 474, 293 (2007) R.P. Butler, S.S. Vogt, G.W. Marcy et al., Astrophys. J. 617, 580 (2004) H. Kjeldsen, T. Bedding, R.P. Butler et al., Astrophys. J. 635, 1281 (2005) C. Lovis, M. Mayor, W. Benz et al., Nature 441, 305 (2006) C. Lovis, F. Pepe, Astron. Astrophys. 468, 1115 (2007) M. Mayor, D. Queloz, Nature 378, 355 (1995) B. McArthur, M. Endl, W. Cochran et al., Astrophys. J. 614, L81 (2004) C. Melo, N.C. Santos, W. Gieren et al., Astron. Astrophys. 467, 721 (2007) C. Mordasini, Y. Alibert, W. Benz, D. Naef, Giant planet formation by core accretion, in Extreme Solar Systems, ed. by D. Fischer, F. Rasio, S. Thorsett, A. Wolsczcan. ASP Conf. Ser. (2008), in press 12. F. Pepe, A.C. Correia, M. Mayor et al., Astron. Astrophys. 462, 769 (2007) 13. F. Pepe, C. Lovis, From HARPS to CODEX: Exploring the limits of Doppler measurements, in: Physics of Planetary Systems, ed. by N. Piskunov et al. Nobel Symposium, vol. 135 (Swedish Royal Academy of Sciences, 2008), in press 14. E.J. Rivera, J.J. Lissauer, R.P. Butler et al., Astrophys. J. 634, 625 (2005)

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Next Generation Deep Redshift Surveys with the VLT Olivier Le Fèvre

Abstract The next generation of deep redshift surveys will need to considerably expand redshift coverage and survey volume to satisfy the need for robust and unbiased statistical measurements of the evolutionary properties of galaxies and large scale structure. To maintain the leading edge of the European community in this field will require the development of new instruments: in particular a ‘NIRMOS-like’ uncooled multi-slit spectrograph expanding the high multiplex ( 200) capabilities to the near infrared, and a ‘MegaMOS’ offering an order of magnitude increase in field of view and multiplex (∼ 10 000) in the visible. Next generation surveys will also require large allocations of time over several years. A change in VLT time allocation policy will be necessary to enable continued support from ESO and the community.

1 Deep Redshift Surveys at the VLT Today 1.1 Historical Perspective Conducting deep redshift surveys within ESO countries started with surveys like ESP with the 3.6 m ESO telescope [1], and the Canada-France Redshift Survey conducted on the CFHT [2, 3]. Other surveys contributed significantly to the field with surveys like the LDSS [4] and DEEP [5]. These redshift surveys were the key observational progress necessary to firmly establish the strong evolution in the life of galaxies since z 1 [6]. The current generation of redshift surveys is based on a new generation of highly efficient multi-object spectrographs on the large 8 to 10 m telescopes. The ESO community has been very active in shapping up these new opportunities. Today, deep spectroscopic surveys at the VLT are at the heart of all studies up to the highest possible redshifts, covering 0 < z < 7. O. Le Fèvre () Laboratoire d’Astrophysique de Marseille, OAMP, CNRS & Université de Provence, BP8, 13376 Marseille, France e-mail: [email protected] A. Moorwood (ed.), Science with the VLT in the ELT Era, Astrophysics and Space Science Proceedings, doi: 10.1007/978-1-4020-9190-2_27, © Springer Science + Business Media B.V. 2009

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1.2 Probing the Universe with Spectroscopic Redshift Surveys Spectroscopic redshift surveys enable redshift measurements of distant sources with accuracies of order 100 to a few hundred km/s. The 3D galaxy distribution field can be build and spectral features can be measured accurately. Deep redshift surveys provide the material to study the formation and evolution of galaxies and large scale structures, and to place constrains on the cosmological world model. The drive is to probe all epochs towards reionization (say z ∼ 10–15), with a large enough number of objects and observed volume to enable a robust measurement of statistical descriptors of the population (luminosity function, star formation rate, merger rate, . . . ) and minimize cosmic variance effects. Access to the distribution of galaxies on large scales 100 Mpc is needed to probe the evolution of the large scale structures, connected to the underlying dark matter distribution. As they are probing large volumes of space, deep surveys are also providing valuable input to constrain parameters in the world model, e.g. from the growth rate of structures [7] or from the evolution of reference scales like baryonic acoustic oscillations [8], complementary to other cosmological probes.

1.3 Instrumentation for Surveys at the VLT At the VLT, a number of instruments can be used today to conduct deep redshift surveys. In the visible domain, VIMOS is a multi-slit survey spectrograph covering 14 × 16 arcmin2 , for up to ∼ 1000 spectra at a time [9]. FORS2 on UT2 offers a multi-slit capability in a 7 × 7 arcmin2 field, with excellent red sensitivity complementary to VIMOS. FLAMES-GIRAFFE is a multi-fiber instrument with up to 130 objects observed in a field 30 arcmin in diameter, with also the capability to observe 15 galaxies in 3D with small IFUs [10]. In the near infrared, there are no survey instruments per say, but SINFONI and ISAAC can be used for single object spectroscopy in the JH and K bands, offering integral field capability with SINFONI [11]. In development, KMOS will expand the capabilities in the IR with 24 arms capable of integral field spectroscopy. MUSE will probe the Universe with a large 1 arcmin2 IFU fed with adaptive optics.

1.4 On-going Large Redshift Surveys at the VLT Several ‘very wide’ to ‘very deep’ surveys are being carried out at the VLT. The wide VIMOS VLT Deep Survey (VVDS) is currently covering 8 deg2 [12], and aiming for 30 deg2 down to IAB = 22.5. The wide zCOSMOS survey covers 2 deg2 at a 2.5× higher spectral resolution than VVDS [13]. A medium-deep survey on 1 deg2 of the UKIDSS-UDS combining VIMOS and FORS observations is in progress (Almaini et al.). Going deeper, VVDS-Deep [14], zCOSMOS-deep [13], and GOODS [15]

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offer complementary observational strategies to probe the high redshift population 1 < z < 5. The deepest surveys are the GMASS with FORS2 [16], and the VVDSUltraDeep (on-going with VIMOS). Taking the example of the VVDS, a total of ∼ 45 000 spectra have been obtained in the 3 survey regimes limited in apparent magnitude to IAB ≤ 22.5, 24, 24.75. A total of 25 000 redshifts and spectra have been made publicly available (http://cencosw.oamp.fr). With as few as possible a priori selection accurate measurements of the properties of galaxies with 0 < z ≤ 5 have been obtained [14], just citing a few: luminosity function [17], star formation rate [18], mass function and stellar mass density [19, 20], correlation function [21, 22], and constraints on cosmology parameters [7]. A new view of the high redshift 1.5 < z < 5 galaxy population has also emerged [23]. More than 40 papers have been published so far, including two in Nature, all coming from a total of less than 35 garanteed nights in 2001–2002. As a follow-up for some of these surveys, one should mention the integral field 3D surveys with Flames-Giraffe and SINFONI, providing insight on the internal dynamics of individual galaxies. Existing or on-going surveys can then be roughly summarized as follows: (i) at z ∼ 1: surveys cover areas of ∼ 10 deg2 , with samples of a few 104 galaxy; with 105 samples hopefully coming soon; (ii) at z ∼ 2–3: surveys with several hundred galaxies in areas of ∼ 1 deg2 are available; and samples of a few 103 galaxies are being observed; (iii) at z ∼ 4–6 a few tens of redshifts have been measured Samples at z ∼ 1 are still 20–100 times smaller than in large local 2dFGRS and SDSS surveys ( 1 million redshifts). The push to larger high redshift samples therefore remains very strong.

2 Next Generation Surveys: Test Cases for the VLT 2.1 SDSS-like Survey at z 1 Large samples in large volumes enable to study galaxy properties globally but also for selected sub-populations (e.g. selected by luminosity, type, density), as demonstrated by local redshift surveys. To be equivalent to the SDSS at z ∼ 1 would require to observe galaxies with a luminosity below L∗ , and to cover scales on order 200 Mpc. What can the VLT and current instrumentation do? To cover 100 deg2 to a magnitude AB ∼ 22.5 with VIMOS would require about 300 nights, with 7 × 106 galaxies observed. This is therefore feasible now. Upgrading VIMOS with red sensitive detectors would improve the efficiency by about 25%.

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2.2 VVDS-like Survey at z 3–6 The epoch 3 < z < 6 is most likely a key epoch in the build-up of galaxies. Current knowledge is based on only several hundred galaxies with confirmed spectroscopic redshifts. The next step is to see how these galaxies are distributed in large scale structures on scales up to 100 Mpc, including up to 105 redshifts for a statistically robust study of sub-populations. A survey of ∼ 105 galaxies in 5 deg2 would require about 300 nights of VLTVIMOS, therefore feasible with existing instrumentation. The addition of near-IR multi-object spectroscopy with an uncooled instrument of the NIRMOS type [25] working in the zY and J bands would be very powerful, following key tracers of star formation activity like Hα to z > 1.

2.3 All-Sky Survey at z 1 The ‘ultimate’ redshift survey is to aim for a complete measurement of the 3D distribution of galaxies over the whole sky. Trying to cover the redshift range up to z ∼ 2 at IAB = 23 with a sampling of 1 in 3 galaxies with the VLT would require about 500 years with current instrumentation. One certainly needs to think of new instruments with higher multiplex! The VLT can certainly compete in this game. A ‘MegaMOS’ spectrograh on the VLT with a field > 1 deg2 , and ∼ 10 000 objects observed simultaneously would allow to observe ∼ 108 galaxies. Feasibility studies making use of new technology (micro-mirror arrays, high density fiber feed, . . . ) are therefore advisable. This would have to be put in perspective of the future space projects, like the SPACE mission proposed to ESA Cosmic-Vision program [24].

2.4 Ultra-deep Probe to z 10 Going to even higher redshifts to z ∼ 10 and getting a few hundred objects to understand what they really are may seem impossible with ground-based 8 m telescopes. Here it is hard to beat a space facility with the infrared background about three orders of magnitude lower than from the ground. JWST and dedicated missions like SPACE [24] will be the premier facilities in this game. VLT, and to a much larger extent the EELT will remain very efficient for integral field observations. Concepts with multiple IFUs will certainly play an important role if the multiplex gain is kept high.

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3 Implications for Next Generation Instruments and VLT Operations Model There are two limiting factors for new surveys on the VLT. One is the lack of a near-IR multi-object spectrograph with a large field and large multiplex. NIRMOS was selected as early as 1996, but technical and managerial reasons led ESO to cancel it, unfortunately without a possibility to capitalize on this study to define a world leading facility. However, a simple and efficient instrument concept optimized for Y and J bands, remains a front runner for many galaxy evolution applications. In another front, the need for very wide areal coverage is pushing first generation instruments to their limits, and it is clearly time to develop ‘MegaMOS’ instrument concepts for very wide field ∼ 1 deg2 , capable to obtain ∼ 104 redshifts. These two instruments will be the key elements in a new strategy. A fundamental support to large surveys is the strong commitment of an organization to carry out a few large programs over many years. In its current way of dealing with large programs, ESO is de facto preventing very large > 100 nights, multi-year programs to get accepted or to get completed. The example of VIMOS is enlightening: it has not been possible to date to complete the highest ranked science program which was the main motivation to build VIMOS: a survey of the Universe to z ∼ 1 with > 105 redshifts. For the VLT and ESO to remain competitive on large survey science will therefore require a change of paradigm to enable the large allocation of observing time required by deep redshift surveys.

References 1. 2. 3. 4. 5. 6. 7. 8. 9. 10. 11. 12. 13. 14. 15. 16. 17. 18. 19. 20. 21.

G. Vettolani, E. Zucca, G. Zamorani et al., Astron. Astrophys. 325, 954 (1997) S.J. Lilly et al., Astrophys. J. 455, 50 (1995) O. Le Fèvre et al., Astrophys. J. 455, 60 (1995) M. Colless, R.S. Ellis, K. Taylor, R.N. Hook, Mon. Not. R. Astron. Soc. 244, 408 (1990) D.C. Koo, R. Guzman, S.M. Faber et al., Astrophys. J. 440, L49 (1995) S.J. Lilly, O. Le Fèvre, F. Hammer, D. Crampton, Astrophys. J. 460, L1 (1996) L. Guzzo et al., Nature, in press D.J. Eisenstein, I. Zehavi, D.W. Hogg et al., Astrophys. J. 633, 560 (2005) O. Le Fèvre, M. Saisse, D. Mancini et al., SPIE 4841, 1670 (2003) H. Flores et al., Astron. Astrophys. 455, 107 (2006) Henri Bonnet, Robert Abuter, Andrew Baker et al., Messenger 117, 17 (2004) B. Garilli et al., in preparation S.J. Lilly, O. Le Fèvre, A. Renzini et al., Astrophys. J. Suppl. Ser. 172, 70 (2007) O. Le Fèvre, G. Vettolani, B. Garilli et al., Astron. Astrophys. 439, 845 (2005) E. Vanzella, S. Cristiani, M. Dickinson et al., Astron. Astrophys. 434, 53 (2005) C. Halliday et al., Astron. Astrophys. (2008), in press, arXiv:0801.1193 O. Ilbert, L. Tresse, E. Zucca et al., Astron. Astrophys. 439, 863 (2005) L. Tresse, O. Ilbert, E. Zucca et al., Astron. Astrophys. 472, 403 (2007) L. Pozzetti, M. Bolzonella, F. Lamareille et al., Astron. Astrophys. 474, 443 (2007) S. Arnouts, C.J. Walcher, O. Le Fèvre et al., Astron. Astrophys. 476, 137 (2007) O. Le Fèvre, L. Guzzo, B. Meneux et al., Astron. Astrophys. 439, 877 (2005)

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O. Le Fèvre B. Meneux, O. Le Fèvre, L. Guzzo et al., Astron. Astrophys. 452, 387 (2006) O. Le Fèvre, S. Paltani, S. Arnouts et al., Nature 437, 519 (2005) M. Robberto et al., Proc. Venice Conf. (August 2007), arXiv:0710.3970 Olivier Le Fèvre, Michel Saisse, Dario Mancini et al., SPIE 4008, 546 (2000)

GUAIX: The UCM Group of Extragalactic Astrophysics and Astronomical Instrumentation J. Gallego, N. Cardiel, S. Pascual, M.C. Eliche-Moral, A. Castillo-Morales, R. Guzmán, A. Gil de Paz, P.G. Pérez-González, J. Gorgas, J. Zamorano and GUAIX Team

1 Introduction The front-ends of astronomical research are linked to new telescope improvements, like the building of 10 m-class (and probably larger in a near future) telescopes, and instruments with increasingly larger data gathering capabilities. In the current situation is quite common that astronomers are faced with a huge amount of data to handle in order to extract the useful information. Our group at the UCM have developed expertise in the design and execution of large scientific projects using the newest technology in 10 m class telescopes (e.g., GOYA, http://www.astro.ufl.edu/GOYA/; OTELO, http://www.iac.es/project/otelo/), and in the development of state-of-the-art software tools for data reduction and analysis. With the aim of taking advantage of the expertise acquired in both fronts, we have formed a new group of extragalactic astrophysics and astronomical instrumentation, GUAIX (http://guaix.fis.ucm.es), which aims at exploiting the synergy of such combination.

2 Pipelines for GTC Instruments At present GUAIX is involved in developing the data reduction pipelines for three instruments for the Spanish 10 m GTC (Gran Telescopio Canarias, http://www.gtc. iac.es), namely EMIR, FRIDA and SIDE (see short descriptions below). One of the main goals in the development of those pipelines is to provide the astronomers with reduced data ready to be employed in their scientific analysis. This, which is the normal situation with data from space observatories, is not so common in ground-based observations, where astronomers are used to self-reduce their own data once back at their home institutions. In order to convince the reluctant astronomers about the reliability of the reduced products provided by the pipelines, it is critical to include a J. Gallego () Departamento de Astrofísica y C.C. de la Atmósfera, Facultad de C.C. Físicas, Universidad Complutense de Madrid, Avda. Complutense s/n, 28040 Madrid, Spain e-mail: [email protected] A. Moorwood (ed.), Science with the VLT in the ELT Era, Astrophysics and Space Science Proceedings, doi: 10.1007/978-1-4020-9190-2_28, © Springer Science + Business Media B.V. 2009

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proper handling of errors (and their propagation) in the data treatment, which apart from facilitating the necessary quality control, provides the critical information to constrain the reliability of the scientific measurements (see [1, 2]). An immediate benefit of the use of instrument-tuned pipelines is that the reduced data can be incorporated into databases with the guarantee that they have been homogeneously derived. GUAIX is directly involved in the development of the data reduction pipelines for three instruments of GTC. Here we briefly describe the main characteristics of those instruments. EMIR (http://www.ucm.es/info/emir) is a wide-field, near-infrared, multi-object spectrograph proposed for the Nasmyth focus of GTC. It will allow observers to obtain up to ∼ fifty intermediate resolution spectra simultaneously in one of the nIR bands Z, J, H, K. A configurable slit unit will be used for target acquisition. EMIR is designed to address the science goals of the proposing team (GOYA project) and of the GTC community at large. EMIR is being built by a consortium of Spanish and French institutions, led by the Instituto de Astrofísica de Canarias. The PI is Dr. F. Garzón. FRIDA (http://www.astroscu.unam.mx/ia_cu/proyectos/frida) is an near-IR camera and spectrograph for the adaptive optics focus of GTC. For integral field spectroscopy, it will use an image slicer. FRIDA will provide spectral resolutions of roughly 1000, 4000, and, uniquely for a diffraction-limited integral-field spectrograph, 30 000. FRIDA is being constructed by the Instituto de Astronomía of the Universidad Nacional Autónoma de México (UNAM), the Instituto de Astrofísica de Canarias, the University of Florida, the Departamento de Astrofísica of the UCM, and the Laboratoire Astrophysique de Toulouse. The PI is Dr. J.A. López. SIDE (Super Ifu Deployable Experiment, http://side.iaa.es) is currently being proposed as a third generation GTC instrument. It is a multi-purpose middle resolution fiber fed spectrograph dedicated to wide field MOS and three-dimensional spectroscopy. It will operate in the visible and near-IR light. This instrument is managed by a consortium of institutions from Spain, México and Florida, headed by the Instituto de Astrofísica de Andalucía (IAA, CSIC). The PI is Dr. F. Prada.

References 1. N. Cardiel, J. Gorgas, J. Gallego, Á. Serrano, J. Zamorano, M.L. García-Vargas, P. GómezCambronero, J.M. Filgueira, Proc. SPIE 4847, 297 (2002) 2. N. Cardiel, J. Gorgas, J. Gallego, Á. Serrano, J. Zamorano, M.L. García-Vargas, P. GómezCambronero, J.M. Filgueira, Rev. Mex. Astron. Astrofís., Ser. Conf. 16, 73 (2003)

VISTA Public Surveys and VLT followup Will Sutherland

The 4-metre VISTA telescope and its 67 Mpixel wide-field near-infrared camera are now nearly complete at Paranal, awaiting delivery of the primary mirror in early 2008. I provide an overview of the six Public Surveys which have been selected by ESO to use most of the VISTA time for the first five years. There is good synergy between the VISTA surveys and the second-generation VLT instruments HAWK-I, X-shooter and KMOS; in the longer term, enhancement of VLT multi-object spectroscopic capability is highly desirable.

1 VISTA System Overview 1.1 Introduction VISTA is a 4-metre telescope optimised for wide-field surveys, sited at Paranal Observatory on its own peak approximately 1.5 km North of the VLTs. It is being constructed by a UK consortium led by UKATC, Edinburgh: when complete it will be the world’s largest dedicated survey telescope, and will become an ESO facility as part of the UK-ESO joining agreement. For at least its first five years it will use a single instrument, the near-infrared camera.

1.2 VISTA Telescope The VISTA telescope [1, 2] has an unusual design in which the telescope essentially forms the fore-optics to a very large near-IR Camera. The fast 4 m f/1 primary mirror and 1.24 m secondary mirror form a modified Ritchey-Chretien system, which is jointly optimised with the IR Camera 3-lens corrector, delivering an f/3.25 focus at Cassegrain. The fast design results in a very compact and rigid telescope structure, a wide Cassegrain field, a large instrument payload and a small enclosure: minor downsides are the tight alignment tolerances between the two mirrors, and the fact that only one instrument can be mounted at a given time. The telescope structure Will Sutherland () Astronomy Unit, Queen Mary Univ. London, London E1 4NS, UK e-mail: [email protected] A. Moorwood (ed.), Science with the VLT in the ELT Era, Astrophysics and Space Science Proceedings, doi: 10.1007/978-1-4020-9190-2_29, © Springer Science + Business Media B.V. 2009

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Fig. 1 Front view of the VISTA telescope. The IR Camera vacuum window is seen at centre, protruding through the hole in the dummy primary mirror. A colour version of this figure is available at dx.doi.org/10.1007/978-1-4020-9190-2_29

(Fig. 1) is an altazimuth mount constructed by Vertex-RSI of Texas, and includes an 84-point active support system for the primary mirror. A high-precision hexapod (NTE, Barcelona) provides 5-axis control of the secondary mirror position for focus and collimation. Both mirrors are polished by LZOS of Moscow, with the secondary now delivered to Paranal and the primary expected in early 2008.

1.3 IR Camera The VISTA IR Camera attaches to the Cassegrain rotator on the back of the primary mirror cell. At 3 metres long, 2.9 tonnes and 67 Mpixels it is the world’s largest near-IR imager [3]. The Camera’s field of view is 1.65 degrees diameter, which is sparse-filled with sixteen Raytheon VIRGO 2k × 2k near-IR detectors (total 67 Megapixels) sensitive from 0.8–2.4 µm. The detectors give 0.34 arcsec pixel scale, thus 0.6 deg2 (2150 arcmin2 ) instantaneous field of view. The Camera has no cold stop, but instead contains a 2 m long cylindrical cold-baffle to block the detectors’ view of ambient-temperature surfaces. This leads to a long cryostat nose protruding through the central hole in the primary, and a very large front vacuum window of 95 cm diameter (Fig. 1). Towards the rear of the cryostat are a 3-lens wide-field corrector, an 8-position filter wheel, two fixed autoguiders and wavefront sensors, and the IR detector system.

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The camera is designed for long-term low-maintenance operations, with 3 closed-cycle coolers and only one internal moving mechanism (the filter wheel). The high QE of the detectors, the silver-coated mirrors and few all-Infrasil transmitting elements (1 window + 3 lenses) lead to an excellent overall system throughput.

1.4 Status At present the enclosure, telescope structure, secondary mirror and IR Camera are essentially complete at Paranal: the primary mirror is still in the final stages of polishing and is on the critical path. Delivery of the primary is expected in March 2008, with science verification approximately 6 months thereafter. Once operational, VISTA’s wide field, high efficiency and single-instrument operation should make it the world’s leading near-IR survey facility, probably until a space mission such as the proposed DUNE or SNAP after 2016.

2 The VISTA Public Surveys Following an iterative process of proposals, selection and merging, the ESO Public Survey Panel has selected six Public Surveys which will share approximately 80% of VISTA’s observing time for the first five years of operation. Each survey will have a series of quality-controlled data releases to the full ESO community. The sky coverage of the six surveys is shown in Fig. 2. Briefly, the six surveys are as follows:

Fig. 2 Approximate sky coverage of the six planned VISTA surveys (see labels), overlaid on a full-sky image from 2MASS. A colour version of this figure is available at dx.doi.org/10.1007/978-1-4020-9190-2_29

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• VHS, the VISTA Hemisphere Survey: this will cover the entire Southern hemisphere in J, Ks bands to Ks,AB ≈ 20.0, with H-band and increased depth in the 5000 deg2 overlap with the planned Dark Energy Survey. • VVV, VISTA Variables in Via Lactea. This is a multi-epoch survey of the Galactic Bulge and part of the plane: among many goals, RR Lyrae will be used to map the detailed 3D structure of the inner Galaxy. • VMC, the VISTA Magellanic Cloud survey, will survey both Magellanic Clouds and the Bridge to approximately the main-sequence turnoff depth, to provide a comprehensive view of stellar populations and the star-formation history. • VIKING, the VISTA Kilo-degree Infrared Galaxy survey. This is an intermediatedepth Ks,AB ≈ 21.2 extragalactic survey of 1500 deg2 in 5 bands, matching the VST-KIDS visible survey to provide a 9-band combined survey. Primary science goals include z > 6 quasars, 3D weak lensing using photometric redshifts, and dark energy. • VIDEO, VISTA Deep Extragalactic Observations. This will survey 12 deg2 in 5 bands to Ks,AB ≈ 23.5 within three Spitzer SWIRE fields. The main science goals are the evolution of galaxies, clusters and AGNs at z ∼ 1 to 3, and photometric redshifts for Spitzer, Herschel, submm and radio sources. • UltraVISTA will survey 0.75 deg2 within the HST-COSMOS field in Y,J, H,Ks to a very deep limit, Ks,AB ≈ 25.5. This will study galaxy mass assembly to z ∼ 4, and high-luminosity Ly-break galaxies at z ∼ 6–8. A narrowband 1.18 µm filter will also be used to search for Ly-α emitters at z ≈ 8.7. More details of these planned VISTA surveys, and also VST visible surveys, are provided in [4] and the ESO website.

3 VLT Followup In addition to the ESO public surveys from VISTA and VST, on a similar timescale there will be numerous other planned imaging surveys including X-ray (E-Rosita), visible (Dark Energy Survey), mid-IR (warm Spitzer, WISE), far-IR (Akari), submm (APEX, Herschel, SCUBA2) and radio (LOFAR, SKA pathfinders). Thus, we can look forward to major advances in wide-area multi-wavelength imaging in the next 5 years, but our capabilities for spectroscopic followup are advancing more slowly: hence classifications in multicolour space and photometric redshifts are likely to play an increasing role, but spectroscopic subsamples will remain essential for many purposes. The VISTA and VST public surveys will probably give rise to numerous requests for VLT followup in three general categories: 1. Deeper and higher resolution followup imaging on selected targets. 2. Spectroscopic confirmation of new rare objects, including high-z quasars, ultracool brown dwarfs, gravitational lenses, etc. Due to low surface densities, these usually require single-target spectroscopy.

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3. Multi-object spectroscopy of large samples, preselected in various ways from multi-wavelength imaging data. Training sets and verification of visible + VISTA photometric redshifts is likely to be one major driver. It is notable that many of the expected VISTA discoveries, including z > 6 quasars and galaxies, z ∼ 1 clusters, z ∼ 1–2 red galaxies (EROs, DRGs and BzKs, etc.), and ultracool brown dwarfs, have key spectral features in the observed wavelength range 0.8–1.3 µm. This provides very good synergy with FORS2 and HAWK-I for followup deeper imaging, X-Shooter for single-target spectroscopy, and KMOS for clustered targets. However for spectroscopy of large samples the VLT instrumentation is less optimal, since VIMOS has low efficiency beyond 0.8 µm, while KMOS will be efficient but its field and multiplex are modest. In the near term, an upgrade or copy of either VIMOS or FLAMES to cover the 0.8–1.3 µm range appears very promising, and is likely to be achievable using ambient-temperature optics at a moderate cost. In the longer term in the era of JWST and ELTs, the VLT will no longer be at the forefront in sensitivity or resolution, but a further enhancement to wide-field spectroscopy can retain a leading role. Next-generation dark energy and Galactic archaeology surveys [5] will probably require a field > 1 degree which cannot be achieved at existing VLT foci, but appears practical using a new VLT Prime-focus corrector with fibre-fed spectrographs and a multiplex ∼ 2500 objects [6]. As in the lively discussion at the meeting, there are clearly challenges making such an instrument fit the available space envelope and retaining backward-compatibility for other instruments and VLTI, but the science payoff is potentially very large and this prospect deserves more detailed study. In summary, VISTA, VST and many other facilities will provide a wealth of wide-field imaging data in the near future, so multi-object spectroscopy is likely to remain among the most important roles in the future of VLT.

References 1. J.P. Emerson, W. Sutherland, A. McPherson et al., ESO Messenger 117, 27 (2004) 2. J.P. Emerson, A. McPherson, W. Sutherland, ESO Messenger 126, 41 (2006) 3. G.B. Dalton, M. Caldwell, K. Ward et al., in: Ground-based Instrumentation for Astronomy, ed. by A. Moorwood, M. Iye. Proc. SPIE, vol. 5492 (2004) pp. 988–997 4. M. Arnaboldi, M.J. Neeser, L.J. Parker et al., ESO Messenger 127, 28 (2007) 5. M. Irwin, in Science with the VLT in the ELT Era, ed. by Alan Moorwood. Astrophysics and Space Science Proceedings (Springer, Netherlands, 2009), p. 193 6. I. Parry, in Science with the VLT in the ELT Era, ed. by Alan Moorwood. Astrophysics and Space Science Proceedings (Springer, Netherlands, 2009), p. 417

Probing Dark Energy with Cosmological Redshift Surveys at the VLT L. Guzzo and the VVDS Consortium

Large redshift surveys of galaxies play a key role in the quest for the origin of cosmic acceleration. In this context, it has recently been pointed out [1] that redshift-space distortions due to galaxy peculiar motions can be used to trace the growth rate of structure f (z) back in time. Coupled to estimates of the expansion history H (z) as those provided by Type Ia supernovae, this can distinguish models with a truly extra “dark energy” component from theories in which the acceleration is explained by modifying the laws of gravity. Current measurements are not accurate enough to distinguish between these two alternatives. According to extended simulations, this will become possible with an order of magnitude increase in the sampled volume and number of redshifts at similar depth, i.e. measuring > 100 000 galaxies over

Including also: M. Pierleoni (MPA Garching) B. Meneux (MPE Garching) E. Branchini (Univ. Roma 3) O. Le Fèvre (LAM Marseille), C. Marinoni (Univ. Marseille), B. Garilli (INAF-IASF Milano), G. Blaizot (Univ. Lyon), G. De Lucia (MPA Garching), A. Pollo (Univ. Warsaw), H.J. McCracken (IAP Paris), D. Bottini (INAF-IASF Milano), V. Le Brun (LAM Marseille), D. Maccagni (INAF-IASF Milano), R. Scaramella (INAF-OA Roma), M. Scodeggio (INAF-IASF Milano), L. Tresse (LAM Marseille), G. Vettolani (INAF-IRA Bologna), A. Zanichelli (INAF-IRA Bologna), C. Adami (LAM Marseille), S. Arnouts (LAM Marseille), S. Bardelli (INAF-OA Bologna), M. Bolzonella (INAF-OA Bologna), M. Bondi (INAF-IRA Bologna), A. Bongiorno (Univ. Bologna), A. Cappi (INAF-OA Bologna), S. Charlot (IAP Paris), P. Ciliegi (INAF-OA Bologna), T. Contini (Toulouse), O. Cucciati (INAF-OA Brera), S. De la Torre (LAM Marseille), K. Dolag (MPA Garching), S. Foucaud (Nottingham), P. Franzetti (INAF-IASF Milano), I. Gavignaud (AIP Potsdam), L. Guzzo (INAF-OA Brera), O. Ilbert (IfA Hawaii), A. Iovino (INAF-OA Brera), F. Lamareille (INAF-OA Bologna), B. Marano (Univ. Bologna), A. Mazure (LAM Marseille), P. Memeo (INAF-IASF Milano), R. Merighi (INAF-OA Bologna), L. Moscardini (Univ. Bologna), S. Paltani (ISDC Geneva), R. Pellò (Toulouse), E. Perez-Montero (Toulouse), L. Pozzetti (INAF-OA Bologna), M. Radovich (INAF-OA Napoli), D. Vergani (INAF-IASF Milano), G. Zamorani (INAF-OA Bologna), E. Zucca (INAF-OA Bologna). L. Guzzo () INAF-Osservatorio Astronomico di Brera, Via Brera 28, 20021 Milan, Italy e-mail: [email protected] L. Guzzo Max Planck Institut für Extraterrestrische Physik, 85748 Garching, Germany L. Guzzo Excellence Cluster “Universe”, 85748 Garching, Germany A. Moorwood (ed.), Science with the VLT in the ELT Era, Astrophysics and Space Science Proceedings, doi: 10.1007/978-1-4020-9190-2_30, © Springer Science + Business Media B.V. 2009

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nearly 108 h−3 Mpc3 . The VLT with VIMOS can play an important role in this endeavour, filling the gap between current surveys and projects foreseen beyond year 2015, based on new ground-based or space-born mega-MOS spectrographs.

1 Cosmic Acceleration and Growth of Structure Observations indicate that we live in a low-density, expanding Universe with spatially flat geometry, that, quite surprisingly, appears to have recently entered a phase of accelerated expansion. This latter conclusion emerges naturally when interpreting the observed Hubble diagram of distant Type Ia supernovae within the standard Friedmann-Lemaitre-Robertson-Walker (FLRW) cosmology [2, 3]. Formally, observations are well reproduced by adding a cosmological constant in the equations of General Relativity (GR), i.e. the term originally introduced by Einstein to obtain a static solution. This has a few disturbing features, and scenarios with evolving “dark energy” density have been proposed, as e.g. quintessence (see e.g. [4] for a review). These variants correspond to modifying the right-hand side of Einstein field equations, i.e. adding an extra contribution in the stress-energy tensor. Alternatively, one can however assume that it is the theory of gravity that needs to be revised and thus modify the left-hand side of the equation. This would imply that the “observed” acceleration is just a cosmic mirage, simply evidencing our still limited knowledge of the laws of Nature (see [5] for a comprehensive review of these variants). These two options cannot be distinguished by looking only at the expansion history H (z) or equivalently at the equation of state w(z) of the extra component (w = −1 being the cosmological constant case). The linear growth of density inhomogeneities provides a way to break the degeneracy. This can be described in the linear regime by the growth rate f = d ln D/d ln a, where D(t) is the growing mode of fluctuations and a = (1 + z)−1 is the cosmic scale factor. A simple form f (z) [Ωm (z)]γ gives an accurate description for a wide range of models [8, 9], with γ depending on the gravity theory (e.g. [10]). f (z) is sensitive to the physics responsible for the cosmic acceleration: scenarios with the same expansion history H (z), but based on a different gravity theory will predict a different f (z) and γ . We have recently shown [1] that measurements of redshift-space distortions at different cosmic epochs represent a very promising way to estimate f (z). Galaxy peculiar motions are a direct consequence of the growth of structure. When redshifts are used to measure galaxy distances, the contribution from peculiar velocities introduces a measurable distortion in the clustering pattern, which is proportional to f (z). Such an anisotropy can be quantified by means of the redshift-space two-point correlation function ξ(rp , π). The compression of ξ(rp , π) at large rp ’s is proportional to the parameter β = f/bL (see [11] for a review), i.e. to the growth rate, modulo the bias factor bL (the ratio of the clustering amplitude of the galaxies to that of the matter). Locally (z 0.15) the 2dFGRS has measured β = 0.49 ± 0.09 [12] for galaxies with bL = 1.0 ± 0.1 [13]. This provides an important local constraint, corresponding to f (z = 0.15) = 0.49 ± 0.14.

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2 Measuring f (z) at z ∼ 1 with the VVDS-Wide Survey The VIMOS-VLT Deep Survey (VVDS) was designed to probe the combined evolution of galaxies and large scale structure to z ∼ 2 using the VIMOS spectrograph at the ESO VLT. It measured so far ∼ 40 000 spectra [14, 15] over its Deep (0.5 deg2 to IAB < 24) and Wide (∼ 8 deg2 to IAB < 22.5) parts. We have used a sub-sample of 5895 galaxies with 0.6 < z < 1.2 (volume V = 6.35 × 106 h−3 Mpc3 ) in the VVDSWide F22 4 deg2 field to measure β at an effective redshift z = 0.77 [1]. We fit ξ(rp , π) with a full distortion model including both linear and non-linear distortions, characterised respectively by β and by the rms pairwise dispersion σ12 . Figure 1 shows ξ(rp , π) estimated using standard methods [16], with superimposed the bestfit model contours, corresponding to β = 0.70 ± 0.26 and σ12 = 412 ± 70 km s−1 . Error bars are obtained from 100 fully realistic mock realisations of the survey [17, 18]. Using the amplitude of mass fluctuations provided by the power spectrum of CMB anisotropies [19], we obtain an effective bias bL = 1.3 ± 0.1 (see [1] for details), and a growth rate f (z = 0.77) = 0.91 ± 0.36. This value is compared in Fig. 1 to model predictions, together with measurements from the 2dFGRS [12] and 2SLAQ surveys [20]. Given the size of the error bars, deviations [5] from the standard cosmological-constant model cannot yet be detected. Interestingly, how-

Fig. 1 Left panel: ξ(rp , π) at z 0.77 from the VVDS-Wide survey, replicated over four quadrants to enhance deviations from circular symmetry. Colours correspond to the level of correlation as a function of the transverse (rp ) and radial (π ) separation of galaxy pairs. The effect of galaxy infall due to the growth of large-scale structure is proportional to the flattening of the purple-blue large-scale levels, with the solid contours corresponding to the best-fitting distortion model with β = 0.70 and σ12 = 412 km s−1 (see Ref. [1] for details). Right panel: Estimates of the growth rate f = βbL compared to predictions from theoretical models: the standard cosmological constant (CDM) model (w = −1) (solid line); an open Ω = 0 model with the same Ωm (long-dashed line, for both cases f (z) Ω(z)0.55 ); two models in which dark energy is coupled to dark matter [6] (upper dashed curves); the DGP braneworld model, an extra-dimensional modification of the gravitation theory [7] for which f (z) Ω(z)0.68 (dot-dashed curve). A colour version of this figure is available at dx.doi.org/10.1007/978-1-4020-9190-2_30

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ever, measurements tend to disfavour an Ω = 0 model, providing an independent indication for the need of a cosmological constant (or modified gravity).

3 Future Prospects for Cosmological Redshift Surveys at the VLT Only ∼ 6000 redshifts have been used to obtain the result discussed above: there are clearly exciting prospects for this technique to become a primary method to study cosmic acceleration. We have used our large set of simulations to predict the gain in accuracy that can be expected from future surveys. We have found that, for a given bias factor, the measurement error on β depends on the mean density n and the volume of the survey V , as σβ ∝ (n0.44 V 0.5 )−1 (i.e. scaling nearly as the inverse square root of the total number of objects), as shown in Fig. 2. We are not accounting here for expected improvements in the distortion model and clustering estimator [17].

Fig. 2 Dependence on the survey size of the rms error on the distortion parameter β, for the VVDS-Wide F22 data [1] and future surveys (IAB = 22.5, ∼ 1-hour exposure, 1 VIMOS pass). The dashed band corresponds to a ±10% error around the fiducial value. Predictions are based on realistic simulations [17] and present two options where a pre-selection of galaxies at z > 0.6 is applied. Using VVDS redshifts and the accurate CFHTLS photometry [21] we have shown that a > 95% complete sample can be constructed using colour- or photometric-redshift pre-selection, doubling the effective number density of measured galaxies. A colour version of this figure is available at dx.doi.org/10.1007/978-1-4020-9190-2_30

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As shown, going below a 10% uncertainty on β is within reach of current VLT instrumentation. A 500-hrs survey with VIMOS could be accommodated within the current framework of ESO Large Programs without impacting normal activity, provided that it can be spread over 8 semesters. Such a program would assemble an unprecedented survey of > 100 000 redshifts at 0.5 < z < 1.2 over 32 deg2 . This would be the basis not only for a key measurement of the growth rate, but also for a number of front-ranked researches that have just been sketched by current surveys, due to lack of statistics and volume at these redshifts. As a by-product, these redshifts will also be precious for calibrating photometric redshifts from the huge imaging surveys that are just starting, as e.g. Pan-STARRS. A more extreme option to measure nearly 1.5 million redshifts over 400 deg2 is also shown. Such a survey would push the error on β below 3% and would be possible only by dedicating one of the VLT UT’s to survey work. Using VIMOS, this would need around 5000– 6000 hours. This could be significantly reduced with a new MOS spectrograph with larger multiplexing (3000–4000) and field of view (∼ 1 deg2 or more) (see e.g. contributions by O. Le Fèvre and B. Nichol). This would allow even larger areas to be covered, such that one would not only trace accurately f (z) within the survey itself, but also measure complementary probes of H (z), as Baryonic Acoustic Oscillations. The question is thus whether the ESO community will be willing, in the ELT era, to invest in having one 8-m “redshift machine” out of the four VLT units. The lesson to be learned from the unique results obtained by the previous-generation 4-m AAT telescope when it was dedicated to surveys with the 2dF facility, is that such a decision would certainly turn out to be a breakthrough investment for cosmology in the next decade.

References 1. 2. 3. 4. 5. 6. 7. 8. 9. 10. 11. 12. 13. 14. 15. 16. 17. 18. 19. 20. 21.

L. Guzzo, M. Pierleoni, B. Meneux et al., Nature 451, 541 (2008) A.G. Riess et al., Astron. J. 116, 1009 (1998) S. Perlmutter et al., Astrophys. J. 517, 565 (1999) M.S. Turner, D.J. Huterer, Phys. Soc. Jpn. 76, 111015 (2007) E.J. Copeland, M. Sami, S. Tsujikawa, preprint, hep-th/0603057 L. Amendola, Mon. Not. R. Astron. Soc. 312, 521–530 (2000) G. Dvali, G. Gabadadze, M. Porrati, Phys. Lett. B 485, 208–214 (2000) E.V. Linder, Phys. Rev. D 72, 043529 (2005) L. Wang, P.J. Steinhardt, Astrophys. J. 508, 483 (1998) R. Lue, R. Scoccimarro, G.D. Starkman, Phys. Rev. D 69, 124015 (2004) A.J.S. Hamilton, in: The Evolving Universe. ASSL, vol. 231 (Kluwer, 1998), p. 185 E. Hawkins et al., Mon. Not. R. Astron. Soc. 346, 78 (2003) L. Verde et al., Mon. Not. R. Astron. Soc. 335, 432 (2002) O. Le Fèvre et al., Astron. Astrophys. 439, 845 (2005) B. Garilli et al., submitted to Astron. Astrophys. (2008) S.D. Landy, A.S. Szalay, Astrophys. J. 412, 64 (1993) M. Pierleoni et al. (2008), in preparation G. De Lucia, J. Blaizot, Mon. Not. R. Astron. Soc. 375, 2 (2006) D.N. Spergel et al., Astrophys. J. Supp. 170, 377 (2007) N.P. Ross et al., Mon. Not. R. Astron. Soc. 381, 573 O. Ilbert et al., Astron. Astrophys. 457, 841 (2006)

The First Galaxies and Galaxy Clusters Eelco van Kampen

Obvious targets for Extremely Large Telescopes are high-redshift galaxies and galaxy (proto-)clusters. However, due to the extremely small field-of-view of such telescopes, finding such objects, notably the rarer proto-clusters, is best be done using other facilities. The very helpful negative K-correction experienced in the far-IR to mm wavebands makes surveys in these wavebands the ideal pathfinders for high-z ELT targets.

1 Introduction The shape of the peak of the dust-emission of many bright, high-z starforming galaxies is such that the observed flux is roughly the same in the redshift interval 1 < z < 10, depending on the exact band [1]. This means that once the sensitivity of a (sub-)mm telescope is good enough to detect these sources at z ≈ 1, one is able to detect all such sources down to their formation epoch, as this is not likely to be beyond z ≈ 10. We use simulations to test how effective a search for proto-clusters, which can then be observed using an ELT, is in the (sub-)mm.

2 Simulation Methods We use an updated version of the phenomenological galaxy formation model of van Kampen, Jimenez and Peacock [2], which is semi-numerical, in the sense that the merging history of galaxy haloes is taken directly from N -body simulations. When dark matter haloes merge, a recipe based on dynamical friction is used to decide how many galaxies exist in the newly merged halo. The most massive of those galaxies becomes the new central galaxy to which gas can cool, while the others become its satellites, with no cooling. The rate at which this hot gas cools and is able to form stars is given by the standard radiative cooling curve. Various feedback mechanisms reheats some of the cooled gas back to the hot phase. The simulation code maintains an internal account of the amounts of gas being transferred between the E. van Kampen () Institute for Astro and Particle Physics, Univ. of Innsbruck, Technikerstr. 25/8, 6020 Innsbruck, Austria e-mail: [email protected] A. Moorwood (ed.), Science with the VLT in the ELT Era, Astrophysics and Space Science Proceedings, doi: 10.1007/978-1-4020-9190-2_31, © Springer Science + Business Media B.V. 2009

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Fig. 1 Mock sub-mm survey with a cluster inserted at z = 2.5, in the 850 micron waveband for a 0.5 × 0.5 deg. survey at the James Clark Maxwell Telescope. The (proto-)cluster is clearly visible in the centre, as a still collapsing overdensity that will become a relaxed cluster at about z ≈ 0.8. A colour version of this figure is available at dx.doi.org/10.1007/978-1-4020-9190-2_31

two phases, and the amount consumed by the formation of stars. The cosmological model adopted is Ωm = 0.3, Ω = 0.7, h = 0.7, σ8 = 0.93, Ωb = 0.02 h−2 . We include two modes of star formation: quiescent star formation in disks, following the Schmidt law with a threshold according to the Kennicutt criterion, and star-bursts during major merger events, which can be either galaxy-galaxy mergers or halo-halo mergers. The sub-mm flux from the dust component is calculated using GRASIL [3], which is a code to compute the spectral evolution of stellar systems taking into account the effects of dust, which absorbs and scatters optical and UV photons and emits in the IR-submm region.

3 An Example Survey and Suggested Follow-up Using the phenomenological galaxy formation model described in Sect. 2 we produced mock observations of proto-clusters in the 850 micron sub-mm waveband. The cluster clearly stands out, especially at the highest flux levels, and provides an excellent set of targets for follow-up observations. Many surveys of order 10 sq. de-

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grees or larger are being planned, and these would be good pathfinders for studying high-z galaxies and galaxy clusters. First of all one could get low-resolution ‘spectra’ for large samples though combining data from SCUBA-2 on the JCMT, PACS and SPIRE on the Herschel Space Observaory, AzTEC, LABOCA, BLAST, and the LMT. Then targets can be found using infra-red IDs from UKIDSS or Ultra-VISTA, for example, or radio IDs from the VLA or the GMRT. We can then obtain high-resolution images in the sub-mm using the (e)SMA or ALMA and in the optical using the ELTs, high-resolution spectroscopy from ALMA, JWST, or the LMT, and redshifts and masses from FMOS (Subaru), KMOS (VLT), for example.

4 Summary The main point made is that the (sub-)mm waveband is a blurry but deep probe of the high-redshift Universe, which means that upcoming (sub-)mm surveys are excellent target finders for ELTs. Future work will focus on how to best select possible ELT targets for a given survey. Acknowledgements This research was supported in part by the Austrian Science Foundation FWF under grant P18493, and the European Commission’s Research Infrastructures activity of the Stucturing the European Research Area programme, contract number RII3-CT-2003-506079 (HPC-Europa).

References 1. A.W. Blain, M.S. Longair, Mon. Not. R. Astron. Soc. 264, 509 (1993) 2. E. van Kampen, R. Jimenez, J.A. Peacock, Mon. Not. R. Astron. Soc. 310, 43 (1999) 3. L. Silva, G.L. Granato, A. Bressan, L. Danese, Astrophys. J. 509, 103 (1998)

Narrow Band Surveys and the Epoch of Reionization B.P. Venemans, R.G. McMahon, I.R. Parry, D.J. King, J. Bland-Hawthorn and A.J. Horton

In the last couple of years tremendous progress has been made in the study of high redshift (z > 4) galaxies. Red-sensitive wide field optical cameras equipped with narrow band filters discovered emission line galaxies up to z ∼ 7, reaching the boundary of optical CCDs. Here we report on the results of a search for z = 7.7 emission line galaxies using DAZLE, a wide field near-infrared VLT visitor instrument, designed to image between the bright night sky emission lines that dominate the sky background at 1.0–1.8 µm.

1 Introduction How, when and over what time scale, galaxies formed are questions at the forefront of work in both observational and theoretical cosmology. In recent years the observational horizon has expanded rapidly and radically for those observing distant galaxies. Large format red-sensitive detectors on wide field imaging instruments, the new generation of 8 m class telescopes and the refurbished Hubble Space Telescope (HST), have pushed the limits to which we can routinely detect star-forming distant galaxies progressively from redshifts of one to beyond z = 6. At the highest redshifts currently accessible, narrow band emission lines searches using the Lyman-α line have pushed from redshifts of 4 [2] progressively to 5.7 [3, 5] and 6.5 [4, 7, 8]. The highest redshift galaxy discovered in a narrow band survey has a redshift of z = 6.96 [6], reaching the boundary of silicon based optical technology.

2 Narrow Band Surveys in the Infrared The first attempts to discover z > 7 galaxies using narrow band Lyman-α imaging on infrared cameras were unsuccessful [1, 10], mainly due to the small field of view of the instrument used (VLT/ISAAC, FOV < 10 arcmin2 ) and the brightness of the B.P. Venemans () Institute of Astronomy, Madingley Road, Cambridge, CB3 0HA, United Kingdom e-mail: [email protected] A. Moorwood (ed.), Science with the VLT in the ELT Era, Astrophysics and Space Science Proceedings, doi: 10.1007/978-1-4020-9190-2_32, © Springer Science + Business Media B.V. 2009

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sky background. In the near-infrared, the broad band sky is 20–50 times brighter than in the optical. In a seminal paper, Maihara [9] showed that at a resolution of R = 17 000 the OH lines are unresolved and moreover between the OH airglow the background sky was one-fiftieth the average flux in the H band. To capitalise on this dark background, one needs to observe the sky at high spectral resolution, e.g. R ∼ 1000 whereupon the strong OH lines render 2% of the spectral region useless but the remaining 98% samples the dark sky. The VLT visitor instrument Dark Age ‘Z’ Lyman-α Explorer (DAZLE) is a wide field (6.8 × 6.8 arcmin2 FOV) infrared camera designed to image between the night sky emission lines. Prior to the DAZLE project it was considered impossible to manufacture large interference filters with a resolution of 1000. As part of the DAZLE design study in collaboration with the Anglo Australian Observatory and Barr Associates we were able to demonstrate that practical filters could be manufactured. By observing in a R ∼ 1000 band, one obtains a two-fold sensitivity gain from working at a spectral resolution that maximises the line to continuum contrast and also minimises the effect of the OH sky emission.

3 First Science Run with DAZLE on the VLT In November 2006 we conducted the first science run with DAZLE on the VLT. Integration of DAZLE onto the Nasmyth platform of Melipal (VLT UT3) was completed successfully prior to the start of the scheduled commissioning nights on October 30th and 31st 2006. Two narrow band filters with a FWHM of 9 Å were installed in DAZLE, one centred at 1.056 µm and one at 1.063 µm, allowing us to search for Lyman-α emission line galaxies at z ∼ 7.7. After the successful commissioning of DAZLE on the VLT, our science programme was carried out primarily on the 9 nights from November 2nd to 10th 2006, accumulating a total on-sky integration time of 69 hrs excluding time spent on calibration. The measured seeing in our images during the run ranged from 0.4 arcsec to 1.3 arcsec. We imaged two deep survey pointings with an exposure time of around 10 hrs per filter. In addition we imaged 6 shallow survey pointings when the seeing was poor (0.8–1.3 arcsec), with an average exposure time of 2.5 hrs per filter. The preliminary on-sky measured sensitivity of DAZLE which includes detector dark current, read out noise, instrument and sky background is roughly 3–5 × 10−18 erg s−1 cm−2 in 10 hours for a source spread over 25 pixels, i.e. 1 arcsec2 . A dark corrected, flat fielded image in the 1.056 µm filter of the GOODSSouth field can be seen in Fig. 1. We are reasonably confident that we were seeing the ‘true’ sky background because we could see rings of very marginally higher background due to expected faint OH lines encroaching on the wings of the filter transmission profile. We also saw that the DAZLE background varied with lunar phase and distance. The brightness of the sky between the OH lines as measured from our images can be as dark as ∼ 19.6 AB mag arcsec−2 , although variations by factors > 3 were found. This sky brightness corresponds to

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Fig. 1 DAZLE image of the GOODS-South field in the 1.056 µm filter. The rings of very marginally higher background due to expected faint OH lines encroaching on the wings of the filter. A colour version of this figure is available at dx.doi.org/10.1007/978-1-4020-9190-2_32

∼ 750 photons s−1 m−2 arcsec−2 µm−1 at 1.06 µm, which can be compared to the Maihara number of 590 photons s−1 m−2 arcsec−2 µm−1 at 1.67 µm. The data obtained in November 2006 is currently being reduced and analysis of the data is in progress. Note that observing in two adjacent narrow band filters allows us to reject various types of objects, such as transient objects (supernovae and moving bodies) and objects with extreme colours (e.g. EROs) that contaminate single narrow band searches for high redshift galaxies [1]. Some preliminary results of our DAZLE run can be seen in Figs. 2 and 3. In Fig. 2 we plot a “colour-magnitude” diagram for objects detected in the 1.056 µm filter image of the GOODS-South field. Objects with an emission line in the 1.056 µm (1.063 µm) filter will have a positive (negative) flux excess. One such object is shown in Fig. 3. The object is detected in the 1.063 µm filter, but absent in the 1.056 µm image. Based on the colours in the COMBO-17 survey, the galaxy has a photometric redshift of zp = 0.606 [11]. Very likely this object is a galaxy at z = 0.62 with the H-α line falling into the 1.063 µm filter.

4 Future Surveys and Instruments In the next few years, much progress is expected in the field of narrow band surveys for high redshift galaxies. Subaru with the wide field imager Suprime-Cam will find many galaxies at redshifts z < 7, while new infrared imagers like VLT/HAWK-I, VLT/DAZLE and VISTA will likely discover the first galaxies in the OH clean regions at 1.06 µm and 1.19 µm corresponding to Lyman-α redshifts of z = 7.7

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Fig. 2 “Colour-magnitude” diagram showing the flux difference in the two DAZLE filters as a function of 1.056 µm flux for objects in the GOODS-South field. The dashed lines show the flux difference for an object with an equivalent width of 10 Å and the solid curves show the significance of the flux difference

and 8.8. In the more distant future, JWST and the ELTs will be able to search for high redshift line emitters using infrared IFUs or tunable filters. Although these telescopes will have high angular resolution, the field of view will likely be modest: a few arcmin on a side. An area where the VLT could contribute is wide field infrared narrow band imaging. In principle the existing DAZLE instrument is capable of searching for galaxies in the redshift range z = 9–12. Note that the gaps in the OH airglow in the H band are narrower than in the J band, meaning that the narrow R = 1000 filter DAZLE approach is essential for these redshifts. A natural upper limit for wide field narrow band imaging with the VLT is the maximum field size accessible from the VLT Nasmyth focus, which is ∼ 28 arcmin in diameter. A 2 × 2 mosaic of IR detectors would give a contiguous field of view of 20.5 × 20.5 arcmin2 with a sampling of 0.3 arcsec pixel−1 . Such an instrument (“DAZLE-2”) would have a survey power 9

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Fig. 3 DAZLE images in the 1.056 µm filter (left) and 1.063 µm filter (right) of an emission line object discovered in the GOODS-South field. The size of the images is roughly 20 arcsec. The object has a photometric redshift of zp = 0.606 [11], so we identify the galaxy as an H-α emitter at z = 0.62. A colour version of this figure is available at dx.doi.org/10.1007/978-1-4020-9190-2_32

times that of the existing DAZLE instrument. In the future such visitor instruments could become increasingly important for the VLT. In this way the VLT would be able to play an important role in the era of ELTs.

References 1. 2. 3. 4. 5. 6. 7. 8. 9. 10. 11.

J.-G. Cuby, P. Hibon, C. Lidman et al., Astron. Astrophys. 461, 911 (2007) E. Hu, R. McMahon, Nature 382, 231 (1996) E. Hu, R. McMahon, L. Cowie, Astrophys. J. 522, L9 (1999) E. Hu, L. Cowie, R. McMahon et al., Astrophys. J. 568, L75 (2002) E. Hu, L. Cowie, P. Capak et al., Astron. J. 127, 563 (2004) M. Iye, K. Ota, N. Kashikawa et al., Nature 443, 186 (2006) K. Kodaira, Y. Taniguchi, N. Kashikawa et al., Publ. Astron. Soc. Jpn. 55, L17 (2003) J. Kurk, A. Cimatti, S. di Serego Alighieri et al., Astron. Astrophys. 422, L13 (2004) T. Maihara, F. Iwamuro, T. Yamashita et al., Publ. Astron. Soc. Pac. 105, 940 (1993) J. Willes, F. Courbin, Mon. Not. R. Astron. Soc. 357, 1348 (2005) C. Wolf, K. Meisenheimer, M. Kleinheinrich et al., Astron. Astrophys. 412, 913 (2004)

Stellar Archaeology and Galaxy Genesis: The Need for Large Area Multi-Object Spectrograph on 8 m-Class Telescopes Mike J. Irwin and Geraint F. Lewis

The origin and evolution of galaxies like the Milky Way and M31 remain among the key questions in astrophysics. The galaxies we see today in and around the Local Group are representatives of the general field population of the Universe and have been evolving for the majority of cosmic time. As our nearest neighbour systems they can be studied in far more detail than their distant counterparts and hence provide our best hope for understanding star formation and prototypical galaxy evolution over the lifetime of the Universe [1]. Significant observational progress has been made, but we are still a long way from understanding galaxy genesis. To unravel this formative epoch, detailed large area multi-object spectroscopy of spatial, kinematic and chemical structures on 8 m-class telescopes are required, to provide the link between local near-field cosmology and predictions from the high-redshift Universe.

1 Overview Two of the key challenges facing modern astrophysics are our understanding of the fundamental nature and properties of dark matter, and of the inter-related complex baryonic processes that led to the formation of large galaxies. The large-scale gravity-driven properties of dark matter guide the gaseous matter into proto-galactic environments, but the dominance of baryonic physics on small scales, through cooling, collapse and star formation, emphasises that understanding only the nature of dark matter provides a very limited view of galaxy formation and evolution. However, both of these aspects can be simultaneously addressed with large scale spectroscopic surveys of the stellar components of the Milky Way and companion galaxies. In addition to providing tests of cosmological models on low-mass scales, such surveys will also unravel the complex process of hierarchical formation wrapped within various Galactic structures. These studies directly compliment proposed major studies of dark energy, and taken as a whole such efforts with large area M.J. Irwin () Institute of Astronomy, Madingley Road, Cambridge, CB3 0HA, UK e-mail: [email protected] A. Moorwood (ed.), Science with the VLT in the ELT Era, Astrophysics and Space Science Proceedings, doi: 10.1007/978-1-4020-9190-2_33, © Springer Science + Business Media B.V. 2009

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multi-object spectrographs are key to providing a new window on the three major energy constituents of the universe and their influence on the Cosmos we see around us.

2 Observational Context Although CDM cosmologies demonstrate the ubiquity of hierarchichal merging as the main driver in galaxy formation and evolution, the discovery of the tidally disrupting Sagittarius dwarf [2] provided the first compelling nearby evidence of this. However, the detailed process by which large galaxies such as the Milky Way arrive at their current state is still largely speculative (e.g. [3–5]) despite recent observational and theoretical progress. Current large scale spectroscopic surveys such as RAVE [6], and the M31 kinematic surveys on Keck [7, 8] highlight the potential of multi-object spectroscopy to probe the structure and properties of the Galaxy and M31 at modest resolution, while FLAMES on the VLT provides a compelling argument for the multiplexed use of mid- to high-resolution spectroscopy, particularly in understanding nearby dwarf galaxies [9]. However, FLAMES, although producing exquisite spectra of hundreds of objects, is at least an order of magnitude shy of the requirements needed for analysis of the Galaxy and M31. All of these efforts are necessarily limited in either coverage or depth. While future proposed Galactic surveys such as Segue-II [10] and the ESA cornerstone space mission GAIA [11] will go some way to alleviate this problem for the Galaxy, they will still leave unexplored a huge volume of parameter space that holds the key to full understanding of Galactic evolution. Tle “field of streams” in Belokurov et al. [12] convincingly demonstrates the complexity of even the halo of the Milky Way as seen in large area photometric surveys. Several giant streams (up to 60 degrees long) from disrupting globular clusters and satellite dwarf galaxies criss-cross the halo providing probes of both the nature of dark matter and the gravitational potential of the Galaxy.

3 The Way Forward The diversity of current and planned surveys illustrates the scientific importance and interest in this subject and 8 m-class large area (1◦ –2◦ diameter) multi-object (1000– 4000) spectrographs would provide a unique opportunity for a major advance in our understanding of the origin, chemical evolution and mass assembly of nearby galaxies. The key advantages of such facilities over existing 8 m-class systems would be the greater multiplexing and wavelength coverage together with the much larger field-of-view. This combination potentially gives almost 2 orders of magnitude improvement over existing facilities and would provide enormous leverage in Galactic

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studies by enabling a range of structural probes in a multi-parameter space sensitive to both dynamical and chemical properties. Galaxies are dynamically and chemically inhomogeneous [9] and it is only by using the higher dimensionality afforded by combining kinematic and chemical measures that a full understanding of their constituent parts will arise. The fossil record of past events is in the spatial distribution, kinematics and chemical fingerprint of the individual stars; to address the nature of dark matter and the baryonic processes within the Galaxy we need to obtain the kinematic and chemical signatures of large populations (≈ millions) of stars.

4 Survey Strategies To achieve the preceding goals hinges on two complementary approaches: a low resolution (LR) survey to yield the kinematics and overall chemical abundance of stars over a large volume, and an overlapping higher resolution (HR) survey with to uncover the detailed patterns in stellar chemistry.1 The synergy of these complementary surveys combined with the astrometric capability of the GAIA mission would provide the most detailed picture of the processes that formed our Galaxy and of the nature of the component parts.

4.1 LR Survey Requirements An LR (R ∼ 5000) mode should focus upon obtaining accurate kinematics and metallicity measures for a large stellar sample. Here the main Galactic survey science (15 < V < 20) would be complementary to the GAIA (V < 15–17) and RAVE (V < 12) large area surveys. GAIA overlap is important since for both full kinematic and HR abundance analyses, good distances and accurate proper motions are required. This suggests that in general matching to GAIA survey depths for the LR survey, i.e. to V ≈ 20, is sensible. However, we also do not want to lose sight of the advantage of an 8 m-class telescope for deeper surveys in selected nearby galaxies. Sufficient velocity accuracy is required to identify streams and substructure against the Galactic background populations [13]. For radial velocity surveys that can be combined with accurate proper motions from, say, GAIA, accuracies ≈ 2–4 km/s are required for disk and halo substructures respectively. Furthermore, although dwarf satellite galaxies typically have velocity dispersions in the range 6–12 km/s, star clusters range from 1 km/s (open clusters) to 10– 20 km/s (massive globulars). As the velocity dispersion of tidal tails of disrupting systems reflects that of the progenitor, detecting and analysing dissolving clusters 1 Although GAIA will revolutionise this field, it will lack detailed abundances and accurate radial velocities to sufficient depth to fully explore Galactic history.

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Table 1 FLAMES-GIRAFFE R = 20 000 EWs compared to UVES R = 50 000 EWs Line y

λ (Å)

FLAMES EW (mÅ)

UVES FWHM (Å)

EW (mÅ)

FeI

6297.84

122.19

0.36

123.2

OI

6300.32

39.99

0.30

41.4

ScII

6300.65

26.80

0.38

25.4

FeI

6301.53

128.92

0.33

126.9

FeI

6302.49

117.73

0.36

101.0

TiI

6303.81

72.33

0.43

79.5

also requires errors of a few km/s. Complications arise from the superposition of smooth foreground populations [14], and LR abundance measures would be needed to differentiate substructure. An efficient mid-resolution (R = 5000) spectroscopic survey, comparable to FLAMES in LR mode, could be used to obtain velocities to ≈ 2 km/s precision and simultaneous proxy [Fe/H] measurements to 0.1 dex via Mgb and NIR CaT regions. Although this would not allow chemical fingerprinting of individual stars it would alleviate the need for an unfeasibly large HR sample and directly allow the tracing of the dissolution and tidal debris from star clusters, where km/s precision and overall abundance indicators are needed.

4.2 HR Survey For an HR mode there appears little to gain by having R ≥ 40 000, with limited wavelength coverage compared to a much larger simultaneous wavelength coverage but at R = 20 000. Results from VLT-FLAMES, Keck-HIRES and Magellan-MIKE, suggest that generic chemical tagging using this strategy is possible and by focusing on redder wavelengths (4800–6800 Å) studying a million stars at HR is eminently achievable. In this regime there are dozens of FeII lines and more than a hundred FeI lines to derive accurate stellar atmosphere parameters: effective temperature; surface gravity; metallicity and microturbulent velocity; using excitation and ionisation equilibrium and known curves-of-growth. Virtually all the most important elements for chemical tagging also have measurable lines in this region. Furthermore, with extended spectral coverage it is feasible to contrast conventional line indices (e.g. Lick—often all that is available for distant galaxies) with detailed line analysis. R = 20 000 is a compromise between resolving distinct lines and wavelength coverage and by working in the “visible” regime the problem of line blending is lessened but still affects some important lines. An important example is the forbidden oxygen line [OI]6300.3 , shown in Fig. 1; the [OI] line is half-blended with ScII6300.7 , but is still distinct enough to quantify using multi-component profile fitting as illustrated in the figure and accompanying Table 1, with the main limitation

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Fig. 1 The region around the [OI]6300 line from a 5400s VLT-FLAMES exposure at R = 20 000 of a metal-rich, [Fe/H] = −0.41, Bulge giant. An unconstrained multi-component Gaussian model fit of the most prominent lines is shown in red. A colour version of this figure is available at dx.doi.org/10.1007/978-1-4020-9190-2_33

for metal-rich stars being the continuum placement. Such HR data will, for the first time, provide the accurate fingerprinting of a large sample of stars and, and coupled with GAIA and LR surveys will allow a detailed dissection of the evolution of the Milky Way.

References 1. 2. 3. 4. 5. 6. 7. 8. 9. 10. 11. 12. 13. 14.

K. Freeman, J. Bland-Hawthorn, Annu. Rev. Astron. Astrophys. 40, 487 (2002) R. Ibata, G. Gilmore, M. Irwin, Nature 370, 194 (1994) M.G. Abadi, J.F. Navarro, M. Steinmetz, V.R. Eke, Astrophys. J. 597, 21 (2003) J.S. Bullock, K.V. Johnston, Astrophys. J. 635, 931 (2005) A.S. Font, K.V. Johnston, P. Guhathakurta, S.R. Majewski, M.R. Rich, Astron. J. 131, 1436 (2006) M. Steinmetz, ASPC 298, 381 (2003) R. Ibata, S. Chapman, M. Irwin, G. Lewis, N. Martin, Astrophys. J. 634, 287 (2005) R. Guhathakurta et al., Astron. J. 131, 2497 (2006) E. Tolstoy et al., Astrophys. J. 617, 119 (2004) http://cosmology.lbl.gov/BOSS/as2_proposal.pdf (2007) M. Perryman et al., Astron. Astrophys. 369, 339 (2001) V. Belokurov et al., Astrophys. J. 642, 137 (2006) A. Helmi, S. White, P.T. de Zeeuw, H. Zhao, Nature 402, 53 (1999) A.G.A. Brown, H.M. Velazquez, L.A. Aguilar, Mon. Not. R. Astron. Soc. 359, 1287 (2005)

Near-field Cosmology with the VLT Steffen Mieske and Helmut Jerjen

With the arrival of wide-field imagers on medium-size telescopes (e.g. SkyMapper, Pan-STARRS) and the future co-existence of LSST with the E-ELT, it is worthwhile to evaluate the scientific potential of a CCD camera with ≈ 1 degree FOV for the VLT. Here we discuss the role that such an instrument could play in resolving persisting fundamental problems in “near-field cosmology”.

1 Science Case Dwarf galaxies in the local universe are easily studied survivors from the epoch of galaxy formation, and thus preferred targets to establish empirical benchmarks for high redshift cosmological studies, in particular calibrating theories of galaxy formation and interpreting the observed galaxy luminosity function. While CDM predictions for structure formation on large scales agree well with the distribution of baryonic matter (galaxies), it is in the low mass regime where strong discrepancies persist between the expected frequency of low-mass dark matter halos and the number of known dwarf galaxies (e.g. [12, 18, 19, 30]). The most prominent place affected by this so-called substructure crisis is the Local Group, even when accounting for the recently discovered ultra-faint dwarf spheroidals [24, 31]. A solution of this problem may well be found in fundamental physics such as warm or self-interacting dark matter [2, 25]. Alternatively, there is no shortage of astrophysical mechanisms that can diminish the accumulation of baryons in low mass dark matter potential wells: long cooling times for primordial gas in small halos [5], galactic winds driven by supernovae and hot stars [3], or pressure support against collapse of the intergalactic plasma after reionization [4, 26]. The present picture is further confused because there are a multitude of reasons why the faint end of the galaxy luminosity function, the optical manifestation of the dark matter mass function and completely governed by dwarf galaxies [1], could deviate from the simple CDM theory expectation. Progress on this fundamental issue is currently limited by observations, not theory as most recent studies of the galaxy luminosity function in Virgo [20, 23, 29] and Fornax [6, 11, 18] demonstrated that the ambiguity in attributing membership status to cluster/group galaxy candidates S. Mieske () ESO Garching, Karl-Schwarzschild-Str. 2, 85748 Garching, Germany e-mail: [email protected] A. Moorwood (ed.), Science with the VLT in the ELT Era, Astrophysics and Space Science Proceedings, doi: 10.1007/978-1-4020-9190-2_34, © Springer Science + Business Media B.V. 2009

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is the prime source of uncertainty on the quest to find the accurate shape, slope, and possible turning point of the galaxy luminosity function. This ambiguity can only be resolved by deriving genuine distances to complete populations of dwarf galaxies.

2 Galaxy Distances from Surface Brightness Fluctuations Especially in dense environments, the faint end slope of the galaxy luminosity function is completely determined by the large number of dwarf elliptical galaxies, stellar systems primarily composed of old stars and having an almost featureless morphology (Fig. 1). An intriguing possibility to directly derive distances to such galaxies is the Surface Brightness Fluctuation (SBF) method, whose theoretical framework was developed by [27]. The method quantifies the statistical pixel-to-pixel variation of star counts across a galaxy image with the major technical advantage of working on unresolved stellar populations. Since these variations normalised to the underlying mean galaxy light are inversely proportional to distance (see Fig. 2), the SBF amplitude can be used as a distance indicator, once the dependence of the amplitude on stellar content (age, metallicity) is corrected for (e.g. [28]). Jerjen and collaborators [7, 9, 10, 21] and Mieske et al. [15, 18] demonstrated that the method works well for low surface brightness dEs as faint as μB,eff = 26 mag arcsec−2 and MB = −10 mag, out to distances of 20 Mpc using 8 m class telescopes.

Fig. 1 Galaxy luminosity functions for a range of environments, broken down into different morphological types [8]. The faint end of the luminosity function in dense environments is dominated by early-type dwarf elliptical galaxies

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Fig. 2 The principle of SBF distance estimates. Two simulated images of early-type galaxies with implemented surface brightness fluctuations, having identical angular size but different distance [14]. The galaxy on the right is at a 4 times larger distance than the galaxy on the left giving a smoother apppearance

3 Statistical Analysis of Dwarf Galaxy Properties Beyond 30–40 Mpc, the SBF method is not efficient anymore for measuring distances to faint dwarf galaxies [14]. However, near-field cosmology clearly should include the study of low-mass substructures over a range of environments of different scale length: inner-cluster distribution, cluster-to-cluster variations, distribution bias with respect to large-scale filaments and voids. To achieve a proper sampling of these different levels of structure scales, one must extend dwarf galaxy studies out to z = 0.05 to 0.1 (see for example [22]), to cover typical filament length scales of a few hundred Mpc. This corresponds to an area on the sky of a few thousand square degrees. From Fig. 3 it is clear that even with relatively short exposures of a few minutes on an 8 m class telescopes, low surface brightness galaxies with MV ∼ −11 mag can be detected. Considering the decrease of angular size with distance, more realistic detection and classification limits for dwarf galaxies are MV ∼ −11 out to 50 Mpc, and MV ∼ −13 out to 200 Mpc (z = 0.05). Going back to Fig. 1, these magnitude limits are well in the regime where the faint end slope α dominates the shape of the galaxy luminosity function. Of course, detecting low-surface brightness dwarf galaxy candidates alone is insufficient to derive their absolute luminosity and constructing a fiducial galaxy luminosity function. Therefore, a deeper SBF survey is required covering a few control regions like dense clusters and loose groups to calibrate secondary distance modulus estimators, such as angular size vs. central surface brightness, and colour (e.g. [18]).

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Fig. 3 Here we compare the image obtained from a 200 s exposure in V of a dwarf galaxy with MV = −11 mag at a distance of 20 Mpc. The left image corresponds to the VLT mirror size (8.4 m), the right image to the VST mirror size (2.6 m). The galaxy is only detectable on the VLT image

4 Scientific Aims and Survey Setup In summary, we advocate the combination of a smaller scale but deep imaging survey with a shallow imaging survey over a much larger area, in order to study the properties of low-mass galaxies in the nearby universe as a function of environment. The scientific topics that can be addressed by conducting such a near-field cosmology survey are the following: • Contrast the galaxy luminosity function with the expected CDM mass spectrum. At which luminosity/mass do baryons decouple from dark matter? How does this depend on environment? This will give crucial input for studies on dark matter phase space properties, reionization, feedback, photoionization. • What are the clustering properties, spatial and angular correlation function of low mass galaxies compared to CDM predictions. What is the origin of satellite galaxies? • Constrain the anisotropy of dark matter distribution on scales ≤ 100 Mpc (Great Attractor). • Morphological segregation/transformation, harassment, as a function of environment. • Synergy with other scientific areas include the study of globular cluster systems, intra-cluster light, far-field cosmology surveys: ISW, cluster counting, weak lensing, BAO. While the shallow, dE identification survey could in principle also be done based on data from future surveys with 4 m class telescopes, the deep distance survey requires the light collecting area of an 8 m class telescope. This in order to directly derive distances to faint dwarf galaxies with the SBF method. For the deep survey we estimate an area of about 500 square degrees would be sufficient to fully cover the most prominent nearby galaxy clusters (d ≤ 50 Mpc) and a substantial portion of the low density field environment. It would require about one hour of total integration time per pointing shared between V and I (or B and R) band exposures,

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in order to derive reliable SBF distances for d ≤ 50 Mpc and MV −14 mag [14]. Most of the time will be used for the red filter, which is generally better suited for SBF measurement [10, 13, 15–17, 28]. The colour information is used to correct the SBF amplitude for stellar population effects. With a wide-field imager (1 sq degree FOV) at the VLT, one would require about 60 nights of observing time for the deep survey. Assuming a 5000 sq degree coverage for the shallow survey, one would require roughly 100 nights of observing time for 200 seconds exposures in two optical filters V and I (or B and R).

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B. Binggeli, A. Sandage, G.A. Tammann, Annu. Rev. Astron. Astrophys. 26, 509 (1988) P. Bode, J.P. Ostriker, M. Turok, Astrophys. J. 556, 93 (2001) A. Dekel, J. Silk, Astrophys. J. 303, 39 (1986) N.Y. Gnedin, Astrophys. J. 542, 535 (2000) Z. Haiman, A.A. Thoul, A. Loeb, Astrophys. J. 464, 523 (1996) M. Hilker, S. Mieske, L. Infante, Astron. Astrophys. Lett. 397, L9 (2003) H. Jerjen, K.C. Freeman, B. Binggeli, Astron. J. 116, 2873 (1998) H. Jerjen, in Encyclopedia of Astronomy and Astrophysics, ed. by Paul Murdin (Institute of Physics Publishing, Bristol, 2000) H. Jerjen, K.C. Freeman, B. Binggeli, Astron. J. 119, 166 (2000) H. Jerjen et al., Astron. Astrophys. 380, 90 (2001) A. Kambas, J.I. Davies, R.M. Smith, S. Bianchi, J.A. Haynes, Astron. J. 120, 1316 (2000) A. Klypin, A.V. Kravtsov, O. Valenzuela, F. Prada, Astrophys. J. 522, 82 (1999) S. Mei et al., Astrophys. J. 625, 121 (2005) S. Mieske, M. Hilker, L. Infante, Astron. Astrophys. 403, 43 (2003a) S. Mieske, M. Hilker, Astron. Astrophys. 410, 445 (2003b) S. Mieske, M. Hilker, L. Infante, Astron. Astrophys. 438, 103 (2005) S. Mieske, M. Hilker, L. Infante, Astron. Astrophys. 458, 1013 (2006) S. Mieske, M. Hilker, L. Infante, C. Mendes de Oliveira, Astron. Astrophys. 463, 503 (2007) B. Moore, S. Ghigna, F. Governato et al., Astrophys. J. Lett. 524, 19 (1999) S. Phillipps et al., Astrophys. J. 493, L59 (1998) R. Rekola, H. Jerjen, C. Flynn, Astron. Astrophys. 437, 823 (2005) S.D. Rosenbaum, D. Bomans, Astron. Astrophys. 422L, 5 (2004) S. Sabatini et al., Mon. Not. R. Astron. Soc. 357, 819 (2005) J.D. Simon, M. Geha, Astrophys. J. 670, 313 (2007) D.N. Spergel, P.J. Steinhardt, Phys. Rev. Lett. 84, 3760 (2000) A.A. Thoul, D.H. Weinberg, Astrophys. J. 465, 608 (1996) J.L. Tonry, D.P. Schneider, Astron. J. 96, 807 (1988) J.L. Tonry et al., Astrophys. J. 546, 681 (2001) N. Trentham, S. Hodgkins, Mon. Not. R. Astron. Soc. 333, 423 (2002) N. Trentham, R.B. Tully, Mon. Not. R. Astron. Soc. 335, 712 (2002) S.M. Walsh, H. Jerjen, B. Willman, Astrophys. J. Lett. 659, 121 (2007)

Chemical Evolution of Local Group Galaxies Gražina Tautvaišien˙e, Doug Geisler and George Wallerstein

Among the most interesting applications of the VLT and ELT are high-resolution spectral investigations of stars in Local Group galaxies. It is of paramount importance to determine detailed abundances of a variety of chemical elements in these stellar systems to explore their chemical evolution and search for its dependence on global parameters like mass and morphological type. In this contribution we give a short overview of the outstanding observational tasks of the VLT and ELT in the next decades.

1 Spiral Galaxies in the Local Group 1.1 Milky Way Since the classic paper by Eggen, Lynden-Bell and Sandage [9] there have been many attempts to model the chemical evolution of the Milky Way ([7, 26, 27, 34], and references therein). Among the remaining open questions we mention the thickdisk origin and evolution, early Galactic chemical evolution, the inner vs. outer halo and radial abundance gradients in the Galaxy. More than twenty years have passed since the thick disk of the Galaxy was discovered [12], however the origin and evolution of this population is still under discussion. Several hundreds of high-resolution spectral abundance analyses of thickdisk candidates have now been carried out ([2, 3, 11, 22, 30, 31, 37] and references therein), however the database of comprehensive abundance analyses still needs to be enlarged. The detailed investigation of stars with [Fe/H] > −0.3 is needed in order to answer a question recently raised by Reddy et al. [31] concerning the presence/absence of the knee in α/Fe at this metallicity in the thick-disk. The larger sample of thick disk stars is also needed to determine radial and vertical gradients in compositions of thick disk stars. The largest telescopes are needed in order to analyse the most metal-deficient stars in our Galaxy. In 2000–2001, the ESO Large Programme “First Stars” lead G. Tautvaišien˙e () Institute of Theoretical Physics and Astronomy, Vilnius University, Gostauto 12, Vilnius 01108, Lithuania e-mail: [email protected] A. Moorwood (ed.), Science with the VLT in the ELT Era, Astrophysics and Space Science Proceedings, doi: 10.1007/978-1-4020-9190-2_35, © Springer Science + Business Media B.V. 2009

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by R. Cayrel collected high resolution, high S/N ratio spectra of the Galactic stars of the lowest known metallicity. Since Hill et al. [14], there are now more than ten papers published in this series. It is crucial to address such questions as how homogeneous was the early halo? What was the timescale for mixing? How does the detailed chemistry of the inner and outer halo differ? It is also important to probe the different halo components ‘in-situ’ in order to avoid possible biases associated with probing only a nearby sample, as is the case so far. The problem of the time variation of radial chemical gradients is far from being settled, either theoretically or observationally (cf. [6, 20, 36], and references therein). Different models predict opposite behaviours of the metallicity gradient, showing the sensitivity of this observable to the parametrisation of the adopted physical processes. Very useful information could be brought by investigations of abundance gradients of different chemical elements.

1.2 Andromeda M 31 is the largest spiral in the Local Group. Owing to its proximity and long history of observations, M 31 has provided a vast amount of observational data which could be used to constrain disk, bulge and halo formation models of M 31 [8, 17, 24, 42]. The initial expectation was that these various components in Andromeda and the Milky Way have had similar formation and chemical evolution histories. However, when investigated in detail, it is seen that e.g. a Kennicutt star formation law with radial dependence, which can reproduce well the disk gas and abundance profile in the Milky Way disk, can not reproduce the M 31 gas profile in the outer part of the disk, where it predicts a steeper abundance gradient [15]. The M 31 outer disk requires a higher star formation efficiency. Until very recently, it was believed that the M 31 halo has about ten times higher metallicity than that in the Milky Way. However, recent very wide-field studies (e.g. [18]) have found that some of the stars which have traditionally been regarded as the halo in M 31 belong to the very extended bulge. Much further study is required to sort out the extent of M 31’s components. A further key difference between these galaxies is the existence of a very substantial thin disk in M 31’s globular clusters [25], which does not exist in the Galaxy. The age and chemical signature of this disk have very important implications for the formation history of M 31. Much more observational data are needed in order to find out and understand all the similarities and differences between the Milky Way and Andromeda. High-resolution spectral analyses of stars in stellar streams which are found both in the Milky Way and M 31 are needed.

1.3 Triangulum Triangulum (M 33), the third largest galaxy in the Local Group, may or may not be a companion to the Andromeda galaxy but probably has the Pisces dwarf as a satellite.

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M 33 is an ideal target to test galactic chemical evolution models because it shows no signs of recent mergers and no presence of a prominent bulge or bar component (cf. [32]). The recent classification of individual stars in M 33 (e.g. [1, 4]) suggests the existence of different stellar populations and therefore of separate episodes of star formation. Recently, Magrini et al. [21] presented a model which is able to reproduce the available observational constraints on the distribution of gas and stars in M 33 and to predict the time evolution of abundances of several chemical elements. The detailed evolution of other chemical elements in this galaxy is of interest as well.

2 Dwarf Spheroidal Galaxies The Sagittarius dwarf spheroidal galaxy is currently being ripped apart and accreted onto the Milky Way [16]. M 54 is one of the most massive globulars in this galaxy and was probably the nucleus of Sgr [5, 35]. Stars in this galaxy span a wide range of metallicities, −1.6 ≤ [Fe/H] ≤ 0, and show quite interesting abundances. For metalpoor stars with [Fe/H] < −1, [α/Fe] ≈ +0.3 similar to Milky Way halo stars, but for more metal-rich stars the ratio of [α/Fe] as a function of [Fe/H] is lower than that in the Milky Way disk by 0.1–0.2 dex [23]. Tautvaišien˙e et al. [38] have noticed that the pattern of the abundance ratios in the Sgr dSph is quite similar to that in the Magellanic Clouds, which might mean that this galaxy before its capture was as massive as the LMC. Detailed abundances only exist for a small number of stars in this galaxy. Most of the recognised members of the Local Group are known to be dwarf spheroidal systems (dSphs). Detailed abundances now exist for at least a few stars in each of the classical dSphs associated with the Milky Way, except for Leo II. It is now imperative to obtain similar data for the host of new dSphs recently discovered by the Sloan survey. In addition, work on more distant dSphs in the Local Group has to be continued, in particular those associated with M 31. Also, it is important to obtain high spectral resolution data for the most metal-deficient stars, which reflect the chemical composition of a galaxy before the onset of supernovae of Type Ia. This would help us to test the CDM hypothesis that large dSph-like objects might be accreted to form spiral galaxies during the very early timescales of their evolution [13]. Finally, the dSphs that have been studied most intensely to date show evidence of multiple populations and it is of great interest to compare the detailed abundances of these populations. The ELT or at least VLT-ESPRESSO are needed for this work.

3 Irregular Galaxies Many models have been put forward for the chemical evolution of the Magellanic Clouds. In order to explain the distinct O/Fe and α/Fe ratios, which are generally found to be lower than in Galactic stars of the same metallicity, different models

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applied a steepened initial mass function (e.g. [33]) or selective winds (e.g. [29]). Since among the α-particle elements oxygen showed a large deficit relative to iron and a similar deficit was also found in Galactic supergiants, Pagel and Tautvaišien˙e [28], unlike previous authors, have applied yields and time delays identical to those that they previously assumed for the solar neighbourhood. They included inflow and non-selective galactic winds as well as bursting star formation rates, and fitted most of the observational data within their substantial scatter. So far, the amount of high quality data for older LMC stars is very limited and almost non-existant for the SMC. Up to now only a few supergiants in each of only four dwarf irregular galaxies have had stellar abundances derived with high dispersion spectroscopy (NGC 6822 and WLM by Venn et al. [41]; Sextans A by Kaufer et al. [19]; and IC 1613 by Tautvaišien˙e et al. [39]). The dwarf irregular systems are often considered as potential galactic building blocks, e.g. of the Milky Way halo, in modern hierarchical galaxy formation scenarios. However, it looks like the currently investigated dwarf irregular galaxies have lived most of their life in isolation and exhibit quite different chemical abundances than the spirals and similar to those of the dSphs. The number of dIrr galaxies investigated has to be comprehensively increased. In addition, larger telescopes will allow us to study the fainter and older giant stars which will tell us about conditions earlier in the history of these galaxies and not just about the current conditions which the supergiants probe.

4 Final Remarks It is possible that the giant spirals, the Milky Way and M 31, in the Local Group formed in a different way. The role of smaller galaxies as building blocks is not defined yet. Only a few stars in the Magellanic Clouds, Sagittarius and several dSphs are known to have metal-deficient stars with [α/Fe] similar to that in the Milky Way halo stars. Even the smallest galactic systems did not form simply. The VLT is making, and the ELT and/or at least VLT-ESPRESSO will continue to make, critical contributions to the investigation of the chemical evolution of Local Group galaxies.

References 1. 2. 3. 4. 5. 6. 7. 8. 9.

M.K. Barker, A. Sarajedini, D. Geisler et al., Astron. J. 133, 1125 (2007) T. Bensby, S. Feltzing, I. Lundström, Astron. Astrophys. 410, 527 (2003) T. Bensby, A.R. Zenn, M.S. Oey, S. Feltzing, Astrophys. J. Lett. 663, 13 (2007) D.L. Block, K.C. Freeman, T.H. Jarrett et al., IAU Symp. 235, 8 (2006) J.A. Brown, G. Wallerstein, G. Gonzalez, Astron. J. 118, 1245 (1999) G. Cescutti, F. Matteucci, P. Francoise, C. Chiappini, Astron. Astrophys. 462, 943 (2007) C. Chiappini, F. Matteucci, D. Romano, Astron. Astrophys. 554, 1044 (2001) A.I. Dias, M. Tosi, Mon. Not. R. Astron. Soc. 208, 365 (1984) O.J. Eggen, D. Lynden-Bell, A.R. Sandage, Astrophys. J. 136, 748 (1962)

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10. A. Ferguson, M. Irwin, S. Chapman et al., in Island Universes – Structure and Evolution of Disk Galaxies, ed. by R.S. de Jong (Springer, Dordrecht, 2007), p. 239 11. K. Fuhrmann, Astron. Astrophys. 338, 161 (1998) 12. G. Gilmore, N. Reid, Mon. Not. R. Astron. Soc. 202, 1025 (1983) 13. D. Geisler, G. Wallerstein, V.V. Smith, D.I. Casetti-Dinescu, Publ. Astron. Soc. Pac. 119, 939 (2007) 14. V. Hill, B. Plez, R. Cayrel et al., Astron. Astrophys. 387, 560 (2002) 15. J.L. Hou, L. Chen, R.X. Chang, in The Three-dimensional Universe with Gaia, ed. by C. Turon, K.S. O’Flaherty, M.A.C. Perryman. ESA SP, vol. 576 (2005), p. 687 16. R.A. Ibata, M.J. Irwin, G. Gilmore, Nature 370, 194 (1994) 17. S.A. Josey, N. Arimoto, Astron. Astrophys. 255, 105 (1992) 18. J.S. Kalirai, K.M. Gilbert, P. Guhathakurta et al., Astrophys. J. 648, 389 (2006) 19. A. Kaufer, K.A. Venn, E. Tolstoy et al., Astron. J. 295, 171 (2004) 20. W.J. Maciel, R.D.D. Costa, M.M.M. Uchida, Astron. Astrophys. 397, 667 (2003) 21. L. Magrini, E. Corbelli, D. Galli, Astron. Astrophys. 470, 843 (2007) 22. I. Mashonkina, T. Gehren, C. Travaglio, T. Borkova, Astron. Astrophys. 397, 275 (2003) 23. A. McWilliam, T.A. Smecker-Hane, in Cosmic Abundances as Records of Stellar Evolution and Nucleosynthesis, ed. by T.G. Barnes, F.N. Bash. ASP Conf. Ser., vol. 336 (2005), p. 221 24. M. Mollà, F. Ferrini, A.I. Dias, Astrophys. J. 466, 668 (1996) 25. H.L. Morisson, P. Harding, K. Perrett, D. Hurley-Keller, Astrophys. J. 603, 87 (2004) 26. B.E.J. Pagel, G. Tautvaišien˙e, Mon. Not. R. Astron. Soc. 276, 505 (1995) 27. B.E.J. Pagel, G. Tautvaišien˙e, Mon. Not. R. Astron. Soc. 288, 108 (1997) 28. B.E.J. Pagel, G. Tautvaišien˙e, Mon. Not. R. Astron. Soc. 299, 535 (1998) 29. L.S. Pilyugin, Astron. Astrophys. 310, 751 (1996) 30. J.X. Prochaska, S.O. Naumov, B.W. Carney et al., Astron. J. 120, 2513 (2000) 31. B.A. Reddy, D.L. Lambert, C. Alliende Prieto, Mon. Not. R. Astron. Soc. 367, 1329 (2006) 32. M.W. Regan, S.N. Vogel, Astrophys. J. 434, 536 (1994) 33. S.C. Russel, M.A. Dopita, Astrophys. J. 384, 508 (1992) 34. M. Samland, G. Hensler, Ch. Theis, Astrophys. J. 476, 544 (1997) 35. A. Sarajedini, A.C. Layden, Astron. J. 109, 1086 (1995) 36. L. Stanghellini, M.A. Guerrero, K. Chunga et al., Astrophys. J. 651, 898 (2006) 37. G. Tautvaišien˙e, B. Edvardsson, I. Tuominen, I. Ilyin, Astron. Astrophys. 380, 578 (2001) 38. G. Tautvaišien˙e, G. Wallerstein, D. Geisler et al., Astron. J. 127, 373 (2004) 39. G. Tautvaišien˙e, D. Geisler, G. Wallerstein et al., Astron. J. 134, 2318 (2007) 40. K.A. Venn, D.J. Lennon, A. Kaufer et al., Astrophys. J. 547, 765 (2001) 41. K.A. Venn, E. Tolstoy, A. Kaufer et al., Astron. J. 126, 1326 (2003) 42. L.M. Widrow, K.M. Perret, S.H. Suyu, Astrophys. J. 588, 311 (2003)

The VLTI as a Tool to Study Eclipsing Binaries for an Improved Distance Scale K. Shabun, A. Richichi, U. Munari, A. Siviero and B. Paczynski

Even in the ELT era, angular resolution in the range of milliarcsecond and less will remain the prerogative of long-baseline interferometers. One area in which milliarcsecond capability is essential is the study of eclipsing binaries, for which the VLTI has already started to produce valuable results and for which we foresee an increase in the statistics as fainter and fainter magnitudes become possible with the introduction of facilities such as FINITO and PRIMA. Long-baseline interferometry at facilities such as the ESO VLTI is beginning to have the capability to measure directly in the range of milliarcsecond and less the angular separation and the angular diameter of some selected eclipsing binary systems, and we have begun to carry out such observations with the AMBER instrument. In the special case of double-lined eclipsing binaries with well-detached components, from radial velocity and light curves it is possible to obtain a full solution of all orbital and stellar parameters, with the exception of the effective temperature of one star, which is normally estimated from spectral type or derived from atmospheric analysis of the spectrum or reddening-corrected photometric colors. In particular, we aim at deriving directly the effective temperature of at least one of the components in the proposed system, thereby avoiding any assumptions in the global solution through the Wilson-Devinney method. We have obtained an independent check of the results of this latter method for what concerning the distance to the system. This represents the first step toward a global calibration of eclipsing binaries as distance indicators. Our results will also contribute to the effective temperature scale for hot stars. The extension of this approach to a wider sample of eclipsing binaries could provide an independent method to assess the distance to the LMC. This can only be achieved by the VLTI, as the ELT will not have the required angular resolution.

K. Shabun () European Southern Observatory, Karl Schwarzschild Strasse 2, 85748 Garching, Germany e-mail: [email protected] A. Moorwood (ed.), Science with the VLT in the ELT Era, Astrophysics and Space Science Proceedings, doi: 10.1007/978-1-4020-9190-2_36, © Springer Science + Business Media B.V. 2009

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Fig. 1 Averaged H- and K-band squared visibilities of δ Ori, observed with AMBER UT1-UT3-UT4 (dots). The model fit (curves): separation 1.9 mas, position angle 75 degrees, diameters of 0.9 and 0.6 mas respectively. Each dot/curve color corresponds to a different baseline

Our goal is to measure a restricted number of EB systems with a in the range 1–5 mas, and θ1 , θ2 0.4–1.2 mas. This will be accomplished by observing each system at least four times, at precise separated phase of the orbital period P corresponding to maximum separation. Time has been granted to observe δ Ori, a = 1.4, θ1 = 0.98, θ2 = 0.55, P = 5.73d (distance 280 pc), in the ESO Period 78, unfortunately only two observational runs have been successfully accomplished in December 2006 and March 2007, the rest two runs have been shifted to the Period 80. Also four additional observations of δ Ori has been granted for the Period 80, as well as the time for the observations of η Ori, a = 1.7, θ1 = 0.84, θ2 = 0.7, P = 7.98d and R CMa, a = 1.1, θ1 = 0.65, θ2 = 0.5, P = 1.13d. The Fig. 1 shows the averaged squared visibilities of δ Ori observed by AMBER at 31.12.06, in the low-resolution LR-JHK mode. The symbols are the visibilities averaged over the wavelengths, H-band 1.72 and K-band 2.26 microns, in reality we have 16 channels in the K and 11 channels in H band. The solid curves in both figures represent the (preliminary) best fit to the data of the separation and position angle of the binary, with respect to the length and position angle of the baselines. Support high quality radial velocity and photometric light-curve for the proposed targets are currently being collected at Asiago Observatory. 49 high resolution Asiago Echelle spectra of δ Ori have already been collected in 10 different nights during the time interval of ESO period 78, the spectra cover from 3800 to 7300 Å at a resolving power 30 000, and we achieved a S/N per pixel on the extracted spectrum in excess of 150.

References 1. B. Paczynski, Space Tel. Sc. Inst., 273 (1997)

Interferometric Investigations of Eclipsing Binaries 2. U. Munari, A. Siviero et al., Astron. Astrophys. 418, L31 (2004) 3. A. Siviero, U. Munari et al., Astron. Astrophys. 417, 1083 (2004) 4. A. Richichi et al., Astron. Astrophys. 431, 773 (2005)

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Part IV

VLT and VLTI Synergy with ELTs

Status of the European ELT Roberto Gilmozzi and Jason Spyromilio

Abstract The EELT project is in Phase B (detailed design), a 3-year, 57.2 M Euro activity that is scheduled to conclude in a Proposal for Construction by end 2009 or early 2010. The starting point for the current phase was a community process leading to the convergence of earlier concepts into a single European project. The telescope baseline is for a fully steerable 42 m, segmented primary, 5 mirror design, fully adaptive system with Nasmyth and Coudé foci. This paper reports on the status of the EELT Phase B.

1 Introduction Following extensive community consultation in early 2006 a basic reference design for the European Extremely Large telescope was developed. The starting point for the design was the toolbox generated by the ELT Science and Engineering working group and its panels, and was presented in detail to the ESO committees and to the European community at the Marseille conference in early December 2006. The outcome was unanimous support for the design and 42-m size. In December 2006 the ESO Council resolved that the EELT should proceed into Phase B with the aim to have a proposal for construction ready to be submitted to the ESO Council in late 2009 or early 2010. The 2007–2009 resources allocated to Phase B are 57.2 M Euro, including approximately 110 FTEs. During phase B, contracts are being placed with industry for the advancement of major subsystems to preliminary design status. Specifically, contracts are currently in place for the development of the main structure, the dome, the adaptive mirrors, the tip-tilt unit and the primary mirror support. Several prototype mirror segments are being procured and polished to the specifications of the project: this will enable our industrial partners to establish robust production processes. Integrated modeling, development of concepts for the control system, the mirror cells and the adaptor rotators are ongoing, as is the design of the secondary unit. For critical subsystems where more than one technology exists or where more than one approach is possible, multiple contracts have been placed. At the present time, more than 50% of the Phase B budget has been committed. The instrumentation phase A and point design studies have been launched. The site evaluation process is ongoing within the FP6 ELT-DS and is supported by the project wherever necessary. R. Gilmozzi () European Southern Observatory, Garching bei München, Germany e-mail: [email protected] A. Moorwood (ed.), Science with the VLT in the ELT Era, Astrophysics and Space Science Proceedings, doi: 10.1007/978-1-4020-9190-2_37, © Springer Science + Business Media B.V. 2009

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During Phase B several reviews shall take place, both internal and external, culminating with a cost review in late 2009 or early 2010 followed by a design review before the completion of the Proposal for Construction. There have been some changes in the schedule, mostly to consolidate interesting new ideas and results from the contractors, but none has impacted the critical path.

2 The Optical Design Three ELT projects are being funded for detailed design: the European ELT, TMT and GMT. Interestingly there exists a diversity of optical designs driven from the variations of the requirements and the baseline engineering solutions adopted by the projects. The GMT is Gregorian, the TMT Ritchey-Cretien (after a period of courtship with the Gregorian design) and the European ELT (after trading off spherical primaries Gregorian and RC designs) has settled for a 5 mirror design. The European ELT has chosen a three powered mirror optical design on the basis that as the fields and the aberrations become ever larger fixing them inside the instruments also becomes increasingly challenging. The 3-mirror anastigmat is in common use in instrumentats but it is not often proposed, nor used, in telescopes. The need for an additional two reflections to get the beam to an instrument friendly location such as a Nasmyth platform results in a loss of efficiency which is irrecoverable. In the EELT the additional reflections are used to embed adaptive optics and field stabilization capabilities and therefore the tradeoff is more complex than simply counting glass-air surfaces. In an adaptive telescope and in the particular case of an extremely large one, the mitigation of risk at the same time as ensuring performance has the EELT to the 5-mirror design. In this design the EELT has an elliptical 42-m segmented primary mirror, a convex monolithic 6-m secondary, an aspheric concave tertiary located in the middle of the primary and two flat relay mirrors as shown in Fig. 1.

Fig. 1 Left: The optical design of the EELT. Right: One potential solution for the EELT mount. A colour version of this figure is available at dx.doi.org/10.1007/978-1-4020-9190-2_37

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The optical design offers excellent image quality across the field of view of the telescope, minimal aberrations in the out of focus laser guide return beams and a natural location (above the quaternary mirror) for a full field atmospheric dispersion compensator. The tradeoff to adopt this design is complex. One the negative side are the necessity for broad wavelength reflective coatings on the mirrors of the telescope as they affect all instruments, the limited field of view of the telescope due to necessity of passing through the hole in the quaternary and the challenge of making the secondary mirror, the large central obstruction (10% of the beam), the mechanical complexity of supporting the quaternary and M5 and many others. On the positive side we can see the excellent image quality, the ability to trade off aberration corrections between active mirrors, the relatively easy upgradability of the system and others. These tradeoffs were undertaken during 2006 in collaboration with industrial partners. For a 42-m adaptive telescope the conclusion was that the dimensions of an adaptive mirror in either a Gregorian or a Ritchey-Cretien design would be prohibitive in cost and risk and the scientific and technological cost of not including adaptive optics into the telescope would be strategically the wrong conclusion. Making multiple smaller telescopes did not address the scientific aims of the project.

3 The Telescope Mount For an enormous telescope such as the EELT a “rocking chair”-like telescope mount with the optical axis below the primary mirror is a natural evolution from large radio telescope structures and with the bulk of the telescope mass located below the primary mirror. This solution is used by the GMT project. For the ELT, and for the TMT as well, the mechanical design teams have been tasked to keep the altitude axis above the primary mirror. This partly ensures that thermal sources do not interfere with the beam and critically allows for multiple foci to be accessed around the telescope “tube”. Two industrial firms were contracted by ESO in 2007 to develop solutions to the mount and evaluate its cost and manufacturability. Both firms developed solutions for a balanced structure with good mirror cell stiffness and plausible eigen frequencies. The control of the telescope mount under wind loading is good enough that a field stabilization mirror can handle the residuals. During 2008 the EELT project will be consolidating the various design inputs. The upper Serrurier truss differs between telescopes, with a tripod like structure being preferred by GMT, where the relatively large gaps between their 8.4-m diameter round segments provides ample room for such a structure. In the solution shown above the EELT has a relatively thin upper structure and in other solutions a stiffer tower like structure is considered to limit the deflections of the secondary mirror unit. The EELT design process includes at all stages a full dynamic modeling of the structure and an analysis of the sensitivity to earthquakes. As can be inferred from the statements above, the open loop tracking performance of the EELT is the dominated by wind-shake. Locked rotor eigen frequencies for such an enormous structure are around 2.5 Hz and therefore the EELT is

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not a priori stiff enough to reject the effects of the wind. Moreover, with a large secondary it is hard to minimize the telescope profile towards the wind. Attention has been given to the issue of wind shake during the design of the system and in the derivation of the specifications of the field stabilization unit. Detailed simulations of the telescope performance when exposed to the wind have shown the need for a field-stabilization capability. Even with sophisticated control strategies it is not possible to reduce the residual image motion at the focal plane to an irrelevant fraction of the diffraction limit using the tracking capabilities of the telescope mount. This has already been the case for the seeing limited operation at the VLT so it is no surprise that for the EELT we face the same challenge. For the ELTs, field stabilization is addressed differently amongst the projects. The GMT project is proposing to segment their secondary mirror thereby matching the primary. Each subunit is comparable to those already under construction for 8-m class telescopes. TMT are using a dome design to dramatically reduce the wind loading on the upper parts of the telescope while not restricting air flow. Postfocal corrections will be needed at all telescopes in any case when observing below the seeing limit. At the EELT the fifth mirror in the optical train is to be used for field stabilization.

4 The Mirrors 4.1 The Primary For the EELT the primary mirror will comprise of 984 segments. In total 1148 will be built, a number that provides for 1 spare per family of six segments. The EELT project is paying detailed attention to the problem of serial production of such a large number of segment assemblies. Both EELT and TMT have concluded that the segment size must remain large to limit the control complexity (fewer actuators, fewer edge sensors, fewer wiffle trees, etc.) without increasing it so much that the total mass of the primary increases the requirements on the mirror cell dramatically. The EELT and TMT projects would not have a single identical segment, however, the two projects have determined that a 1.45-m point-to-point hexagonal segment will meet their requirements. This shared baseline increases the potential supply base for the assemblies and reduces the amount of development required in the areas of support and control of the segment positions and shapes. The EELT project has contracted two firms to deliver 7 prototype full sized segments. The support of the primary mirror segments is based on the solutions of developed for the Keck telescopes and evolved for the GTC, namely a wiffle tree support. The principal challenge is to load the segment only axially without introducing lateral forces as the telescope inclines down to 70 degrees. As each segment is an off axis aspheres a small displacement, or rotation, relative to their perfect location appears as an optical aberration. In addition to these effects, residual aberrations from polishing will be corrected using motors to rotate the wiffle tree arms about their

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pivot points, thereby redistributing the loads within the wiffle tree. Phasing the primary mirror requires that the segments be moved in piston and tip-tilt relative to each other. Three actuators are used to move the segment and its support against the mirror “cell”. To provide the phasing with the misalignment information inductive edge sensors have been developed with the FP6 ELT design study program. In addition, other novel technologies for sensing the differential location of the segments are being examined in collaboration with industrial firms.

4.2 The Secondary The convex secondary is planned to be a low thermal expansion glass or glass ceramic mirror. This 6-m diameter mirror may be as thin as 100-mm assuming it required the same flexibility as the VLT primaries. The design of the mirror cell is the subject of an industrial study currently in the tendering phase. For the purposes of establishing a baseline ESO has developed a concept that requires the mirror cell to support the mirror and to correct its shape while large actuators move the cell with respect to the spider to maintain the telescope collimation. In Fig. 2 the ESO concept can be seen.

4.3 The Adaptive Relay Unit The tertiary, quaternary and fifth mirrors are mounted in the adaptive relay unit or (ARU). The tertiary mirror is a 4.2-m mirror. Depending on the requirements for an adjustable back-focal distance, the tertiary may be active. The 2.5-m quaternary mirror is flat. It provides the wide field adaptive functions. Two preliminary design studies have been funded with industrial firms and they include prototyping parts of the M4

Fig. 2 Left: Conceptual design of the secondary mirror in its cell. Center: The adaptive relay unit. Right: One conceptual design of the M4 unit. A colour version of this figure is available at dx.doi.org/10.1007/978-1-4020-9190-2_37

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Fig. 3 Left: Conceptual design of the field stabilization unit. Right: One potential solution for the EELT dome. A colour version of this figure is available at dx.doi.org/10.1007/978-1-4020-9190-2_37

system. The studies are focused on different technologies for the actuation of the mirror and also provide for different densities of actuators. As has already been discussed the wind-shake of the telescope mount will be compensated by the field stabilization M5 mirror that is also flat. A design study with an industrial firm is underway for the electromechanical unit. including a scale one prototype. The provision of the mirror is the subject of multiple studies both in house and in industry. The M5 of the EELT is comparable in size to the SOFIA primary that was light-weighted to approximately 120 Kg/m2 . The requirements on the M5 mean that will need to be 2 to 3 times lighter than that (per unit area) while at the same time maintaining a high stiffness.

5 The Dome The telescope will need to be protected from the sun and other weather during the day. For this a dome will be built. During observations providing ventilation while minimizing the wind speed at the level of the secondary is another key function of the dome. In addition, providing a maintenance friendly environment for the daytime activities while not becoming a maintenance issue in itself are extra requirements on the dome design. A number of dome designs are under investigation. A box-like enclosure has been studied while two independent firms are detailing the design of a spherical dome. A calotte design, as adopted by TMT, provides excellent wind protect. However, such a design was not investigated by the EELT project as it is not considered as an optimal solution to the challenge of a dome. The EELT dome provides a crane with a 20-ton capacity and includes a lifting platform that allows for access to the Nasmyth platforms. The same lifting platform provides the access for servicing the secondary mirror unit. The dome airconditioning is designed so that the ambient air temperature at the beginning of each night is reflected on the inside of the dome. A slight over pressure on the inside of the dome helps to keep the dome clean.

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6 Control The EELT control strategy is being considered as an evolution of the ESO experience in deploying active optical systems (from the NTT to the VLT). For the EELT the analysis of the system is following an integrated approach, including partnering with European institutes working in this area. The scope of each system is also clearly identified to units with clear interfaces to other subsystems, thereby minimizing global complexity. The Paranal field experience following many millions of active optics corrections already performed and more that half a dozen adaptive optics systems in common use is fully encapsulated in the EELT project. In addition, the experience of combining the 8-m telescopes for the purposes of interferometry is providing the EELT project with valuable ideas regarding the importance of vibrations and the difficulties of working so close to the diffraction limit.

7 Conclusions The EELT project is in its detailed design phase and is required to submit a construction proposal in late 2009 or early 2010. For most subsystems industrial contractors are establishing technical solutions and costs while ESO concentrates on the overall system design and engineering. The project is an active collaboration with the scientific and industrial community and the eventual success of this enterprise will be a credit to the long development in Europe of knowledge centers and to a collaborative spirit.

The Science Case for the European ELT Isobel Hook

1 Introduction The European ELT (E-ELT) project has developed dramatically over the last year and has moved into the detailed design phase. The science case has been developed over many years with much of the community involvement and meetings being sponsored by the EC-funded OPTICON network [1]. The activity involves over 100 astronomers from around Europe. In December 2005 ESO created 5 working groups whose primary role was to re-assess various aspects of the E-ELT following the review of the OWL project. The Science Working Group (SWG) produced a report in April 2006 [2] describing over 30 science cases and their feasibility with telescopes in the range 30 m to 60 m diameter. From among these cases, the SWG selected a sub-set of “Prominent Science Cases” as follows: • Exo-planets: – Direct detection; – Radial velocity detection. • Initial Mass Function in stellar clusters. • Stellar disks. • Resolved Stellar Populations: – Colour magnitude diagrams; – Abundances; – Detailed abundances and kinematics. • Black Holes. • The physics of galaxies. • Metallicity of the low-density IGM. • The highest redshift galaxies. • Dynamical measurement of the Universal expansion. In mid 2006 the ESO and OPTICON ELT SWGs merged into a single European ELT SWG, which provides ongoing scientific input to the E-ELT project. One of the tasks of the SWG is to develop a Design Reference Mission. I. Hook () University of Oxford Astrophysics, Denys Wilkinson Building, Keble Road, Oxford OX1 3RH, UK e-mail: [email protected] A. Moorwood (ed.), Science with the VLT in the ELT Era, Astrophysics and Space Science Proceedings, doi: 10.1007/978-1-4020-9190-2_38, © Springer Science + Business Media B.V. 2009

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1.1 The Design Reference Mission The goals of the Design Reference Mission (DRM) are to (i) produce a set of example science proposals and simulations and (ii) assess science output and assist with tradeoff decisions. For each science case, observing proposal is written which sets out the science goals and observing modes required. The PI also provides the information required to simulate the observations (such as target brightness, spatial distribution, etc.). Simulations are then carried out by staff at ESO and/or in some cases by members of instrument study teams. The results of the simulations are then passed back to the PI for an assessment of whether the science goals are met. In some cases this will provide feedback that helps guide specifications of the instrumentation, AO systems, etc. The process can then be iterated as the design is optimised. Initially three “Demonstrator cases” were selected to be the first for simulations, and these are well underway. Following this a wider set of proposals will be simulated (based on prominent cases). The three “Demo Cases” are: • Direct Detection of Exo-planets: extreme contrast observations; • Stellar populations: I-K Colour-Magnitude Diagrams; • Galaxy mass assembly: multi-IFU resolved Spectroscopy.

2 Science Cases A very brief description of the prominent science cases is given below. Initial results from the DRM are included where appropriate. This list of cases is certainly not complete, and it is likely that some of the most exciting science that the ELT produces will be things we cannot predict today. By considering the cases below we aim to build a flexible system that covers a wide parameter space and that can adapt as new scientific discoveries are made.

Exo-planets Many planets have been detected by indirect methods such as radial velocity and transit searches. Recently, high-contrast, adaptive-optics (AO) imaging has begun to produce direct detections of faint companions to nearby young stars. However known exo-planets show a wide range of properties and are generally very different from planets in our own Solar System (see review [3]). The ELT will make dramatic advances in both direct and indirect detection of exo-planets. Direct Detection of Exo-planets: DRM Demo Case #1 Direct detection provides physical information such as mass, orbit, temperature and composition of the exo-planet. Very high spatial resolution and contrast are required,

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of order 109 between the parent star and planet. A lot will be learned from the future VLT/SPHERE and Gemini Planet Imager instruments [4, 5]. However the ELT will have the fundamental advantage of larger diameter and hence higher resolution. A planet-finding instrument, EPICS, is being studied for the European ELT [6]. Simulations carried out by the EPICS team for the case of photon-limited observations and assuming a perfect coronograph show that (for example) in a 10 hour observation with a 42 m telescope, Jupiter-like planets can in principle be detected at 15 pc from us and Earth-like planets detected at 5 pc. However in practice the number of suitable stars within 5 pc of us is very limited and hence the probability of detecting an Earth-like planet is low. In addition various systematic effects such as gaps or misalignment between mirror segments, and the effect of speckles will impact the detectability of planets. These effects are currently being studied during the initial part of the EPICS Phase-A study.

Indirect Detection: Radial Velocity The radial velocity technique has provided the majority of exo-planet candidates to date. The increased collecting area of an ELT will allow this technique to be extended towards lower mass planets including Earth-like planets around solar type stars. It will allow the discovery of new planets and the follow up of planets detected from other techniques (such as transit searches) in order to determine their physical parameters. Taking into account the strategies to overcome the main stellar intrinsic limitation of RV estimates, i.e. acoustic oscillation modes and activity-related jitter, to reach the required accuracy (< 10 cm/s), each star would require ∼ 100 h of observations over ∼ 4 years. These observations require a high resolution, ultrastable optical spectrograph.

2.1 Stellar Disks The existence of circumstellar disks around young stellar objects is now firmly established. The diffraction limit of a 42 m telescope at 10 µm corresponds to about 1 AU at for a stellar system at 20 pc from us, (approximately 5 times higher resolution than JWST), opening up the possibility of detecting structure in stellar disks that are predicted by some models of planet formation. Imaging in the optical and near-IR probes the scattered light from the disk while the thermal IR is sensitive to the dust emission. Spectroscopy will be used to probe dynamics and chemical processing in the disk. A range of spectral resolution will be used to study dust, gas and ices (e.g. silicates, molecular gas and even organic materials).

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2.2 The Initial Mass Function The shape of IMF is dependent on the processes of star formation, and provides a key diagnostic between star formation models. Massive stars in clusters provide feedback of energy and chemically enriched material into the ISM, yet the shape of the IMF at the bright end is hard to measure, particularly in forming stellar clusters which are both crowded and often highly extincted. The spatial resolution and sensitivity of the E-ELT at near and mid-IR wavelengths will allow determination of number densities and brightnesses of massive stars embedded in clusters despite huge amounts of dust extinction. At the low-mass end the sensitivity of the ELT to point sources will allow measurement of the low-mass end of the IMF in a range of environments, to test for variations in the IMF and the existence of a low-mass cutoff. E-ELT Spectroscopy may be used to characterise objects from wide-field imaging surveys from other telescopes.

2.3 Resolved Stellar Populations Measuring the age and chemical composition of individual stars within a galaxy reveals the various stellar populations that it contains, and hence provides a history of the galaxy. Detailed studies of this type are currently possible within the Milky Way and Local Group galaxies, but to reach a representative sample, including giant ellipticals, requires the resolution and sensitivity of an ELT. There are two main probes: (i) imaging to provide colour-magnitude diagrams from which ages and crude metallicity measures can be made, and (ii) spectroscopy to provide chemical abundances and kinematics. (i) Colour-magnitude diagrams (photometry)—DRM demo case #2 Simulations of this case have been carried out as part of the DRM. The specific goal chosen was to obtain colour-magnitude diagrams in K, I-K for galaxies at various distances. The stellar distributions (magnitudes and number densities on the sky) were provided by E. Tolstoy, and simulated images were produced by J. Liske (ESO) assuming a 42 m telescope and using point spread functions predicted for Laser-tomography AO. The images were run through photometry software and the colour-magnitude diagram reconstructed. Initial results suggest that in a reasonable observing time the position of the tip of the Red Giant Branch can be recovered for galaxies at a distance modulus of ∼ 31 mag, i.e. just at the distance of Virgo. The Horizontal Branch can be recovered to ∼ 28.5 mag, e.g. in M83. The main sequence turnoff (required for unambiguous age determination) can be located to a distance modulus of ∼ 26 mag, i.e. to the distance of Cen A. These results are critically dependent on the delivered Strehl (in particular at shorter wavelengths) and also on the method used for photometry experience with current and upcoming AO systems will test these assumptions and methods.

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(ii) Spectroscopic chemical abundances and kinematics Moderate resolution spectroscopy of the Calcium Triplet (860 nm) can provide kinematics and a metallicity estimate for individual stars. It is expected that the ELT will be able to make such measurements on the brightest stars at the tip of the Red Giant Branch in galaxies at the distance of Virgo. Higher resolution (R = 20 000–70 000) optical spectra would provide abundances and ages of early (old), metal-poor stars in the Milky Way halo and in nearby systems, e.g. using the Uranium/Thorium age indicators. This places strong requirements that the ELT operates at the bluest possible range (short-ward of 4000 Å).

2.4 Black Holes Only a few black holes have accurate mass measurements, and key questions remain such as “How common are they?” and “Why do their masses relate to the mass of their host galaxy bulges?”. The spatial resolution of the ELT will allow the sphere of influence of black holes to be probed at large distances from us. It is expected that the ELT can study 109 M black holes to 100 Mpc (and perhaps to high redshift, depending on surface brightness limitations), which will allow a statistical sample to be constructed. A 106 M black hole, similar to that in our Milky Way, could be studied at the distance of the Virgo cluster. These observations require high spatial resolution, ideally with an IFU, and sensitivity to faint surface brightnesses.

2.5 The Physics and Mass Assembly of Galaxies out to z ∼ 6 (DRM Demo Case #3) One of the main goals in astronomy is to understand formation of galaxies, a current hot topic being the relevance of feedback processes (from SNe and AGN). Recent observations have demonstrated the power spatially resolved spectroscopy (via AOfed integral field spectrographs) to measure physical parameters of distant galaxies [7]. To date this has been possible only for relatively bright galaxies. With an AOassisted multi-IFU instrument on an ELT it should be possible to extend these studies to large statistical samples of galaxies with a wide range of luminosities, and to significantly higher redshifts. This case has been simulated as part of the DRM. The specific case studies is a kinematic study of ∼ 1000 massive 2 < z < 6 galaxies. The diagnostic chosen was whether or not a merger could be distinguished from a rotating disk. Simulations were carried out by M. Puech (ESO). The input is based on IFU observations of nearby galaxies and models of high-z galaxies. Provisional conclusions indicate that reliable kinematic studies of super-L* galaxies can be made to z ∼ 5 with a 42 m ELT, and to z ∼ 2 for galaxies with 0.1 M∗ .

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2.6 Early Type Galaxies at z > 1 Early type galaxies are the most massive galaxies, contain most of stellar mass and are tracers of mass assembly and the growth of structure. However they are faint, red, and have no emission lines in the optical part of the spectrum (in the observed frame), hence it is difficult with current technology to obtain high quality spectra. In order to study star formation histories, masses and evolution of scaling relations requires a large sample of such galaxies. Such as sample may be found using near-IR surveys (such as will be provided by VISTA) but to obtain redshifts and physical diagnostics will require multi-object spectroscopy with an ELT over the wavelength range (0.6–2.5 µm). Note that this case does not require high-order AO but does require an instrument with a wide patrol field (of order 25–100 sq arcmin).

2.7 Very High-Redshift Galaxies The highest redshift galaxies known at present, with z ∼ 7 and possibly higher, have been found with major investments of observing time on the world’s largest telescopes (e.g. [8, 9]). Despite these heroic efforts the signal-to-noise that can be obtained is low. In order to find and study galaxies at higher redshift will require a combination of JWST (well suited to finding targets) and ELT (well suited to spectroscopy and high-resolution imaging). High resolution imaging will provide an assessment of whether the systems are relaxed or merging, and whether they are gas-rich. Spectroscopy will provide physical diagnostics and may even allow detection of He lines (and indicator of Population III stars). R > 1000 spectroscopy on an ELT will be competitive with JWST in terms of exposure time but requires large FOV and high multiplex (∼ 100).

2.8 The Low Density IGM The metal enrichment of the Intergalactic Medium (IGM) at high redshift depends on the density of sources ejecting metals during the re-ionisation epoch, and the presence of galactic super-winds at later times (z < 5). Hydrodynamic simulations predict CIV column densities of 109 –1010 cm−2 for regions with over-density of about unity [10]. To detect such CIV column densities requires a factor ∼ 100 improvement in the detection limit of CIV doublets over that achievable with existing telescopes-initial estimates show that this could be achieved in ∼ 20 hrs on a 42 m ELT for a V = 16 QSO. This type of observation requires high resolution spectroscopy (R ∼ 100 000) in the optical.

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2.9 Cosmic Expansion One of the biggest questions in physics today is “What is the Dark Energy that is driving the acceleration of the Universe?”. An interesting possibility that appears to be feasible with an ELT is direct measurement of cosmic expansion in real time, by repeat observations of QSO absorption lines over a period of 10–20 years. The tiny change in redshift (of order a few cm/s/yr) caused by the acceleration of the Universe over this period is in principle detectable, and provides a independent dynamical measurement of the expansion that can be compared to geometrical measurements from supernovae or the BAO method [11]. CODEX is a proposed ultra high-stability, high-resolution (R ∼ 150 000) spectrograph for such measurement (see talk by J. Liske at this meeting). This instrument could be used for radialvelocity studies of exo-planets.

2.10 Some Requirements on Instrumentation in the ELT Era Many of the cases discussed above would benefit greatly from observations from other facilities, either 8 m class optical/IR telescopes such as VLT or smaller survey telescopes such as VISTA, or facilities operating at other wavelengths (e.g. ALMA). In terms of ground-based observations the main theme that emerges from these cases is a need for wide field surveys. Wide field infra-red surveys: Imaging surveys are required to find the targets for several science cases, such as very high-redshift galaxies, z > 1 early-type galaxies, low-mass objects (for IMF measurement) and new rare objects in general. Spectroscopic surveys are also required for initial classification of the sources (and/or redshift measurement in some cases). For cases requiring relatively narrow-field, deep surveys JWST may provide the sources, but for wider-field surveys groundbased 4 m and 8 m class telescopes would be better suited. In order to obtain spectroscopic classification of large numbers of sources, a high-multiplex multi-object spectrograph would be needed, with a wide field of view to match the imaging surveys as much as possible. Wide field optical surveys: Two of the prominent cases above require large samples of bright QSOs (CODEX, IGM studies). These could be found with a wide-field optical survey on a relatively small telescope (2–4 m). Identification and redshift measurement would require a corresponding spectroscopic survey.

References 1. I. Hook (ed.), The Science Case for the European Extremely Large Telescope: The Next Step in Mankind’s Quest for the Universe (OPTICON & ESO, 2005) 2. Report of the ELT Science Working Group (April 2006), available at http://www.eso.org/sci/ facilities/eelt/publications.html

232 3. 4. 5. 6. 7. 8. 9. 10. 11.

I. Hook S. Udry, N.C. Santos, Annu. Rev. Astron. Astrophys. 45, 397 (2007) K. Dohlen et al., SPIE 6269E, 24 (2006) B. Macintosh et al., SPIE 6272E, 18 (2006) C. Vérinaud et al., SPIE 6272E, 19 (2006) R. Genzel et al., Nature 442, 786 (2006) M. Iye et al., Nature 443, 186 (2006) D. Stark et al., Astrophys. J. 663, 10 (2007) R. Cen, K. Nagamine, J.P. Ostriker, Astrophys. J. 635, 86 (2005) L. Pasquini et al., IAU Symp. 232, 193 (2006)

GRB Afterglows in the ELT Era David Alexander Kann and Sylvio Klose

Afterglow phenomenology on a statistical basis is a substantial tool to get insight into the physical processes at work (cf. [4, 8]). Since several years we have been undertaken such an approach by gathering and analyzing the largest possible optical/NIR data set in a systematic way [1–3, 6, 7]. In Zeh et al. [6] we analyzed all optical afterglows searching for supernova light appearing at late times and discovered that the data indicate that all long bursts are related to supernova explosions. In Zeh et al. [7] and Kann et al. [1] we investigated the light curve shape and the spectral energy distribution of all optical afterglows known in the pre-Swift era in a systematic way. Finally, in the most recent and most comprehensive publications [2, 3] we used our data base to discuss the properties of the afterglows of short bursts in comparison to the long burst sample, as well as comparing the long GRB afterglows of the pre-Swift with those of the Swift era. Figure 1 shows the light curves of a sample of optical afterglows of long GRBs from pre-Swift (gray) and the Swift era (black) selected for good light curve coverage and known redshift, complete until the end of August 2007. The data are corrected for Galactic extinction and, where possible, for host galaxy contribution, but otherwise as observed. In red, we plot the light curves of short GRBs, in all cases, afterglows were detected (square data points), but additional upper limits are given too (downward pointing triangles). Clearly, the afterglows of short GRBs are much fainter than those of long GRBs, and high S/N or high-resolution spectroscopy of these dim cosmic beacons will require the light-gathering power of an ELT. The inset figure shows early light curves of GRBs that had rapid spectroscopic observations, within 0.1 days. The fastest observations were obtained with VLT UVES in Rapid Response Mode [5]. This mode, which allows the follow-up of transients within minutes, will keep the VLT competitive far into the ELT era as a powerful tool to study the early evolution of the afterglows of the most powerful explosions in the Universe.

References 1. D.A. Kann et al., Astrophys. J. 641, 993 (2006) D.A. Kann () Thüringer Landessternwarte Tautenburg, Tautenburg, Germany e-mail: [email protected] A. Moorwood (ed.), Science with the VLT in the ELT Era, Astrophysics and Space Science Proceedings, doi: 10.1007/978-1-4020-9190-2_39, © Springer Science + Business Media B.V. 2009

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Fig. 1 GRB optical afterglow light curves with well-sampled data up to the end of August 2007. Pre-Swift long GRB afterglows are grey, Swift era ones are black. Short GRB afterglows are red, with detections marked by squares, and additional upper limits by triangles. The inset shows early light curves of long GRBs with rapid spectroscopic follow-up. A colour version of this figure is available at dx.doi.org/10.1007/978-1-4020-9190-2_39

2. 3. 4. 5. 6. 7. 8.

D.A. Kann et al., Astrophys. J. (2007), submitted D.A. Kann et al. (2008), in preparation A. Panaitescu et al., Mon. Not. R. Astron. Soc. 366, 1357 (2006) P. Vreeswijk et al., Astron. Astrophys. 468, 83 (2007) A. Zeh, S. Klose, D.H. Hartmann, Astrophys. J. 609, 952 (2004) A. Zeh, S. Klose, D.A. Kann, Astrophys. J. 637, 889 (2006) B. Zhang, Nature 444, 1010 (2006)

On the Way to an E-ELT Instrumentation Plan Sandro D’Odorico, Mark Casali and Vincenzo Mainieri

1 From the VLT to the European ELT The European Extremely Large Telescope (E-ELT) project is very much based on the heritage of the ESO Very Large Telescope, the array of four 8 m optical telescopes located on the Paranal peak in the Atacama desert. The VLT started operation in 1998 and with its complement of 11 instruments permanently on line has proved to be the most powerful facility for ground-based astronomy at optical and infrared wavelengths (see the contribution by A. Moorwood in these Proceedings). In the case of the VLT, an Instrumentation Plan Proposal was distributed to the community as early as 1989 and this approach, with corrections and upgrades introduced on the way, has been very effective in developing a coherent set of instruments in collaboration between ESO, Universities and Institutes in the ESO member countries. In the case of the E-ELT, ESO is now coordinating a series of instrument studies with the goal to define a first generation of instruments to be included in the proposal for construction due at the beginning of 2010. In the last 25 years the collecting power of the ground-based telescopes has increased dramatically with as many as 15 new 8–10 m telescopes in or about to enter into operation. The acquisition of photons from faint cosmic sources using larger apertures has a major impact on many field of astrophysics but two other parallel developments had an equally positive influence: first, array detectors for both UVVisual and NIR wavelengths have become regularly available in size up to 2K × 2K pixels (for CCD up to 4K × 2K) and with QE above 60% over the whole range of sensitivity: second, through the introduction of active optics pioneered by ESO at the NTT in the late 80’s, the optical image quality of the telescope has matched the best seeing achievable at the high quality sites (that is down to 0.3 ). In the last decade an additional capability has been demonstrated at the telescope: the acquisition of images corrected for atmospheric turbulence by use of adaptive mirrors (e.g. NACO, SINFONI at the VLT). In the case of the E-ELT, it is planned to build a telescope with an associated Adaptive Optics system that would be able to perform at the diffraction limit (θ = 1.22λ/D) over a moderate field with high Strehl S. D’Odorico () European Organisation for Astronomical Research in the Southern Hemisphere, K. Schwarzshild Str. 2, 85748 Garching bei Muenchen, Germany e-mail: [email protected] A. Moorwood (ed.), Science with the VLT in the ELT Era, Astrophysics and Space Science Proceedings, doi: 10.1007/978-1-4020-9190-2_40, © Springer Science + Business Media B.V. 2009

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(concentration of light in the diffraction peak) in the NIR. A class of scientific objectives which relies on the angular resolution (e.g. planetary discs, regions of star formation, active galactic nuclei) will benefit from this unique advantage. Apart of the angular resolution, the diffraction-limit capability will result in larger gain in the sky-limited observations of any stellar-like sources (stars, GRB, SN, QSO): the time to reach a given S/N will be proportional to D −4 . Somewhat lower but still very significant will be the gain in the observations of sources like high z galaxies which are not stellar but can be made up by sub-seeing knots.

2 Instrument Locations in the E-ELT Baseline Concept The current opto-mechanical concept of the 42 m E-ELT (B. Delabre, 2008, Astron. Astrophys., in publication) foresees a five-mirror design feeding two f/16 Nasmyth foci and a coudé type focus. Mirrors 3, 4 and 5 are located in the ARU tower. M4 is an adaptive mirror and M5 fast tip/tilt mirror. Figure 1 shows the telescope structure. Each Nasmyth platform can host up to four instrument focal stations. These can be fed directly by mirror 5 of the telescope or be positioned on the sideway of the optical axis and be fed by an additional mirror located in the adaptor module. For observing modes which require a full AO correction a dedicated optomechanical module hosting up to two additional AO mirrors has to be inserted in the platform between the telescope and the instrument. Figure 1 shows the possible location of a conceptual Multi Conjugate Adaptive Optics (MCAO) system. Instruments which require a large volume and/or a very stable and thermally controlled

Fig. 1 View of the 3-D model of the baseline E-ELT. A colour version of this figure is available at dx.doi.org/10.1007/978-1-4020-9190-2_40

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environment will be installed at a coudé focus in the telescope basement and be optically fed by a train of highly efficient mirrors from the Nasmyth. The telescope-instrument interfaces are being studied as part of the E-ELT Phase B Study which will lead to the E-ELT construction proposal. Whatever the details of these interfaces and of the instrumentation observatory infrastructure, it can be conceived that up to 7 instruments of first generation could be eventually installed at the E-ELT and be operated in a flexible schedule mode like as the VLT.

3 Instruments under Study for the E-ELT Based on the work of the ELT Science and ELT Instrumentation Working Groups1 completed in 2Q 2006 on the parallel Instrument Small Studies carried out under the FP6 ELT Design Study, on the technical work for the 42 m telescope concept at ESO in 2006 and on a first assessment of the resources for the instrument development within the E-ELT Program, a roadmap to the definition of an ELT Instrumentation Plan was presented by ESO to the community on November 29–30, 2006 at a dedicated meeting in Marseille (France) and incorporated in the proposal for a Phase B E-ELT endorsed by the ESO Council in December 2006. The cornerstones of the future E-ELT instrumentation Plan were identified as follows. The scientific objectives of highest priority will dictate the choice of the first generation instruments. The instruments will be distributed in a permanent way among the different focal stations to permit rapid switching and the execution of a broad range of scientific programs at any time. As in the case of the VLT most of the instrument design and construction work will be carried out by the community. ESO will play a coordinating role and it will develop critical subsystems and hardware and software standards. When necessary it will also support specific projects with overall system engineering and management and it will take the responsibility of the final integration and testing in Europe before shipment to the telescope. The paradigm of exchanging guaranteed observing time for staff effort has been consequently adopted for most VLT instruments and it resulted in a strong commitment by the community to the VLT success. It is the starting assumption for the E-ELT instrumentation effort. An instrumentation budget has been developed under these assumption and foresees the construction of 5–6 instruments of first generation with the associated AO modules. According to these guidelines ESO has launched 6 Phase A studies on instruments and two on complementary post-focal Adaptive Optics systems in 2007. The instruments to be studied have been taken from those identified as of highest priority and are briefly described below. The final reports will have to include the calculated scientific capabilities, the cost, the required FTE effort and a construction schedule. The studies will eventually 1 http://www.eso.org/sci/facilities/eelt/publications.html.

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be the basis for the Technical Specs and Statement of Work of the first instrument construction contracts, should their priority and feasibility be confirmed by 2009. In addition to the 6 Phase A studies, ESO identified the need to carry out two additional Concept Studies during this phase. These will be centered on observing modes which were not yet identified as of highest priority but are needed to explore the potential future capabilities of the telescope. Based on the results of the Phase A and Concept studies, on the refinement of the scientific priorities at that date and on the budget envelope, the consolidated list of the first generation E-ELT instruments will be released in 2010 as part of the E-ELT Construction Proposal. Wide-Field, Multi IFU, NIR Spectrograph + AO. The scientific case of this instrument is focused on the 2D spectroscopy of galaxies at redshifts from z = 1 to the epoch of re-ionization over a large field of view (up to 5 arcmin in diameter). AO simulations carried out during the OWL and FP6 ELT-DS instrument studies have shown that an efficient instrument of this type can be built if a field-distributed MOAO system can be realized. A Phase A study of this instrument (acronym EAGLE) and the associated AO system is being executed by a Consortium of LAM, LESIA, GEPI, ONERA (France) and ATC Edinburgh, CfAI Durham (UK). The P.I. is J.G. Cuby from LAM. High Resolution, High Stability Visual Spectrograph. The two key scientific cases of this instrument are the kinematical measurement of the expansion of the Universe and the detection of planets from the radial velocity wobbling of the parent star down to Earth masses. OWL and FP6 ELT DS studies of an instrument of this type (acronym CODEX) have identified a concept which now needs to be fully validated by a Phase A study. The study of this high resolution, high stability spectrograph is carried out by a Consortium of Institutes (Geneve Obs., IAC Tenerife, INAF Trieste and Brera, IoA Cambridge) coordinated by ESO with L. Pasquini as P.I. High Angular Resolution Camera. The main scientific case of this instrument rests on the study of stellar like objects (population in selected stellar systems, nuclei of galaxies) at the highest possible angular resolution (close to Diffraction Limit at NIR wavelengths) over a moderate field (≥ 30 arcsec). In addition to the adaptive M4 in the telescope this requires two other deformable mirrors in a MCAO module. The Phase A study of this instrument (acronym MICADO) is carried out by a Consortium of Institutes (MPE, MPIfA, Munich Observatory, INAF Padova, NOVA) with R. Genzel as P.I. The Consortium was selected following an open Call for Proposal in the community. E-ELT Planetary Imaging Camera and Spectrograph. To directly detect and take spectra of exoplanets requires angular separation of the star from the orbiting planet, efficient ways to suppress the star light with a coronograph and the selection of suitable wavelength bands. The E-ELT offers the unique opportunity in this field due to its uniquely large diameter and optical quality. Conceptual studies of an instrument to make these observations have been carried within the OWL studies and the FP6 ELT-DS instrumentation effort. A key subsystem in this instrument concept is the associated extreme Adaptive Optical module (XAO) that must deliver

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diffraction limited images of the stellar sources at a very high Strehl (> 95%). ESO has coordinated a large team of institutes in the previous study phases and plans to continue to do so because of the strict relation between telescope, AO and instrument performance. During the instrument (acronym EPICS) study in this phase it is planned to validate the instrument concept by a number of additional studies on key issues such as coronography, focal plane WSSs, differential imaging. The studies will also explore the requirements and limitations coming from some of the telescope technologies (e.g. the M1 segmentation). The study of this instrument and its associated Extreme Adaptive Optics system is carried out by a Consortium of Institutes (LAOG, LESIA, LAM, ONERA, MPIA Heidelberg, INAF Padova and ETH Zurich) coordinated by ESO. M. Kasper is the Project Scientist. Single Field Wide-Band Spectrograph. This instrument has not been studied for OWL or within the FP6 ELT-DS effort. It has been introduced in 2007 to meet a pressing request by the Science Working Group to study a basic, single field spectrograph which could operate from day one, first in seeing limited mode, then in GLAO mode and eventually in LTAO or MCAO. The scientific motivation is the need to access at the earliest possible time and in advance of international competition targets which are accessible to an ELT only trough its unique collecting power and angular resolution. A major synergy is expected with ALMA and the JWST. Again based on the primary scientific drivers identified by the E-ELT SWG, the main operation range of the instrument will be the near infrared but the study will also explore the possibility of a parallel visual-red arm to provide in one shot the full coverage of the spectrum. The Phase A study of this instrument (acronym HARMONI) is carried out by a Consortium of Institutes (Univ. Oxford, ATC Edinburgh, CRA Lyon, DAMI Madrid, IAC) with N. Thatte as P.I. The Consortium was selected following an open Call for Proposal in the community. Mid InfraRed Instrument. The scientific case for this instrument encompasses the study of protoplanetary disks, the physics of the regions of star formation and the dusty nuclei of galaxies. The observations will be complementary to those by JWST in the 3–20 mm bands. The space telescope will go deeper because of the absence of the atmospheric and thermal contributions of the telescope but the ground-based instrument will have a higher angular resolution and collect more photons in the bands that are transmitted by the atmosphere. An instrument concept for these bands had been developed within the OWL activities and the FP6 ELT-DS instrument studies and needs to be optimized to the E-ELT interfaces and constraints. An open Call for Proposal for a Phase A study of a MID Infrared instrument has been released in November 2007 and it is expected to select a Consortium in April 2008. Post-Focal Adaptive Optics Systems. In conjunction with an array of up to 6 Laser Guide stars and the associated WF sensors, the telescope AO system (M4 + M5) will contribute both to keep an excellent image quality of the telescope and to provide atmospheric correction for the ground layer turbulence. Apart from the high resolution, ultra-stable spectrograph, all the current high priority E-ELT instruments require high Strehl or Encircled Energy values. This calls for higher order AO systems (specifically MCAO, MOAO and XAO). During the Phase B these post-focal instrument-oriented AO systems will be studied in parallel to the corresponding instrument. For MCAO, the Phase A study is carried out by a Consortium

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Table 1 E-ELT and TMT E-ELT

Properties

TMT

Wide-field, multi IFU NIR spectrograph + AO

Properties NIR multi-slit spectrometer-imager

EAGLE

Wavelength range: 0.8–2.5 µm

(1st gen

Patrol field ≥ 5

IRMS

FoV: 2.27 diameter imaging

(clone of MOSFIRE @ Keck)

candidate)

Multiplexing > 20

(early

Multiplex factor: 46

Spectral R = 5000 (R > 15 000)

light)

Spectral resolution: R = 3270–4500

IQ: > 30% EE in 100 mas

Spectral coverage: all of Y, J, H or K

High-resolution visual spectrograph

High-resolution visual spectrograph

CODEX

Wavelength range: 0.40–0.69 µm

HROS

Wavelength range goal: 0.31–1.0 µm

(1st gen

(goal 0.35–0.72 mm)

(1st

R = 50 000 (1 slit)

candidate)

R > 120 000; Stability: < 5 cm/s

decade)

Mid-IR imager and spectrograph + AO

Mid-IR imager and spectrograph + AO

Wavelength range: 7–20 µm

MIRES

Wavelength range: 8–18 µm

(goal 3.5–27 µm)

(1st

FoV: 10

candidate)

FoV ≥ 30

decade)

5000 < R < 100 000

diameter; R: tbd

continued on next page

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MIDIR (1st gen

E-ELT

Properties

TMT

Single field wide-band spectrograph

Properties IFU NIR spectrometer and imager

HARMONI

Wavelength range: 0.8–2.4 µm

IRIS

Wavelength range: 0.8–2.5 µm

(1st gen

(goal 0.5–2.4)

(early

FoV < 2 for IFU, 10 × 10 for Imaging

candidate)

FoV: tbd; R ∼ 4000 (R ∼ 20 000)

light)

R = 4000 (J, H, K)

Operation with GLAO, LTAO,

AO module NFIRAOS

MCAO (tbd) High angular resolution camera MICADO (1st gen candidate)

Wavelength range: 0.8–2.4 µm FoV:

≥ 30

Operation with GLAO, LTAO,

Wide-field optical MOS—seeing limited WFOS

Wavelength range: 0.34–1.0 µm

(early

FoV: 40.5 arcmin2

light)

R = 500–5000

On the Way to an E-ELT Instrumentation Plan

Table 1 (continued)

MCAO Planet imager and spectrograph

Planet imager and spectrograph

EPICS

Coupled to XAO, coronograph

PFI

Coupled to XAO, coronograph

(1st gen

Wavelength range: 0.6–1.8 µm

(1st

Wavelength range: 1–2.5 µm

R > 50 in Y to H bands

decade)

R = 50 full FoV, R = 500 partial FoV

candidate)

Polarimetry

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of INAF (Bologna-Arcetri-Padova) and ONERA. The P.I. is E. Diolaiti. An open Call for Proposal for a study of the Laser Tomography Adaptive Optics (LTAO) module will be released within January 2008 with the goal to grant a study contract by April 2008. Instrument Concept Studies. In addition to the six Phase A instrument studies identified above, two additional concept instrument studies will be launched after an open Call for Proposals in January 2008 as part of the preparation of the 1st generation instrument selection in 2009. The studies will be centered either on instruments not included among the six of highest priority or on some of the instrument concepts identified in the document ESO/STC 430 but not yet explored in previous studies. They will complement the Phase A studies by probing unexplored potential capabilities of the telescope.

4 E-ELT Versus TMT Instrumentation The TMT project has presented at the end of 2007 a construction proposal supported by an extensive document on the science objectives.2 Unlike the E-ELT, the TMT Ritchey-Cretien, three mirror optical design does not foresee any adaptive mirror in the telescope. Concerning the instrumentation, the TMT project has identified beside an AO module to be located on the Nasmyth platform (NFIRAOS) three “early light” instruments plus five “1st decade” instruments. In Table 1, the preliminary specifications of the six E-ELT instruments now under study are compared with the corresponding TMT suite. On the E-ELT side two new instrument concepts have still to be identified (see Sect. 2) and on the TMT side two of the “1st decade” instruments, a Near Infrared Echelle Spectrometer (NIRES) and a large field Infrared Multi Object Spectrometer (IRMOS) have not been included in the table. The bulk of the instrument suite of the two projects are very similar. The most notable difference is the inclusion in the TMT suite of the Wide-Field Optical Multi Object Spectrograph, designed to operate at seeing limit down to the UV. The E-ELT has the potential advantage to support all observations with the ground-layer correction provided by the adaptive mirror in the telescope optical train.

2 http://www.tmt.org/foundation-docs/TMT-DSC-2007-R1.pdf.

From ESPRESSO to CODEX J. Liske, L. Pasquini, P. Bonifacio, F. Bouchy, R.F. Carswell, S. Cristiani, M. Dessauges, S. D’Odorico, V. D’Odorico, A. Grazian, R. Garcia-Lopez, M. Haehnelt, G. Israelian, C. Lovis, E. Martin, M. Mayor, P. Molaro, M.T. Murphy, F. Pepe, D. Queloz, R. Rebolo, S. Udry, E. Vanzella, M. Viel, T. Wiklind, M. Zapatero and S. Zucker

Abstract CODEX and ESPRESSO are concepts for ultra-stable, high-resolution spectrographs at the E-ELT and VLT, respectively. Both instruments are well motivated by distinct sets of science drivers. However, ESPRESSO will also be a stepping stone towards CODEX both in a scientific as well as in a technical sense. Here we discuss this role of ESPRESSO with respect to one of the most exciting CODEX science cases, i.e. the dynamical determination of the cosmic expansion history.

1 Introduction CODEX (= COsmic Dynamics EXperiment) is a concept for an extremely stable, high-resolution optical spectrograph for the European Extremely Large Telescope (E-ELT). The science case for CODEX encompasses a large range of topics, including the search for exo-planets down to Earth-like masses, primordial nucleosynthesis and the possible variation of fundamental constants. However, its prime science driver is the exploration of the universal expansion history by detecting and measuring the cosmological redshift drift using QSO absorption lines. This is also one of the 9 ‘prominent’ science cases chosen by the E-ELT Science Working Group to be among the highlights of the entire E-ELT science case. A description of the CODEX project as a whole was given by [3]. The recognition that an ultra-stable, high-efficiency, high-resolution optical spectrograph would not only be an extremely valuable instrument for the E-ELT but also for the VLT led to the development of the ESPRESSO concept (= Echelle Spectrograph for PREcision Super Stable Observations, see L. Pasquini’s contribution to these proceedings). Again, there are a large number of applications for such an instrument, as several other contributions to these proceedings have highlighted. Hence, there is also a very strong science case for ESPRESSO, including e.g. detailed studies of the intergalactic medium and of stellar abundances. J. Liske () ESO, Karl-Schwarzschild-Str. 2, 85748 Garching, Germany e-mail: [email protected] A. Moorwood (ed.), Science with the VLT in the ELT Era, Astrophysics and Space Science Proceedings, doi: 10.1007/978-1-4020-9190-2_41, © Springer Science + Business Media B.V. 2009

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However, apart from these scientific drivers, ESPRESSO will also fulfil another role: CODEX will represent a major development in high-resolution optical spectrographs compared to existing instruments such as UVES and HARPS. In order to achieve its science goals CODEX will have to deliver an exceptional radial velocity accuracy and stability, i.e. 2 cm s−1 over a timescale of ∼ 20 yr. Although the basic design concepts are already in place, several of the sub-systems needed to achieve the CODEX requirements do not currently exist. However, they will be implemented in ESPRESSO for the first time. Hence, in many respects ESPRESSO will be a CODEX precursor instrument that will allow us to test and gain experience with the novel aspects of these instruments. Here we will discuss ESPRESSO in the context of its CODEX precursor role, both in a technical sense as well as with respect to the main CODEX science driver, which we briefly describe next.

2 Cosmic Dynamics The discovery from type Ia SNe that the universal Hubble expansion appears to have begun accelerating at a relatively recent epoch, and its profound implications for fundamental physics have sparked an intense interest in the observational exploration of the Universe’s expansion history. Several methods for measuring the Hubble parameter H (z) already exist but all of them are either geometric in nature or use the dynamics of localised density perturbations. The simplest, cleanest and most direct method of determining the expansion history, however, is to observe the dynamics of the global Robertson-Walker metric itself. One way to achieve this is by measuring the so-called redshift drift, i.e. the tiny, systematic drift as a function of time in the redshifts of cosmologically distant sources (Fig. 1). This effect is directly caused by the de- or acceleration of the universal expansion and can hence be used to determine its history. A measurement of this effect would be able to provide evidence of the acceleration that is entirely independent of the SNIa results or any other cosmological observations, and that does not require any cosmological or astrophysical assumptions at all. It would also provide H (z) measurements over a redshift range inaccessible by other methods. Recently, [1] found that a 42 m telescope would indeed have the photon collecting power to detect the redshift drift by monitoring the redshifts of QSO absorption lines over a timescale of ∼ 20 years (Fig. 1), providing a strong motivation for a CODEX-like instrument at the E-ELT.

3 ESPRESSO as a CODEX Pathfinder 3.1 Technical Aspects The UVES and HARPS experiences have allowed us to identify a number of properties that a spectrograph must feature in order to deliver exceptional radial velocity

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Fig. 1 The solid lines show the redshift drift as a function of redshift in velocity units for two different combinations of ΩM and Ω as indicated, and a Hubble constant of H0 = 70 km s−1 Mpc−1 . The grey shaded areas result from varying H0 by ±8 km s−1 Mpc−1 . The three sets of ‘data’ points show Monte Carlo simulations of a redshift drift experiment with CODEX/E-ELT using a total of 4000 h observing time and a total experiment duration of 20 yr. Each of the sets of points corresponds to a different implementation of the drift experiment pursuing different observational goals. See [1] for more details. A colour version of this figure is available at dx.doi.org/10.1007/978-1-4020-9190-2_41

accuracy and long-term stability, including: simultaneous wavelength calibration, a fully passive instrument with zero human access, located inside a vacuum tank which is itself located inside a nested environment in an underground facility that allows progressively more precise temperature and pressure control, and the fluxweighted timing of observations with sub-second precision. Temperature control of the CCD will be particularly important, while high system throughput will also be of the essence. Some of these concepts are already established. However, two of the most critical aspects are also those requiring the most R&D: light scrambling and wavelength calibration. At a resolution of 150 000 a typical pointing accuracy of ∼ 0.05 arcsec corresponds to an error of 100 m s−1 necessitating a scrambling gain of ∼ 5000 in order to reach 2 cm s−1 accuracy. Hence, in addition to any fibre we will require a dedicated scrambling device to ensure that a photon’s position on the CCD only depends on its wavelength but not on its position on the entrance aperture. Current wavelength calibration sources such as ThAr lamps and I2 cells are suboptimal in several respects, their non-uniformity and lack of long-term stability be-

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ing among the concerns. However, a new concept for wavelength calibration has recently emerged. A ‘laser frequency comb’ system provides a series of uniformly spaced, very narrow lines whose absolute positions are known a priori with a relative precision of ∼ 10−12 (see A. Manescau’s contribution; [2]). Neither of the two systems above currently exist, but they would be developed for ESPRESSO. Being able to test and validate them ‘on the sky’ would provide valuable experience and input for further improvements.

3.2 Scientific Aspects The scientific goals of CODEX are sufficiently removed from current observational reality to require validation and demonstration of feasibility of all aspects of data handling and analysis. This includes data acquisition strategies (e.g. minimum and maximum viable exposure times), the tracking of CCD distortions, and the accuracy of flat-fielding, sky subtraction and scattered light corrections. The extraction of the cosmological signal from the data will also require testing. Issues include how to deal with QSO variability and the accuracy of the conversion to the cosmological reference frame. In addition ESPRESSO will allow us to determine the currently unknown intrinsic widths of the narrowest metal absorption lines in order to reassess their usefulness for the drift experiment, and to look for any sources of astrophysical noise so far overlooked. Data on QSOs collected with ESPRESSO for other scientific purposes would allow us to address all of the above issues. We estimate that ∼ 200 hours of observations of the brightest known QSOs would provide an end-to-end system verification, from data acquisition to signal extraction, at the level of ∼ 30–40 cm s−1 .

4 CODEX + ESPRESSO = z˙ ? Since ESPRESSO will have characteristics similar to those of CODEX the question arises whether ESPRESSO can be used to make a start on the redshift drift experiment. The idea is that since ESPRESSO would be available several years before CODEX data appropriately collected with ESPRESSO could serve as a ‘zeroth’ epoch measurement, effectively extending the time baseline of the experiment for a few years, thereby improving the final result without delaying it. Figure 2 shows the comparison of the cosmological constraints in the Ω –ΩM plane expected from a drift experiment using CODEX only (coloured ellipses) and CODEX + ESPRESSO (solid contour), where we have assumed that ESPRESSO would take up operation 8 years before CODEX. Evidently the extension of the time baseline by 8 years is not enough to offset the lack of the VLT’s photon collecting power compared to the E-ELT: the improvement of the constraints is only quite modest, with the lower limit on Ω increasing by ∼ 20 per cent.

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Fig. 2 Expected constraints in the Ω –ΩM plane from the redshift drift experiment with CODEX/E-ELT described in detail in [1]. The coloured ellipses show the joint 68 and 90 per cent confidence regions that result from a total integration time of 4000 h and a total experiment duration of t = 20 yr. The hashed region indicates the 95 per cent lower limit on Ω . The solid contour shows the improvement of the 68 per cent confidence region that results from the additional investment of 4000 h of observing time using ESPRESSO on the VLT in its ‘SuperHarps’ mode (i.e. at R ≈ 150 000 and using one UT), assuming that these observations take place ∼ 8 yr before the start of the CODEX observations. The 95 per cent lower limit on Ω of the combined experiment is shown has the horizontal line above the shaded region. Flat cosmologies and the boundary between current de- and acceleration are marked by solid black lines. The dark shaded region in the upper left corner designates the regime of ‘bouncing universe’ cosmologies which have no big bang in the past. A colour version of this figure is available at dx.doi.org/10.1007/978-1-4020-9190-2_41

References 1. J. Liske et al., Mon. Not. R. Astron. Soc., submitted (2008) 2. M. Murphy et al., Mon. Not. R. Astron. Soc. 380, 839 (2007) 3. L. Pasquini et al., Messenger 122, 10 (2005)

First Results of AQuEye, a Precursor ‘Quantum’ Instrument for the E-ELT C. Barbieri, G. Naletto, E. Verroi, C. Facchinetti, T. Occhipinti, A. Di Paola, E. Giro, P. Zoccarato, G. Anzolin, M. D’Onofrio, F. Tamburini, G. Bonanno, S. Billotta, C. Pernechele, P. Bolli, V. Da Deppo and S. Fornasier

Abstract In September 2005, we completed a study (QuantEYE, the ESO Quantum Eye, [D. Dravins et al. in Proceedings from Meeting ‘Instrumentation for Extremely Large Telescopes’, held at Ringberg Castle, July 2005 ed. by T. Herbst, 2005]) in the frame of the projects for the 100 m Overwhelmingly Large (OWL) telescope. The main goal was to demonstrate the possibility of approaching with the existing technology the picosecond time resolution needed to bring quantum optics into the astronomical domain, measuring the statistics of the photon arrival time and demonstrating the feasibility of a modern version of the Hanbury Brown Twiss Intensity Interferometry (HBTII). To gain experience with such a novel instrument, we have built and put in operation a prototype (AQuEye) for the Asiago 1.8 m telescope. Here we present its main characteristics.

1 Introduction Photons are very complex entities, carrying more information than extracted in astronomical applications with conventional techniques of imaging, spectroscopy and polarimetry. According to Glauber, Arecchi, Mandel, etc. seminal papers from 1960 onwards, arbitrary states of light can be specified by first, second, and higher order correlation functions with respect to position r and time t . Starting from these general principles, we have taken up the attempt to measure second-order correlation functions with respect to time (statistics of the photon arrival times, entropy characteristics of the photon stream, photon correlation spectroscopy) and to space (HBTII). In addition to time resolution, a crucial important parameter is the diameter of the telescope. First, because the above mentioned correlations are fully developed on time scales of the order of the inverse optical bandwidth. For instance, with the very narrow band pass of 1 Å in the visible, through a definite polarization state, typical time scales are ∼ 10−11 seconds (10 ps), and only very large collecting areas can insure a sufficient photon rate in such a short time interval. Second, the amplitude of the second order correlation function increases with the square of the telescope area (not diameter), so that a 40 m telescope will be 256 times more C. Barbieri () Department of Astronomy, University of Padova, Padova, Italy e-mail: [email protected] A. Moorwood (ed.), Science with the VLT in the ELT Era, Astrophysics and Space Science Proceedings, doi: 10.1007/978-1-4020-9190-2_42, © Springer Science + Business Media B.V. 2009

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sensitive to such correlations than a 10 m telescope. In conclusion, only the future Extremely Large Telescopes (ELTs) can bring quantum optical effects within the astronomical domain. Therefore, a ‘quantum’ instrument for an ELT, able to push the time resolution towards the limits imposed by Heisenberg’s uncertainty principle and to cope with GHz count rate, might have the same scientific impact of opening a new window. We might call this new Astronomy with the name of Quantum Astronomy, or also Photonic Astronomy, to reflect the increasing role of the photon in the optical sciences.

2 QuantEYE Design On the instrumental side, one needs a very fast single-photon detector with high quantum efficiency, linear up to the GHz count rate, running continuously (no gating) and keeping the time tagging capability at the level of 10 ps, for the several hours needed to reach faint astrophysical sources. A CCD-type matrix of such detectors would allow imaging, greatly beneficial to identify the ‘quantum’ star in a complex field such as a nebula, and to eliminate seeing and transparency fluctuations. This ideal detector does not yet exist, although technology is rapidly advancing. The devices available at the time of the QuantEYE study came as single units of small dimensions. Their properties drove the conceptual opto-mechanical design (see [1, 2]). The 100 m pupil of OWL had to be divided in a large number of small sub-pupils. The baseline solution was a fixed area, non-imaging photometer made by a focal reducer (an inverted Cassegrain telescope) at the focus of OWL, plus a 10 × 10 lenslet array feeding 10 × 10 single-photon detectors via fibers. The study also demonstrated that, although the amount of data produced by QuantEYE is really large, modern technology can overcome this problem without unduly large difficulty.

3 AQuEye, the Little Brother of QuantEYE To gain experience for such highly unconventional instrument, we have built a prototype of QuantEye, named AQuEye (the Asiago Quantum Eye, see [3, 4]) for the 182 cm Copernicus Telescope at Asiago—Cima Ekar. Needless to say, such a photometer, even if mounted on a small telescope, can produce data with an exceptionally high dynamic range, limited only by photon statistics. To speed up work, and remain into fairly strict resources, we took advantage of the existing Asiago imaging spectrograph (AFOSC), which provides not only many ancillary functions (shutter, field acquisition and rotation, guiding, controls, etc.), but also the intermediate pupil of the telescope. Following the QuantEYE design, the telescope pupil is divided in 4 sub-pupils. AQuEye thus behaves essentially as a fixed aperture, 4 simultaneous channel, photometer. The optical performances are very good at all wavelengths from 420 to 750 nm, with more than 85% Encircled Energy inside the 50 µm pixel of the detector.

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Fig. 1 AQuEye mounted at the 182 cm telescope. A colour version of this figure is available at dx.doi.org/10.1007/978-1-4020-9190-2_42

As detectors, we acquired Geiger-mode Single Photon Avalanche Photodiodes (SPADs), 50 µm diameter active area. These devices have a quantum efficiency around 50%, can tolerate full day light, are thermoelectrically regulated, their dark count is around 50 Hz, the integrated timing circuit gives a time stamping accuracy better than 50 ps. On the adverse side, in addition to their ‘single pixel, small dimensions’ configuration, is the dead time of ∼ 80 ns (12 MHz max count rate from each SPAD), which is only partly overcome by pupil splitting. The overall acquisition and control system consists of a VME crate mounted on the telescope and connected to a personal computer located in the control room via a fiber optics link. Each detector provides a string of pulses read by a Time-ToDigital (TDC) unit with an internal clock of approximately 40 GHz. Taking into account the several jitter sources in the overall system, we can claim a precision of each time stamp better than 50 ps. All time tags are permanently recorded in a 1 Terabyte external memory. The raw data (namely photon arrival times) can then be binned in arbitrary time intervals. Time-binning is done a posteriori, with integration steps varied at will, so that no information is lost. An external Rubidium oscillator, disciplined by a GPS receiver (in future, a GALILEO receiver will be implemented), provides an accurate time reference to the system. At the telescope, AQuEye is mounted in place of the standard CCD camera of AFOSC (see Fig. 1).

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Fig. 2 The autocorrelation function of time stamps of photons from the pulsar in Crab nebula, as observed in December 2007 in a 1 minute long string of data. A colour version of this figure is available at dx.doi.org/10.1007/978-1-4020-9190-2_42

4 Astronomical Performances Using the measured SPAD QE, AFOSC efficiency, internal optics and filter transmissions, we derive the following performances: 1. On the bright side, the limiting factor is the 8 MHz maximum count rate permitted by the electronics (TDC firmware + VME + fiber optics link), corresponding to a 5th mag star. 2. On the faint side, the Field of View diameter is 3 arcsec, namely an area of 4.71 arcsec2 . The Ekar sky brightness in average conditions is V = 19 mag/ arcsec2 , and the FoV sky background corresponds to a star of V = 17.3. The Dark current is 50 e/s, equivalent to a star of V = 20, so that the sky sets the limiting factor. The linear regime then goes approximately from V = 5.0 to V = 18.0 per channel, so that the 33 ms pulsar in the Crab Nebula is well within the AQuEye capabilities. AQuEye has been used at the telescope in several occasions since June 2007. In addition to the usual engineering tests, we have observed several interesting rapidly variable stars and the pulsar in the Crab nebula. Figure 2 shows the autocorrelation function of the time stamps detected from it, in only one minute of data collection in a night of mediocre seeing. These first AQuEye data have shown that the expectation in terms of sensitivity and timing accuracy are well confirmed, that the software can correctly handle the outputs of the four detectors even at high data rate, and that novel astrophysical

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results can be obtained. The next crucial step, already under way, is time synchronization among two distant telescopes in view of Hanbury-Brown Twiss Intensity Interferometry. To this end we have started a pilot experiment with the Ljubljana Observatory. In summary, although the 182 cm telescope is too small to detect truly quantum effects, with AQuEye we are acquiring convincing evidence that quantum astronomy, in its different aspects (photon stream statistics, photon correlation spectroscopy, HBTII) can be reached with existing technology (to be sure, pushed to its limits) with the E-ELT. Technology of single photon, fast, low noise detectors for visible and near-IR is rapidly advancing, thanks also to the strong push of classical and quantum communication needs, so that by the time ELTs will operate, astronomers cannot miss the capability to utilize the light from celestial objects in such a novel way. Acknowledgements Work partly financed by ESO, the University of Padova, the Italian Ministry for University and Research, and the Galileo Supervising Authority. Thanks are due to A. Cadez and A. Bianchini for useful discussions, and to I. Agnoletto, I. Capraro, L. Lessio and A. Sponselli for help in the laboratory and at the telescope.

References 1. C. Barbieri, V. Da Deppo, M. D’Onofrio, D. Dravins, S. Fornasier, R.A.E. Fosbury, G. Naletto, R. Nilsson, T. Occhipinti, F. Tamburini, H. Uthas, L. Zampieri, QuantEYE, the quantum optics instrument for OWL, in The Scientific Requirements for Extremely Large Telescopes, ed. by P. Whitelock, B. Leibundgut, M. Dennefeld. IAU Symp., vol. 232 (Cambridge University Press, Cambridge, 2006), pp. 506–507 2. C. Barbieri, D. Dravins, T. Occhipinti, F. Tamburini, G. Naletto, V. Da Deppo, S. Fornasier, M. D’Onofrio, R.A.E. Fosbury, R. Nilsson, H. Uthas, QuantEYE, a high speed photometer pushed to the quantum limit for extremely large telescopes, J. Mod. Opt. 1, 1–10 (2006) 3. C. Barbieri, M. Belluso, S. Billotta, P. Bolli, G. Bonanno, V. Da Deppo, A. Di Paola, M. D’Onofrio, C. Facchinetti, E. Giro, S. Marchi, F. Messina, G. Naletto, T. Occhipinti, F. Pedichini, C. Pernechele, F. Tamburini, E. Verroi, M. Zaccariotto, P. Zoccarato, Aqueye, a singlephoton counting photometer for astronomy, invited talk, in Single Photon Workshop, INRIM Torino (September 2007) 4. C. Barbieri, S. Billotta, P. Bolli, G. Bonanno, A. Di Paola, C. Facchinetti, E. Giro, S. Marchi, G. Naletto, T. Occhipinti, C. Pernechele, E. Sain, M. Zaccariotto, P. Zoccarato, First results of AQUEYE, the fast multichannel photometer for the 182 cm telescope at Cima Ekar, in VI JENAM, Meeting European Astronomical Society, Erevan, Armenia, August 2007 5. D. Dravins et al., QuantEYE: The quantum optics instrument for OWL, in Proceedings from Meeting ‘Instrumentation for Extremely Large Telescopes’, held at Ringberg Castle, July 2005, ed. by T. Herbst (2005)

The E-ELT: A Chance to Measure Cosmic Magnetic Fields K.G. Strassmeier and I.V. Ilyin

Magnetic fields affect the evolution of structure in the Universe and drive solar and stellar activity which is a key to life’s origin and survival. However, our understanding of how cosmic magnetic fields form and evolve is currently very limited. Our close-up look at the Sun has enabled the creation of approximate dynamo models (which took us 100 years), but none yet predict the level of magnetic activity of the Sun or any other star. Therefore, major progress requires the understanding of solar, stellar and galactic magnetism in general and that in turn requires a population study. These population studies are now being thought about or some attempts are underway. The current 4–12 m class telescopes eventually provide the targets for future ELT magnetic-field observations of more complex and distant objects, even of the early Universe, and of more exotic objects like habitable extraterrestrial planets.

1 Cosmic Magnetic Fields: An E-ELT Science Highlight? Magnetism is one of the four fundamental forces in nature. Understanding the Universe is impossible without understanding magnetism, which plays an important role in the formation of celestial bodies and their evolution (e.g. [7]). Despite its importance, the structure, the evolution and the origin of a cosmic magnetic field are among the most important questions in modern astrophysics. At the E-ELT wavelengths, magnetic fields can only be measured indirectly, seen as the result of interaction with the propagating electromagnetic waves. Magnetic fields break the degeneracy of electronic configurations in an atom and splits its energy levels depending on the magnetic quantum number. As the result, depending on the atomic-level configuration and its coupling, a spectral line is split into a number of components and the light observed in each component has a different polarization state. The linearly polarized π -components originate from the transitions with equal magnetic quantum numbers, the other combinations result in blue- and red-shifted σ -components of the Zeeman profile which are circularly polarized in opposite sense depending on the orientation of the magnetic vector and their separation is a function of the field strength. The electromagnetic wave that K.G. Strassmeier () Astrophysical Institute Potsdam (AIP), An der Sternwarte 16, 14482 Potsdam, Germany e-mail: [email protected] A. Moorwood (ed.), Science with the VLT in the ELT Era, Astrophysics and Space Science Proceedings, doi: 10.1007/978-1-4020-9190-2_43, © Springer Science + Business Media B.V. 2009

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Fig. 1 Expected light curve of a Saturn around a distant Sun. From Dyudina et al.

corresponds to each of the Zeeman components emerges at the stellar surface and propagates through the atmosphere where it gains a different polarization state due to Faraday rotation as the result of ferromagnetic resonance. The change of the polarization state results also in a complex transformation of the absorption coefficients across the atmospheric layers (e.g. [2, 10, 13, 16]).

2 Spectropolarimetry of Earth-sized Extrasolar Planets Once a number of Earth-sized planet candidates were found, the main question will be how to further characterize them. One possible avenue is to try to separate the planetary light from the host-star light by means of either interferometry and/or coronagraphic (differential imaging) techniques. If these detections can be repeated at several wavelengths, a spectral energy distribution of the extra-solar planet is extractable as recently achieved with HST and Spitzer observations of the transiting planet of HD189333 (Udry, private communication), see also, e.g. Deming et al. [3]. First spectropolarimetry of the host star was obtained also just recently [11] and showed the presence of a magnetic field on the host star and the importance of it for the orbiting planet. The ultimate goal from the current perspective is to obtain a spectrum of the planet, similar to as we are taking spectra of the planets in our own Solar System. Obviously, this is a task for space projects like DARWIN, TPF a.o. but could be addressed by ELTs as early as, say, 2017, i.e. the projected first light of the E-ELT (while above space missions are for the 2020 to 2025 time scale). We point out a possible alternative route that stems from the original suggestion by Seager et al. [15] and others (e.g. [18]). It is based on the detection of the linearlypolarized signal of the light reflected off the surface of an orbiting planet (Fig. 2). It

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Fig. 2 Expected light (left) and linear-polarization (right) variations as a function of the orbital angle of a planet with respect to the viewing angle. From Seager et al. [15]

is expected to be δP /P ≈ 10−6 , i.e. 1 part per million of the total flux is polarized (assuming a solar twin as host star). A variant of this approach is to search for periodic line asymmetries in Stokes U and Q as compared to Stokes V, that is what we would want to propose. We expect that the Q&U spectral-line profiles of the host star would be contaminated by and modulated with the orbital period of the planet while the amplitude of the signal would be—additionally—dependent on the aspect angle of the planet with respect to the host star and the line of sight. Optimal extraction and signal-enhancement techniques like Least Squares Deconvolution [4] or Principal Component Analysis [1] will come into the play here. While the line asymmetry would be expected to be of the order of 10−4 even for a hot Jupiter, the radial velocity of the line centroid is rather large (proportional to the mass ratio star:planet). The confusion of this signal with the signature of the magnetic field of the (solar-type) host star is a major interpretation problem (as well as with the molecular-line contribution to the spectrum and its likely variability). However, the modulation of the linear signal should not repeat in circular polarized light if it is just due to reflected starlight while, if it is due to surface magnetic fields, all four Stokes parameters would be modulated equally. While there are many strings attached to such a concept (e.g., as mentioned, the molecular-line polarization), it is clearly worth to be addressed.

3 Performance Goal for an E-ELT Spectropolarimeter Any measurement of the polarization state will suffer from instrumental effects. The goal is to identify and to minimize these effects. Firstly, the phase delay in the retarder is a function of wavelength, even modern super-achromatic retarders [14] allow essentially just to minimize the effect. Secondly, even a small misalignment in the optical axis of the retarder with respect to the axis of the birefringent element induces a spurious polarization then measured in the spectra and possibly interpreted wrongly. And thirdly, misalignment of the optical axis of the birefringent element

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Fig. 3 The 2 × 8.4 m LBT will be soon the largest telescope with a spectropolarimeter. However, its light-gathering surface of 110 m2 will be dwarfed by the staggering 1380 m2 of the 42 m E-ELT. A colour version of this figure is available at dx.doi.org/10.1007/978-1-4020-9190-2_43

with respect to the reference parallactic axis of the target results in the corresponding error in the azimuthal angle of the derived linear polarization. More advanced analysis of polarimeter components with the so-called Müller-matrix method revealed additional terms which we consider to be important components in the contribution to the total uncertainty of the measured polarized spectra (e.g., [5, 8, 9]). Recently, Sun and Adlou [17] studied birefringence induced by mechanical stresses and anti-reflection coatings of optical lenses. The effect was found to be of order less than 10−7 , just a factor of ten below the signal from an extra-solar planet. Fortunately, most of these errors can be canceled by combining (subtracting) two circularly polarized spectra obtained with two orthogonal orientations of the retarder at 0◦ and 90◦ : the residual error is a combination of second order of the retardation and misalignment errors. Addition of the two spectra provides an estimate of the total residual error in the V profile. Similarly, in the case of linear polarization, a cross-talk between Stokes Q and U is induced by the misalignment error of the birefringent element itself and cannot be canceled-out or reduced with subsequent observations with other angles of the birefringent element. Optical cross-talk in linear polarization measurements may occur in the case of a slit spectrograph where the two polarized beams may partially overlap due to poor seeing, which is not the case though with a fiber-feeded spectrograph.

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Therefore, the proper calibration of a spectropolarimeter and the correct interpretation of the calibration results are essential to obtain the highest possible precision for magnetic-field measurements. Our aim in the case of CODEX [12] would be to achieve a relative precision of about 10−6...7 for Stokes QUV in spectra with the highest signal-to-noise ratio obtainable. The two pairs of polarized spectra are to be recorded at the same time with a single spectrograph and CCD. Two subsequent exposures are required to obtain circularly polarized spectra at two angles (0◦ and 90◦ ) of a quarter-wave plate. Two exposures with two angles (0◦ and 45◦ ) of the birefringent element (with the quarter-wave plate retracted) would be necessary to obtain the linearly polarized spectra (the other pair of angles at 90◦ and 135◦ may also be used for verification). Such a QUV mode covering the wavelength ranges 0.4–1 µm, 1–2.5 µm, and 2.5–5 µm feeding three spectrographs simultaneously would provide a limiting magnitude of V ≈ 16 mag or K ≈ 12 m ( G2 star) at δP /P = 10−2 . This shall allow access to bright quasars, globular cluster K dwarfs, brown dwarfs in nearby Open Clusters, Solar-System bodies as faint as Triton and all the way to the Galactic-Center O and AGB stars.

References 1. 2. 3. 4. 5. 6. 7. 8. 9. 10. 11. 12. 13. 14. 15. 16.

17. 18.

T. Carroll, M. Kopf, I. Ilyin, K.G. Strassmeier, Astron. Nachr. 328, 1043 (2007) E. Degl’Innocenti, M. Landolfi, Polarization in Spectral Lines (Kluwer, Dordrecht, 2004) D. Deming, J. Harrington, G. Laughlin et al., Astrophys. J. 667, L199 (2007) J.-F. Donati, M. Semel, B. Carter, D.E. Rees, A. Collier Cameron, Mon. Not. R. Astron. Soc. 291, 658 (1997) J.-F. Donati, C. Catala, J. Landstreet, Espadons, in Proceedings of the 5th CFHT Users Meeting, ed. by P. Martin, S. Rucinski (CFHT, 2001), p. 50 U.A. Dyudina, P.D. Sackett, D.D.R. Bayliss, S. Seager, C.C. Porco, H.B. Throop, L. Dones, Astrophys. J. 618, 973 (2005) B.M. Gaenslera, R. Beck, L. Feretti, New Astron. Reviews 48, 1003 (2004) A. Hofmann, K.G. Strassmeier, M. Woche, Astron. Nachr. 323, 510 (2003) I.V. Ilyin, Stellar atmosphere synthesis, PhD thesis, University of Oulu (2004) J. Jefferies, B.W. Lites, A. Skumanich, Astrophys. J. 343, 920 (1989) C. Moutou, J.-F. Donati, R. Savalle et al., Astron. Astrophys. 473, 651 (2007) L. Pasquini et al., CODEX Design Document (ESO, 2005) D.E. Rees, A gentle introduction to polarized radiative transfer, in Numerical Radiative Transfer, ed. by W. Kalkofen (Cambridge University Press, Cambridge, 1987) A.V. Samoylov, V.S. Samoylov, A.P. Vidmachenko, A.V. Perekhod, J. Quant. Spectrosc. Radiat. Transf. 88, 319 (2004) S. Seager, B.A. Whitney, D.D. Sasselov, Astrophys. J. 540, 504 (2000) J.O. Stenflo, Polarized radiation diagnostics of stellar magnetic fields, in Astrophysical Spectropolarimetry, ed. by F. Moreno-Insertis, F. Sánchez. XII Canary Islands Winter School (Cambridge University Press, Cambridge, 2002), p. 55 L. Sun, S. Adlou, SPIE 6288, 62890H (2006) D. Veras, P.J. Armitage, Astrophys. J. 620, L111 (2005)

The Experience from VISIR and the Design of an ELT Mid-infrared Instrument E. Pantin, R. Siebenmorgen, H.U. Käufl and M. Sterzik

Abstract VISIR is the VLT mid-infrared (mid-IR) Imager and Spectrometer. It provides data at high spatial and spectral resolution in the N (8–13 µm) and Q (16– 20 µm) atmospheric windows. VISIR observations have provided unique constraints on targets such as central regions of AGNs, or protoplanetary disks. Its successor on an ELT will provide data with unique spatial resolution (0.05 arcsec) and competitive sensitivity (50 µJy sources detectable in 1 hour), thus allowing e.g. to better characterize exoplanetary disks and exoplanets. We review here a selection of scientific contributions thanks to VISIR and then, in the light of the VISIR experience, we discuss the stakes for a mid-IR instrument at the European ELT.

1 VISIR Instrument VISIR has been commissioned in April 2004 on the VLT unit telescope #3 (aka Melipal) at Paranal. Since then it operates routinely to deliver mid-infrared (midIR) images and spectra in the atmospheric N-band ( = 300, 300 and 30 000) and Q-band ( = 1500 and 15 000) [1]. VISIR data are diffraction limited if the optical seeing is below ∼ 0.8 arcsec. The high spectral resolution mode remains unique in the southern hemisphere; warm H2 pure rotational lines at 8.02, 9.66, 12.28, 17.03 µm or the [NeII] line at 12.81 µm are of particular interest.

2 A VISIR Pot-Pourri of Scientific Results 2.1 Extragalactic Infrared Astronomy Obvious prime targets for VISIR are galaxies hosting active nuclei (AGNs). AGNs contain a central black hole surrounded by either accreted or ejected material. Studying these materials constrains the physical mechanisms acting at the level of the inner black hole. The close-by environment of the AGN (viscous accretion disk) is usually heated to 100 K and more, while being, in some cases obscured by colder E. Pantin () DSM/DAPNIA/SAp, CE Saclay, UMR 7158, France e-mail: [email protected] A. Moorwood (ed.), Science with the VLT in the ELT Era, Astrophysics and Space Science Proceedings, doi: 10.1007/978-1-4020-9190-2_44, © Springer Science + Business Media B.V. 2009

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Fig. 1 Left panel: NGC1068 as observed with VISIR (underlying image) and MIDI (red dot, green and blue contours). MIDI detects an intermediate scale (2 arcsec) mid-IR emission (green dashed contour); VISIR resolves it into several knots. Right panel: velocity resolved VISIR data of the active galaxy NGC7582 using the high-resolution spectrograph of VISIR at 12.8 µm ([Ne II] line). The velocity curve of inner obscured regions allows to determine the mass of the central black hole. A colour version of this figure is available at dx.doi.org/10.1007/978-1-4020-9190-2_44

dust. Mid-IR observations are thus well suited to study this kind of objects. The left panel of Fig. 1 shows VISIR high angular resolution observations (0.2 arcsec) of the very inner region (AGN) of NGC1068 [2]. VISIR observations providing information at intermediate spatial scales are complementary to MIDI1 [3]. NGC7582 is another AGN. VISIR observations in high spectral resolution ( = 25 000) allowed to derive an estimate of the mass of the black hole (5.5 · 107 M , Fig. 1, right panel [4]). Mid-IR measurements allow to extend a technique already used in the optical or near-infrared range, to strongly obscured galactic nuclei.

2.2 Low-Mass Companions and Protoplanetary Disks Indi B is the closest brown dwarf binary system known. The source separation is only 0.73 arcsec. The Spitzer space observatory is thus unable to discriminate the individual components. VISIR observations, although penalized by a much lower sensitivity, resolve them and put meaningful constraints on the physical characteristics of both components [5]. HD97048 is a young (3 Myr) intermediate-mass Herbig star in the Chameleon cloud, surrounded by a protoplanetary disk of dust and gas. VISIR direct imaging in the 8.6 and 11.3 µm PAH bands have resolved the structure of the disk for the first time ([6] and Fig. 2). The disk is thick, dense, and has a flaring upper surface as predicted by [7]. The measured flaring index (1.26 ± 0.05) is strikingly close to the 1 MIDI is the 10 µm instrument of the VLT interferometer providing spatial information on scales a small as 0.02 arcsec.

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Fig. 2 Left: the HD97048 dust disk resolved with VISIR imaging at 8.6 µm. The determination of the decentering of the isophotes as a function of distance to the star allows to precisely derive the flaring index of the disk. Right: the 0–0 S(1) H2 emission line detected in HD97048 (0.75 arcsec wide slit overplotted on left panel). A colour version of this figure is available at dx.doi.org/10.1007/978-1-4020-9190-2_44

expected value of 9/7 = 1.28 obtained when assuming hydrostatic equilibrium of the disk. VISIR high spectral resolution observations [9] confirmed the presence of warm gas (0.01 to 1 Jupiter masses) by detecting the H2 emission line at 17.035 µm (see Fig. 2). Maybe more important, the gas to dust mass ratio (3000 to 14 000) in the inner 35 AU deviates significantly from the canonical value of 100; maybe due to planet formation inducing dust depletion.

3 Experience from VISIR Based on our experience with VISIR, we realized that: 1. Relevant quality control parameters are the precipitable water vapor content, the conversion factor (ADU/Jy), the background level, and sensitivity. 2. The data themselves contain valuable information for calibration (e.g. sky lines for wavelength calibration). 3. Weather parameters are relatively stable over few hours; thus only a limited number of standard star interlaced observations are really needed for precise photometric calibration. VISIR suffers from various limitations and errors, such as detector striping, flatfieldability, stitching of consecutive low-resolution spectroscopy settings, or background errors. See [8] for a detailed description of these effects. Background errors are understood as an excess of noise at mid and low spatial frequencies probably due to a too small chopping frequency (0.25 Hz). Future mid-infrared instruments should benefit from large format (1k × 1k) more stable detectors and internal flat-field calibrators. The ELT will likely be unable to provide M2 chopping. Recalling that without chopping the precision needed for stability is of the order of

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Fig. 3 Limiting distances when characterizing a giant exoplanet in low-resolution ( = 100) spectroscopy on an ELT, as a function of planet mass. Two cases are considered: an age of 0.1 Gyr (more favorable because young planets are intrinsically hotter), and age of 1 Gyr. A colour version of this figure is available at dx.doi.org/10.1007/978-1-4020-9190-2_44

10−4 –10−5 , the study of alternatives for background cancellation techniques is unavoidable and already ongoing.

4 Observing in the Mid-infrared Range on an ELT The expected sensitivity2 of a mid-IR instrument on an ELT opens some new perspectives e.g. in the field of exoplanetary science. Low-resolution spectroscopy of giant planets on orbits of a few AUs around the parent star becomes feasible to better characterize atmospheric composition and mass (cf. Fig. 3). In case “super-Earth planets” with radii of the order of Jupiter around nearby stars exist, they could be detected directly up to distances in the range 5–10 pc. Finally, an ELT mid-IR instrument, given its sensitivity and spatial resolution, is a perfect tool to accurately study protoplanetary and debris disks. In particular, structures such as gaps could be revealed. The performance of a mid-IR instrument depends significantly on the observatory site. While N band observations are less sensitive to altitude, both Q band sensitivity and spectral coverage are highly dependant on it (e.g. at 24 µm the atmospheric transmission for 4500 m altitude is twice higher when compared to the Paranal). In order to be competitive with cryogenic infrared space observatories, 2 Mid-IR observations are back-ground-noise limited, so s/n scales with the square of the telescope diameter.

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high spatial and spectral resolution shall be the drivers for ground-based mid-IR instruments. The ELT diffraction limit can be maintained in the mid-IR range at a moderate cost, i.e. using relatively simple adaptive optics system when compared to near-IR requirements. Most of the science cases described in this paper deal with targets angularly very close to stronger sources (stars, central engines of AGNs); thus differential observing techniques are a must. This implies very high constraints on the temporal stability of the components (telescope, AO front-end, instrument). On an ELT, a typical 1 Jy source would be 3 times the background level; spurious saturation effects would probably then arise. Stability and accuracy can be recovered when using coronographic devices or simultaneous differential imaging.

5 Conclusions and Perspectives All science cases, requiring either high spatial or spectral resolution, including high contrast, can be uniquely addressed by ground-based mid-IR astronomy. Although mid-IR data are generally more demanding in terms of data reduction efforts, the results can lead to unique scientific break-throughs, sometimes unachievable at any other wavelength. A mid-IR ELT instrument will combine high spatial resolution with meaningful point source sensitivity (0.05 arcsec/50 µJy in 1 h). The performance of such an instrument could be largely increased w.r.t. contrast and spatial resolution, especially involving differential observations, if devices such as four quadrants coronographs, or dual-band imaging would be implemented. With these provisions such an instrument would be not only “the perfect machine” to study dusty disks, but also the direct characterization of exoplanetary systems becomes feasible.

References 1. 2. 3. 4. 5. 6. 7. 8.

P.O. Lagage, J.W. Pel, M. Authier et al., Messenger 117, 12 (2004) E. Galliano, E. Pantin, D. Alloin et al., Mon. Not. R. Astron. Soc. 363, L1 (2005) A. Poncelet, C. Doucet, G. Perrin et al., Astron. Astrophys. 472, 823 (2007) M. Wold, M. Lacy, H.U. Käufl et al., Astron. Astrophys. 460, 449 (2006) M. Sterzik, E. Pantin, M. Hartung et al., Astron. Astrophys. 436, L39 (2005) P.O. Lagage, C. Doucet, E. Pantin et al., Science 314, 621 (2006) S.J. Kenyon, L. Harmann, Astrophys. J. 323, 714 (1987) E. Pantin, L. Vanzi, U. Weilenman, Some (little) things about VISIR, in The 2007 ESO Instrument Calibration Workshop (Springer-Verlag, in press) 9. C. Martin-Zaïdi, P.O. Lagage, E. Pantin et al., Astrophys. J. 666, L117 (2007)

HARMONI: A Narrow Field Near-infrared Integral Field Spectrograph for the E-ELT Matthias Tecza, Niranjan Thatte, Fraser Clarke and David Freeman

1 Introduction We present a concept for a narrow field near-infrared (NIR) integral-field spectrograph (IFS) for the E-ELT called HARMONI—the High Angular Resolution Monolithic Optical and Near-infrared Integral field spectrograph. This IFS is ideally suited to exploit the near diffraction limited performance of the adaptive optics systems at NIR (and red visible) wavelengths of the E-ELT at first light. Although such an E-ELT instrument does represent a substantial increase in scale, no significant technology development is needed to realise this instrument. The instrument will be robust, with few moving mechanisms, designed to act as a workhorse in the early days of the E-ELT. In the era where the James Webb Space Telescope will already be operational, the clear advantages of the E-ELT are its larger collecting area and higher spatial resolution. A narrow field NIR IFS is ideally suited to exploit these capabilities, making HARMONI a unique first-light instrument which will have a dramatic impact in several key areas of observational astrophysics research; from the highest redshift galaxies, through AGN and supermassive black holes, to star and planet formation in our own galaxy.

2 Design Study Funded by the UK Science and Technology Facilities Council we conducted a small pre-study of potential designs for a narrow-field NIR IFS. This study was motivated by the fact that such an IFS, although identified as a high-priority potential first-light instrument in the instrumentation plan for the E-ELT [1], was noticably not part of either the suite of small studies conducted under the aegis of the EC financed FP6 ELT design study, nor the OWL instrument design studies, carried out by ESO and its member states. We used our extensive experience with the design of SINFONI at M. Tecza () Astrophysics, University of Oxford, Denys Wilkinson Building, Keble Road, Oxford, OX1 3RH, United Kingdom e-mail: [email protected] A. Moorwood (ed.), Science with the VLT in the ELT Era, Astrophysics and Space Science Proceedings, doi: 10.1007/978-1-4020-9190-2_45, © Springer Science + Business Media B.V. 2009

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the VLT and SWIFT for the Palomar 200-inch Hale telescope to arrive at a potential workable design for such an instrument which we present here. In the meantime ESO released a call for proposals for a Phase A study of a Single Field, Wide Band Spectrograph to which we responded leading a consortium of British, French, and Spanish universities and institutes and were successful with our bid.

2.1 Initial Requirements Our design, elaborated below, is built upon a set of baseline instrument parameters. Essentially, these are derived by scaling the specifications of the successful SINFONI instrument to the E-ELT but will be refined according to the requirements of the wide variety of science cases where we expect HARMONI to have a large impact. However, this preliminary set enables us to test the feasibility of potential designs for the NIR IFS. Furthermore, it proves a rough estimate for the volume, mass and level of complexity of the instrument. For our candidate design, we adopted the following instrument parameters: • Total number of spaxels ∼ 16 000. • Wavelength range 0.85–2.45 µm. • Largest spaxel scale of 0.05 (Nyquist sampled for GLAO), and a smallest spaxel scale of 0.005 (for LTAO/MCAO). • Rectangular, 2 : 1 aspect ratio, field of view, with 88 × 176 spaxels, or 8.8 × 4.4 at the largest, and a tenth of that at the finest spaxel scale. • Spectral resolution ∼ 4000, allowing one NIR band (J, H or K) to be dispersed on to 2000 detector pixels. • Instantaneous wavelength coverage λ/λ ≈ 0.25, corresponding to 2000 detector pixels per spectrum. As HARMONI could be a first-light instrument, we also imposed that the instrument should utilise proven technology to avoid time-expensive R&D programs and to minimise risk.

2.2 Integral Field Unit Given the need for simultaneous spectral coverage encompassing at least one NIR atmospheric window (J, H or K) and possibly more, the candidate design presented below uses an image slicer for its integral field unit. Achieving such a large simultaneous wavelength coverage with a lenslet array spectrograph would be very difficult, and a slicer design also makes more efficient use of a given detector real estate. The latter could be critical given the large expected outlay in detectors required by this instrument.

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Demagnifying Image Slicer The design of the HARMONI slicer is based on the de-magnifying image slicer [4] that is used in the Oxford SWIFT spectrograph [3] where the lens mosaic has been replaced by a set of re-imaging mirrors, so as to achieve a fully achromatic design. The slicer is simple to manufacture, as both the slicing and the pupil mirrors are flat and with identical spherical re-imaging mirrors. This design also ensures extremely high throughput and fill factor, as the slicer can be fabricated as a monolithic unit, using optical contacting techniques. The slicer is composed of four identical sub-units. Unlike the SWIFT slicer, each slicer sub-unit has an off-axis angle added to both the pupil and slicing mirrors, so that the re-imaged slit is above (or below) the slicer stack, which allows the incoming light to reach the slicer stack unvignetted. This geometry has the added advantage that it provides the means to stack the sub-units into a single large slicer stack, without the need for field splitting optics. Figure 1 shows a perspective view of the entire HARMONI image slicer. The design provides a number of clear advantages: • Flat slicing and pupil mirrors can be fabricated more easily and with superior optical surface finish than more complex optical shapes if the image slicer is built using glass materials which is especially important for operation at visible wavelengths. • There is no need for field splitting and relay optics, making a distinct improvement to the throughput. The throughput goal for HARMONI is 35% on average, including the detector Q.E.; similar to that achieved by SINFONI, requiring as few optical components in the path as possible. • The slicer is fully achromatic, allowing operation over an extended wavelength range.

Fig. 1 Left: Perspective view of a single HARMONI image slicer consisting of 44 slices each 88 pixels long. Each individual slice is 1.3 mm wide and 57.2 mm long. With a magnification of 0.1 the total length of the exit slit is ≈ 266 mm. Right: Perspective view of the entire HARMONI slicer assembled from four identical slicers as show on the left. The top two slicers produce their exit slits above, the bottom two slicers below the incoming beam. A colour version of this figure is available at dx.doi.org/10.1007/978-1-4020-9190-2_45

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• Every slice has an identical optical train through the slicer, thus ensuring uniformity of aberrations, and easy instrument characterisation. • The fill factor is very high, both on the sky, and at the detector, with over 95% of all detector pixels used to detect source/sky photons. • The 2 : 1 aspect ratio of the FoV is well suited for nodding-on-IFU for accurate background subtraction.

2.3 Preoptics The slicer design requires an entrance pupil that lies 300 mm behind the slicer focal plane. The preoptics match this pupil location and also includes the scale changer which delivers the coarse GLAO and the diffraction-limited LTAO/MCAO imagescales. The scale changing is done by a system of achromatic lenses. Additionally, the preoptics contain cylindrical mirrors that provide a 2 : 1 anamorphic magnification, independent of image scale. This ensures Nyquist sampled spectra are obtained in a single exposure, removing the need for spectral dithering. Figure 2 shows the preoptics in the narrow field mode.

Fig. 2 Perspective view of the HARMONI preoptics in the narrow field mode. A group of four scale changing lenses provides a scale of 5 mas per pixel while two cylindrical mirrors create a 2 : 1 anamorphic magnification. In the wide field mode an achromatic doublet replaces the four lenses of the narrow field mode. A colour version of this figure is available at dx.doi.org/10.1007/978-1-4020-9190-2_45

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Fig. 3 Perspective view of the HARMONI image slicer feeding the spectrograph modules each comprising of a three mirror collimator, two fold mirrors, a transmission grating, and a six lens camera. Only the upper two of the four spectrographs are shown, the other two spectrographs are arranged symmetrically below the slicer. A colour version of this figure is available at dx.doi.org/10.1007/978-1-4020-9190-2_45

2.4 Spectrograph Using the slicer concept described above, we have designed four identical back-end spectrograph modules that will collimate, disperse, image and detect the light from the four pseudo-slits formed by the image slicer. Each spectrograph will provide ∼ 4000 spectra, each about 2000 pixels long, using two HAWAII 2 detector arrays. Figure 3 shows a schematic view of the image slicer sub-system, together with the four spectrographs. The f/6 collimator is a 3 mirror design, comprising 2 spherical and 1 aspherical mirror, with a focal length of 1500 mm, providing an elliptical beam of 120 × 240 mm. The dispersing element is a transmission grating. This permits a compact design with the camera placed close to the grating. The camera has a focal length of 420 mm and a f-ratio of f/1.8. It consists of 6 lenses with diameters between 150–300 mm and provides a ±5◦ field.

References 1. S. D’Odorico, E-ELT Instrumentation. Presented at the Conference Towards the European ELT, Marseille (2006) 2. F. Eisenhauer et al., Proc. SPIE 4841, 1548 (2003) 3. M. Tecza et al., New Astron. Rev. 49, 647 (2006) 4. M. Tecza et al., Proc. SPIE 6273, 62732L (2006)

Which Synergies Between LBT/LINC Nirvana and Future ELTs? L. Labadie, T.M. Herbst, S. Egner, M. Brix and M. Kürtser

We give in this paper a short overview of the LBT and LINC-Nirvana interferometer. We present two specific cases of the instrumental issues that were addressed within the project, namely the requirement for a wide field AO correction and the importance of measuring the flexure and vibrations on the telescope. Their importance in the context of future ELTs is also discussed.

1 The Large Binocular Telescope: A Pre-ELT Facility The Large Binocular Telescope (LBT) installed on Mount Graham, Arizona is a unique facility, which supports two 8.4-m primary mirrors jointly moved by an alt-azimuth mount. The two beams of the LBT can be combined interferometrically thanks to the LINC-Nirvana instrument, which is currently under integration at MPIA, Heidelberg. In its coherent combination configuration, LINC-Nirvana will be equivalent to a 23-m telescope, i.e. among the largest telescopes worlwide with half of the diameter of the future 42-m European ELT. The synthesized telescope LBT naturally appears as a pre-ELT facility. The technical issues inherent to this unique facility will have to be faced again to some extent in the context of ELTs. In the next sections, we detail a number of key points which were considered as critical for LBT/Linc-Nirvana and replace them into a larger context.

2 Telescope and Instrument Overview 2.1 The Large Binocular Telescope We briefly review here some of the main capabilities of LBT/LINC-Nirvana. Further details can be found in previous papers [1–3]. The binocular telescope is installed in L. Labadie () Max-Planck Institut für Astronomie, Königstuhl, 17, 69117 Heidelberg, Germany e-mail: [email protected] A. Moorwood (ed.), Science with the VLT in the ELT Era, Astrophysics and Space Science Proceedings, doi: 10.1007/978-1-4020-9190-2_46, © Springer Science + Business Media B.V. 2009

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Fig. 1 Left: general view of the LBT dome on Mount Graham, Arizona. The enclosure has a cubic shape. On this image, only the blue prime focus camera is installed on the left eye. Right: the binocular telescope inside the dome. The two 8.4-m mirrors are 14.4-m center-to-center and supported by an alt-azimuth mount. A colour version of this figure is available at dx.doi.org/10.1007/978-1-4020-9190-2_46

a 25 m3 cubic dome that opens a further 20 m in the horizontal direction for observations (white arrows in Fig. 1). The LBT itself is a compact telescope/interferometer with its two 8-m primary mirrors telescopes (see Fig. 2) that will host 10 instruments. For comparison, 13 instruments including the interferometers MIDI, AMBER and PRIMA are installed on the VLT site on Paranal. The center-to-center separation between the two dishes is 14.4 m, which provides a partially filled aperture of 23 m of diameter. The Gregorian design of each telescope gives access to the prime foci to perform seeing-limited wide field imaging in the visible. The relatively small f-number of the primary (∼ f/1.4) gives a large field-of-view over a small area: the two prime focus camera, LBC Blue and LBC Red, can image a ∼ 27 × 27 FOV on a mosaic of CCD with ∼ 100 mm side.

2.2 LINC Nirvana on LBT LINC-Nirvana is a Fizeau interferometer with the images from each aperture coherently combined on the focal plane of the instrument. If the condition of homotheticity between the input and output pupils [4] is respected,1 every point-like source will produce a interferometric fringe pattern2 at the spatial frequency D/λ modulated by the 8.4-m PSF. The advantage of the adopted combination scheme is the wider fieldof-view that can be preserved compared to classical interferometers. LINC-Nirvana functions as a true imager delivering high angular resolution over a 10 × 10 fieldof-view. The diffraction-limited resolution is, respectively, 10, 15 and 20 mas in the 1 This

condition requires that the output pupil is a directly scaled version of the input pupil that is it keeps the ratio between apertures size and separation. 2D

is the diameter of the synthesized aperture and λ the wavelength.

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J, H and K bands. However, as for any large telescope, diffraction-limit operation requires a high performance adaptive optics system to compensate for the atmospheric turbulence.

2.3 The Adaptive Optics System on LINC-Nirvana Conventional adaptive optics systems are able to correct the wavefront corrugations with a good Strehl ratio only over a small field-of-view of a few arcseconds. This is a limiting factor for science programs that aim at studying large scale structures. The alternative is to implement multi-conjugated adaptive optics (MCAO), which makes use of several guide stars, and possibly several deformable mirrors, to analyze the atmospheric turbulence over a larger field-of-view. Several approaches are possible for implementing MCAO (Ground Layer AO, Layer-Oriented AO, etc.) and a starting point on this technique can be found in Beckers (1993) or Ragazzoni et al. (2002) [6, 7]. This technique is obviously more complex than single AO systems, but in the case of LINC-Nirvana using MCAO is clearly mandatory to obtain a good fringe contrast over a larger FOV. The MCAO system on LINC-Nirvana is based on eight to twelve pyramid wavefront sensors operating in the visible [5]. The correction of the ground layer turbulence can be achieved over 6 FOV by coupling the WFS to the deformable secondary mirror. The internal 2 FOV that is used for interferometry is also corrected for higher atmospheric turbulence, namely at ∼ 4 and 10 km, by conjugating the DM to the appropriate altitude. Such an adaptive optics system for wide-field correction will clearly be a major step towards large telescope operation. Meanwhile, the technique was demonstrated on-sky recently with the first spectacular K-band images from MAD at the VLT which delivered a high Strehl correction over a 2 × 2 FOV [8].

2.4 Constraints on the Telescope Image Quality for the Coherent Combination High resolution imaging will be possible in Fizeau mode only if the PSFs originating from each telescope are precisely superposed on the detector plane. The relative tilt between the PSFs can be induced by differential flexure that must then be controlled to some level. The plots of Fig. 2 show a simulation of the effect of the relative tilt on the PSF superposition. The effect is, as expected, more pronounced at shorter wavelengths and, the advantage of a 22.8-m aperture is lost as soon as the tilt exceeds few pixels3 Table 1 gives the specifications on the tilt in J and K bands. These specifications appear quite strong, however a high performance tip/tilt system can compensate for the relative shift of the PSFs. 3 The

pixel scale is 5 mas/pixel.

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Fig. 2 Simulation on the effect of relative shift between PSFs Table 1 J and K band specifications on the relative shift between the PSFs. The shift in the focal plane is considered before entering the collimator optics. The plate scale giving the shift on the detector is 600 µm/arcsec Specification J band

Specification K band

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Flexure of the telescope structure Another critical aspect of the LBT is the flexure constraints that the different components have to support. Such a large structure can suffer from vibrations at different frequencies, which can be particularly severe for interferometric combination. Although a piston mirror will compensate for residual OPD, it is necessary to understand the extent of the different external constraints. A vibration campaign has been conducted at the LBT to measure the vibrations on the primary mirror. A second campaign will allow us to measure the displacement between primary and secondary as a crucial step for interferometry. Figure 3 shows the acceleration spectrum of the primary, showing three clear peaks below 80 Hz. At this stage, we know that this remains within the frequency range of the piston mirror.

3 Synergies with ELTs Some of the issues presented in the previous sections will probably have to be addressed for larger facilities like the E-ELT. The use of wide field adaptive optics is mandatory on an ELT as soon as science cases requiring large FOV are considered. But any type of new technology (MEMS, etc.) planned for ELTs will become a valuable add-on only once we have understood most of the aspects linked to the first

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Fig. 3 Measured acceleration spectrum for the LBT primary. A colour version of this figure is available at dx.doi.org/10.1007/978-1-4020-9190-2_46

generation of MCAO systems. From this point of view, LINC-Nirvana represents the first challenging “testbed” in the coming years. Among others, the vibration issue will be another key point for future extremely large telescopes. In the last few years, it was shown that unwanted vibrations are very damaging for interferometric facilities like Keck or VLTI. Thus, the question was then tackled at LBT in a preventive way. Because any ELT primary mirror will be composed of many segments, we can already anticipate that flexure and vibration issues will be part of the challenges to be addressed. Finally, the size of the LBT itself has implied a new approach in terms of telescope mounting and dome construction never considered before. The importance of the project is also very demanding in terms of management. We presented here only few aspects of the instrumental constraints of the LBT and its interferometric Fizeau imager LINC-Nirvana. However, in the context of future ELTs, positive synergies can certainly result from the experience gained at the different levels presented here.

References 1. T.M. Herbst, P.M. Hinz, Interferometry on the large binocular telescope, in New Frontiers in Stellar Interferometry, Proc. SPIE, vol. 5491 (2004), p. 383 2. T.M. Herbst, R. Ragazzoni, A. Eckart, G. Weigelt, The LINC-NIRVANA interferometric imager for the large binocular telescope, in Ground-based Instrumentation for Astronomy. Proc. SPIE, vol. 5492 (2004), pp. 1045–1052 3. P. Bizenberger, E. Diolaiti, S. Egner, T.M. Herbst, R. Ragazzoni, D. Reymann, W. Xu, LINCNIRVANA: Optical design of an interferometric imaging camera, in: Ground-based and Airborne Instrumentation for Astronomy. Proc. SPIE, vol. 6269 (2006), p. 62690D 4. J.R.P. Angel, J.M. Hill, P.A. Strittmatter, P. Salinari, G. Weigelt, Interferometry with the large binocular telescope, in Astronomical Interferometry. Proc. SPIE, vol. 3350 (1998), pp. 881–889 5. R. Ragazzoni, T.M. Herbst, W. Gässler et al., A visible MCAO channel for NIRVANA at the LBT, in Adaptive Optical System Technologies II. Proc. SPIE, vol. 4839 (2003), pp. 536–543 6. J.M. Beckers, Annu. Rev. Astron. Astrophys. 31, 13 (1993) 7. R. Ragazzoni, E. Diolaiti, J. Farinato et al., Astron. Astrophys. 396, 731 (2002) 8. ESO press release 19/07 (30/03/2007) at http://www.eso.org/public/outreach/press-rel/pr-2007/ pr-19-07.html

TMT Science and Instruments David Crampton, Luc Simard and David Silva

Abstract To meet the scientific goals of the Thirty Meter Telescope Project, full diffraction-limited performance is required from the outset and hence the entire observatory is being designed, as a system, to achieve this. The preliminary design phases of the telescope and the first light adaptive optic facility are now approaching completion so that much better predictions of the system performance are possible. The telescope design and instrumentation are summarized in this presentation, with a brief description of some of the scientific programs that are foreseen.

1 Introduction The TMT Project is rapidly moving towards construction of a 30 m telescope that is being designed from the outset as a system that will deliver diffraction limited images at wavelengths longer than 1 micron. Many science programs will thus realize the D4 advantage in point source sensitivity inherent in such a telescope. Although a primary diameter, D, greater than 30 m would offer an even larger gain, our analyses indicate that 30 m is the optimum balance between science benefit, cost, technological readiness and schedule at the present time. Using components that are currently, or soon-to-be available, it is possible to achieve images with a high Strehl ratio at wavelengths greater than 1 µm with a 30 m telescope and, although the instruments are challenging, they are feasible. The telescope will be a Ritchey-Chretien design with a f/1 primary mirror. The latter will be composed of 492 1.4 m segments. Instruments will be located on two large Nasmyth platforms, addressed by an articulated tertiary mirror. This will enable rapid (less than 10 minutes) switching between on-sky observations of targets with different instruments (less than 5 minutes with the same instrument). Adaptive optic (AO) systems, including a laser guide star facility, are being integrated into the observatory system, with plans to employ multi-conjugate AO (MCAO), multi-object AO (MOAO), ground layer AO (GLAO), mid-IR AO (MIRAO) and “extreme” AO (ExAO). The telescope will produce a 20 diameter field for seeing limited observations from the ultraviolet (0.3 µm) to the MIR (28 µm). Aerodynamic studies demonstrate excellent performance of the Calotte style enclosure, providing venting to minimize image degradation within the enclosure while not inducing unacceptable wind buffeting of the telescope itself. David Crampton () TMT Project Office, 2636 East Washington Blvd, Pasadena, CA 91107, USA e-mail: [email protected] A. Moorwood (ed.), Science with the VLT in the ELT Era, Astrophysics and Space Science Proceedings, doi: 10.1007/978-1-4020-9190-2_47, © Springer Science + Business Media B.V. 2009

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The design and development phase of the TMT project will reach completion in early 2009. Science operations are expected to begin in the last half of the next decade.

2 Instrument Suite The TMT Scientific Advisory Committee (SAC) identified a comprehensive suite of eight instruments required to tackle the science that they envisage for the first decade of operation. The proposed instruments span the discovery space in wavelength, spatial resolution, spectral resolution (R) and field-of-view/slit length. They also define a number of important TMT subsystem parameters such as the physical sizes and weights of instruments that the observatory should be designed to accommodate. Six of the instruments use built-in AO systems or use NFIRAOS, the facility MCAO system, to exploit the diffraction-limited capability of TMT. The other two are seeing-limited but could utilize AO to improve their observing efficiency. The instruments also exploit the entire wavelength range of TMT, from 0.31 to 28 µm; they include a high contrast instrument; and instruments with a wide variety of field sizes, up to 20 in diameter. Thus the instrument suite is representative and suitable for defining general instrument requirements that should provide enough flexibility to accommodate future instruments. Feasibility studies of the instruments were carried out in 2005–2006. Nearly two hundred scientists and engineers at forty-six US, Canadian and French institutions were involved in these studies, which were reviewed by panels of international experts. Most importantly, perhaps, at this stage, these studies demonstrate that the instruments are feasible, albeit challenging. The science cases and operational concept documents of these studies highlight and document the tremendous scientific potential of TMT. A diagram demonstrating how the instruments could be arranged on the Nasmyth platforms is shown in Fig. 1.

2.1 Early Light Instruments The instrument suite has been divided by the SAC into “early light” and “first decade” instruments for a variety of pragmatic reasons, mostly to do with funding constraints, commissioning practicalities, and technological readiness. The early light suite consists of IRIS, a NIR integral field spectrograph and imager working at the diffraction limit, WFOS, a widefield multiobject spectrograph, and IRMS, a multislit NIR spectrograph and imager fed by the facility MCAO system. While bringing the early light instruments on-line is clearly the top priority of the TMT instrumentation program, the ultimate goal is to bring the full suite into operation over the first decade of operations. Many global observatory design decisions and choices were therefore made with this in mind. The very ambitious WFOS originally requested by SAC has been scaled back to reduce cost, risk and commissioning complexity and make it suitable for early light. Likewise, the early light IRIS

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Fig. 1 Diagram showing how the instruments, located on two large Nasmyth platforms, will be addressed by the articulated tertiary mirror. See text for meaning of the acronyms. IRMS will be located on the side port of NFIRAOS (where NIRES-B is shown) during the first years of operation. The “instrument” labeled APS is the system used to align and phase the 492 segments of the primary mirror and the optical system as a whole. A colour version of this figure is available at dx.doi.org/10.1007/978-1-4020-9190-2_47

configuration has also been kept as simple as possible yet still retaining the ability to meet its key science objectives.

InfraRed Imaging Spectrograph (IRIS) IRIS, located on the bottom port of NFIRAOS, will be able to conduct diffractionlimited imaging and integral field spectroscopic observations in the 0.8–2.5 µm wavelength region. The SAC has consistently ranked IRIS as the top instrument priority for TMT, partly because of its ability to utilize the exquisite images delivered by NFIRAOS (which will be only 7 milliarcsec in the J band). The imager will have a field of view of at least 15 and IRIS is expected to incorporate several different plate scales for integral field spectroscopy yielding fields of view up to 2 . Astrometric measurements with precisions of order 100 microarcsec are foreseen. The versatility of IRIS allows it to address a very broad range of science problems. Some of the prime science drivers include observations of the first luminous objects in the Universe, supermassive black holes in the cores of distant galaxies, relativistic effects at the Galactic Center, resolved stellar populations in the crowded fields of galaxies out to the Virgo galaxy cluster, and high-resolution imaging and spectroscopy of planets and satellites in the Solar System.

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InfraRed Multi-slit Spectrograph (IRMS) Some form of multi-object NIR spectroscopy is another essential capability for early light. Understanding the so-called “First Light” objects in the Universe, the origin and evolution of galaxies and other objects detected by JWST and ALMA will require spectra of many extremely faint objects in the NIR, and multiplexing will thus be essential. Although a fully multiplexed deployable IFU system using MOAO was judged to be too risky and expensive for an early light instrument, fortuitously a clone of the MOSFIRE multislit instrument, currently being built for Keck, provides a very exciting interim capability. Although MOSFIRE will be a seeing-limited instrument for Keck, it can be easily adapted for use in an AO mode with NFIRAOS, providing an exceedingly powerful capability for TMT at low risk and modest cost. When optimized for widefield mode, NFIRAOS will deliver images to IRMS that will produce almost an order of magnitude gain in encircled energy within narrow (160 mas) slits over the entire of 2 diameter field.

Wide-Field Optical Spectrograph (WFOS) A number of key TMT science programs will be best conducted with a powerful survey instrument operating at optical wavelengths. These programs include the determination of the baryonic power spectrum, the tomography of the intergalactic medium, the determination of the dynamical states of stellar populations in nearby galaxies, the dark matter distribution in elliptical galaxies, and the star formation history in local galaxies. A high multiplexing capability (several hundreds of spectra) over a relatively large field-of-view is required to sample large cosmological volumes with sufficient target density to probe the range of desired physical scales. The spectrograph should also provide good image quality and moderate spectral resolutions over the wavelength range 0.31–1.1 µm.

2.2 First Decade Instruments InfraRed Multiple Object Spectrometer (IRMOS) IRMOS, as envisaged by SAC, is simultaneously perhaps the most ambitious and the most exciting of the TMT instruments. Its goal is to deliver 2D integral field spectroscopy of many objects over a 5 field of regard using MOAO to deliver quasidiffraction-limited spatial resolution in the NIR. Key IRMOS science includes the physical properties of galaxies (internal velocity fields, star formation rate, chemical abundance) at the epoch of peak galaxy assembly (z ∼ 2–3) and the physical conditions in star-forming regions. IRMOS will allow many outflows around newlyforming stars to be resolved and their interaction with the interstellar medium to be studied.

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Near-Infrared Echelle Spectrometer (NIRES) NIRES is a straightforward NIR echelle spectrograph for the 1–2.4 µm wavelength range that also fits behind NFIRAOS. In fact, NIRES could effectively be a clone of existing or planned diffraction-limited NIR spectrographs (e.g. Keck NIRSPEC). The NIRES feasibility study clearly identifies the enormous potential of such an instrument on TMT for programs such as understanding the physics of gamma-ray bursters, probing the intergalactic medium at high redshifts, and delivering precision radial velocities of late type stars for planet searches.

High Resolution Optical Spectrometer (HROS) High-resolution optical spectrographs have occupied center stage in recent years thanks to a range of exciting work: the very productive Doppler searches for exoplanets, the measurements of chemical abundances in absorbing intergalactic matter along the lines-of-sight to distant quasars and the surveys of the outer reaches of the Milky Way in search of the most metal-poor stars. HROS is fundamentally a seeing-limited, high-resolution optical spectrometer, although it is recognized that its performance could be enhanced by the use of LTAO (Laser Tomography AO). HROS will provide a spectral resolution of R = 50 000 (1 arcsec slit) or R ≥ 90 000 (image slicer) over the wavelength range: 0.31–1.0 µm.

Mid-Infrared Echelle Spectrometer (MIRES) MIRES brings diffraction-limited, high spatial-resolution imaging and highresolution spectroscopy in the thermal infrared (5–25 µm) to the TMT instrumentation suite. The MIRES science case features many fascinating objectives: the origin of stellar mass, the exploration of the inner parts of protoplanetary disks, astrochemistry, and the deposition of pre-biotic molecules onto planetary surfaces. The spectral resolution of MIRES will be 5000 < R < 100 000 with a diffraction-limited slit.

Planet Formation Instrument (PFI) The Planet Formation Instrument (PFI) is focused on the direct detection and characterization of extrasolar planets. PFI is unique among the instruments in that it places significant requirements on the telescope optics (primary pupil shape, secondary support structure, segment edge and reflectivity) and vibrational environment, and these requirements have been factored in the telescope design. PFI will build very strongly on the heritage being gained by the “planet finder” instruments that are currently being designed for Gemini and the VLT. The contrast requirements are 10−8 at 50 milliarcsec, goal of 10−9 at 100 milliarcsec (parent star magnitude I < 8).

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3 TMT Science It is obviously extremely difficult to foresee what scientific programs will be carried out by TMT and many of the discoveries made will be unanticipated or even unimagined at present. However, the SAC has provided examples of potential programs to guide the development of the observatory. Some of these were mentioned above in the descriptions of each instrument but, in practice, several instruments and attributes of the Observatory will contribute to solving most of the major questions. Two examples of this, and how TMT might be used to address them, are summarized here.

Fundamental Physics and Cosmology Various probes of Dark Energy will be undertaken by TMT, including precise measurements of the expansion history and power spectrum of the Universe at low and high redshifts using supernovae. The latter will also be used along with Gamma Ray Bursters (GRB) and Super Massive Black Holes to study the physics of extreme objects. Problems such as determining whether there are variations in the fundamental physical constants require high spectral resolution observations with signal-to-noise that only an extremely large telescope can provide. The nature of Dark Matter will be examined using probes such as the determination of 3D orbits of stars near the Galactic Center, measurement of the baryonic power spectrum through tomography of the InterGalactic Medium (IGM), and the determination of the kinematics of stars in dwarf galaxies using precision spectroscopy with WFOS, HROS, NIRES and IRIS, and precision astrometry with IRIS and WIRC. For studies of the sources of first light and cosmic reionization, TMT will have strong synergy with JWST and future 21 cm surveys. Although it is anticipated that JWST will be able to detect the brightest such sources, TMT should go at least one magnitude fainter and perhaps much more, depending on their size. At slightly lower redshift, IRIS, IRMOS and NIRES will study detailed properties of first galaxies and their influence on the IGM. WFOS and IRMS spectra of large samples of distant galaxies and the intervening IGM, coupled with 3D spectroscopy of samples of galaxies with IRMOS will elucidate how galaxies acquire gas, how star formation proceeds, and the effects of supernovae, SMBH and active star formation on the formation of galaxies. Some of these projects gain not only through the huge sensitivity gain provided by the diffraction limited images delivered by TMT but also because study of fields in the galactic center and nearby galaxies are fundamentally limited by confusion (crowding). In addition, all projects will gain from the rapid target acquisition and rapid switching between instruments that is being designed into the system. Observation of GRBs, especially, will benefit from the fact that TMT will rapidly slew and acquire targets, set up active and adaptive optics systems and be ready to begin observation with any instrument in less than 10 minutes.

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Formation of Stars and Planets HROS and NIRES will be used to detect and measure orbits and masses of exoplanets. PFI will be able to directly image and characterize a complementary sample of planets. MIRES, using the high angular resolution and high sensitivity of TMT, will be able to study the atmospheres of planets and protoplanetary disks. These studies will help answer questions related to our place in the Universe: how and when planets are formed and how might life arise on such planets.

4 Summary TMT will provide a major advance in mankind’s ability to probe a very broad range of physics of the Universe. Partial construction funding has just (December 5, 2007) been announced by the Gordon and Betty Moore Foundation, ensuring that the project will be able to transition smoothly into the construction phase. Many more examples and details of programs that will be enabled by TMT are discussed in the Detailed Science Case which can be found at http://www.tmt.org (see Foundation Documents) along with several other documents and references that describe the Observatory in much more detail than is possible here. Acknowledgements The TMT Project gratefully acknowledges the support of the TMT partner institutions. They are the Association of Canadian Universities for Research in Astronomy (ACURA), the California Institute of Technology and the University of California. This work was supported as well by the Gordon and Betty Moore Foundation, the Canada Foundation for Innovation, the Ontario Ministry of Research and Innovation, the National Research Council of Canada, the Natural Sciences and Engineering Research Council of Canada, the British Columbia Knowledge Development Fund, the Association of Universities for Research in Astronomy (AURA) and the U.S. National Science Foundation.

Part V

VLT Synergies with ALMA and JWST

The Atacama Large Millimeter/Submillimeter Array Leonardo Testi

Abstract In this contribution I briefly describe the ALMA project and its current status and timeline. ALMA has been designed and is being built to allow to achieve transformational science in the coming decades. The combination of ALMA, the fully mature VLT/VLTI and later the E-ELT will be a unique asset offered by ESO to the european astronomical community.

1 The ALMA Project The Atacama Large Millimeter/submillimeter Array (ALMA) has been designed and is being built to become the leading ground based observatory at millimeter and submillimeter wavelengths in the foreseeable future. Following the successful experiments in the eighties of millimeter interferometry with a limited number of antennas mostly located at a relatively low altitude, in Europe, North America and Japan were independently initiated plans to built new generations of (sub)millimeter observatories capable of overcome the limitations imposed by the low altitude, limited collective area and frequency coverage, and poor image fidelity. These plans were later merged to become the ALMA project. The highest level science requirements for ALMA are: 1. The ability to detect spectral line emission from CO or C II in a normal galaxy like the Milky Way at a redshift of z = 3, in less than 24 hours of observation. 2. The ability to image the gas kinematics in protostars and protoplanetary disks around young Sun-like stars at a distance of 150 pc (roughly the distance of the star forming clouds in Ophiuchus or Corona Australis), enabling one to study their physical, chemical and magnetic field structures and to detect the tidal gaps created by planets undergoing formation in the disks. 3. The ability to provide precise images at an angular resolution of 0.1. Here the term precise image means representing within the noise level the sky brightness at all points where the brightness is greater than 0.1% of the peak image brightness. This requirement applies to all sources visible to ALMA that transit at an elevation greater than 20 degrees. L. Testi () ESO, Karl Schwarzschild Str. 2, 85748 Garching bei München, Germany e-mail: [email protected] A. Moorwood (ed.), Science with the VLT in the ELT Era, Astrophysics and Space Science Proceedings, doi: 10.1007/978-1-4020-9190-2_48, © Springer Science + Business Media B.V. 2009

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Fig. 1 Panoramic views of the Chajnantor plateau, the ALMA Operations Site, approximately looking to the north (top panel) and to the south (bottom panel). The altitude of the plateau is 5000 m, and it is located close to the city of San Pedro de Atacama in northern Chile. The semi-active Lascar volcano (just over 5900 m peak) is visible at the left side of the top panel. A colour version of this figure is available at dx.doi.org/10.1007/978-1-4020-9190-2_48

The ability to achieve these goals translates in an instrument design with a large collecting area (∼ 7000 m2 ), a wide frequency coverage in the millimeter and submillimeter wavelength range, a large number of antennas for efficient and reliable aperture synthesis, a superb angular resolution, and a site with favorable atmospheric conditions. ALMA will initially be composed of 54 twelve meter and 12 seven meter diameter antennas located on the Chajnantor altiplano at an altitude of 5000 m in northern Chile (see Fig. 1). The plateau offers a relatively constant altitude site for arranging the ALMA antennas in several different configurations. In the most compact configuration the longest available baseline will be of only ∼ 150 m, in the most extended it will be up to ∼ 15 km. To fully exploit the excellent conditions on the site for sub-millimeter observations, ten receiver bands covering all the atmospheric transparency windows from 30 GHz to 1 THz are planned (see Fig. 2). The six highest priority bands (3, 4, 6, 7, 8 and 9) will be available from the start of full science operations at the end of 2012, a decision based on the review of an ongoing R&D program is pending on band 10, and a limited number of band 5 receivers will also be available as part of an EC-FP6 funded program. The remaining frequency bands will be added to the array in a subsequent phase.

2 Transformational Science with ALMA The combination of the excellent site, antenna and receiver quality imply that ALMA will achieve routinely sub-mJy sensitivity in the continuum in all fully available bands. The large bandwidth (8 GHz) and flexible correlator will allow a wide range of spectral line applications.

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Fig. 2 Bottom panel: atmospheric transparency at the Chajnantor site in good (but not exceptional) conditions. Top panel: frequency coverage of the ALMA frequency bands. A colour version of this figure is available at dx.doi.org/10.1007/978-1-4020-9190-2_48

With these characteristics ALMA will be one of the most versatile observatory in the word with applications ranging from observations of the Sun and Solar System objects, of stars and the Interstellar Medium in our Galaxy and beyond, to a probe of the formation and evolution of galaxies. The ALMA Design Reference Science Plan [2] is a collection of almost 150 science programs that could be executed with ALMA during the first few years of full operations and it is used as a reference to ensure that the observatory will meet its scientific objectives. The scientific programs of the worldwide astronomical community were recently reviewed and discussed at an international conference “Science with the Atacama Large Millimeter Array: A New Era for Astrophysics” [4, p. 313]). The conference book contains a comprehensive description of many science programs that will be made possible with ALMA, here I just show two examples closely related with the ALMA top science goals and presented at the conference by Carilli et al. [1] and Wolf [4]. In Fig. 3 the potential of ALMA in detecting molecular and atomic gas at high redshift is shown [1]. The increase in frequency coverage and sensitivity of ALMA will allow to detect and study the molecular gas at high redshift in a similar fashion as it is done today in the local universe. The [CII]λ158 µm line, which is expected to be the main coolant of the warm neutral ISM in galaxies, will be redshifted into the ALMA band and will be a powerful probe of the star formation rate of high redshift galaxies. For galaxies with star formation rates above ∼ 10 M /yr, this line is expected to be detectable in a few hrs with the baseline ALMA system in the

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Fig. 3 Left panel: frequency coverage and spectral sensitivity of ALMA compared with the spectral energy distribution and spectrum of the ultraluminous galaxy Arp220 redshifted to z = 5 and with the most sensitive current millimeter array (PdBI) and the next generation of centimeter wave telescopes. Right panel: simulated ALMA spectrum of the molecular gas in J1148 + 5251 at z = 6.42 for an integration time of 24 hrs (adapted from [1]). A colour version of this figure is available at dx.doi.org/10.1007/978-1-4020-9190-2_48

redshift ranges z ∼ 1 to 8 and beyond z ∼ 10. When the full complement of Band 5 receivers will be installed, the line will be easily detectable in star forming galaxies also in the redshift range 8–10. The molecular lines emission of the cold ISM are expected to be fainter than the [CII] line, nevertheless, in a reasonable integration time, it will also be possible to study the molecular content of high redshift galaxies. As an illustration, Fig. 3 shows a simulated 24 hrs integration time ALMA spectrum for a rich molecular environment at redshift ∼ 6.4 [1]. ALMA simulated observations of a planet forming protoplanetary disk around nearby solar mass stars are shown in Fig. 4 [4]. The ALMA simulated observation at the highest frequency and with the longest available baselines show the effect of

Fig. 4 Simulated ALMA images of forming planets in circumstellar disks. The simulations are for a 1 MJ protoplanet at 5 AU from a 0.5 M star at 50 pc (left) and 100 pc (right) from the Sun. The observing frequency is 900 GHz using the most extended ALMA configuration (adapted from [4]). A colour version of this figure is available at dx.doi.org/10.1007/978-1-4020-9190-2_48

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the protoplanet on the circumstellar disk structure and the material accreting onto the protoplanet. The simulations are for face-on disks around the nearest young stellar objects (50 and 100 pc for the two figures). The simulations assume long Earth rotation synthesis observations (∼ 8 hrs) and excellent weather conditions and calibration performances.

3 ALMA in the Context Full science operations with ALMA are foreseen for 2012, the initial years of ALMA full operations will thus coincide with the VLT 2nd and 3rd generation instruments and with the second generation VLTI instruments. In Fig. 5, adapted from Kurz et al. [3], we show the ALMA parameter space, in terms of frequency coverage and spatial resolution, as compared with other facilities in the coming decade. In terms of spatial resolution, ALMA is an excellent match to the complement of optical and infrared capabilities offered by the Paranal observatory. In fact, this is one of the main design goals of ALMA, the third of the top level science requirements in Sect. 1 was designed to ensure that ALMA will deliver at millimeter wavelengths data of comparable quality as what will be achieved with large optical telescopes equipped with Adaptive Optics systems as well as HST and NGST from space. The wavelength coverage of ALMA is an excellent complement to the optical and infrared for a wide number of science cases. Just to give a couple of examples, in the extragalactic and high-z context, ALMA will allow to probe the interstellar

Fig. 5 Frequency-Angular resolution parameter space of ALMA as compared with some current and future leading observatories [3]. A colour version of this figure is available at dx.doi.org/10.1007/978-1-4020-9190-2_48

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medium of galaxies while optical/infrared studies are mostly sensitive to the stellar component; in the case of circumstellar disks, ALMA will be sensitive to the bulk of the material throughout the disk, while optical and near infrared studies are a key probe of the disk atmosphere and the interaction between the inner disk and the star. A growing number of studies in various fields of astrophysics are showing the importance of combining observations at different wavelengths. The combination of ALMA and the VLT/VLTI, which will be later joined by the E-ELT, will be an exceptional asset for European astronomers.

References 1. C.L. Carilli, F. Walter, R. Wang, A. Wootten, K. Menten, F. Bertoldi, E. Schinnerer, P. Cox, A. Beelen, A. Omont, Studying the first galaxies with ALMA, Astrophys. Space Sci. 313, 307– 311 (January 2008) 2. M. Hogerheijde, The ALMA design reference science plan, Messenger 123, 20 (March 2006) 3. R. Kurz, S. Guilloteau, P. Shaver, The Atacama Large Millimetre Array, Messenger 107, 7–12 (March 2002) 4. S. Wolf, Detecting protoplanets with ALMA. Astrophys. Space Sci. 313, 109–112 (January 2008)

Observational Cosmology with the ELT and JWST Massimo Stiavelli

1 Dark Areas in Observational Cosmology The three hottest themes in modern observational cosmology are all qualified by the word “dark”: Dark Matter, Dark Energy and Dark Ages. Dark Matter and Dark Energy are studied by a variety of techniques both from the ground and from space and we expect that the James Webb Space Telescope (JWST, [1]) will contribute to these areas (e.g. [2]) but will not be dominant. In contrast, the study of the Cosmic Dark Ages is one of the four main themes of JWST and I expect its contributions to this field to be major. The Cosmic Dark Ages [3] are the epoch of cosmic history bracketed by the recombination of Hydrogen at redshift z ∼ 1300 and its reionization at z ∼ 6–7. The three most important milestones during the Dark Ages are the formation of the first (Population III) stars [4, 5], the formation of the first galaxies [6], and the reionization of Hydrogen [7]. In the following section I will review our present observational knowledge in these areas.

2 Observing the Dark Ages with Present Instrumentation The increasingly high optical depth shortwards of Lyman α in the spectra of SDSS QSOs at redshift z ≥ 6 [7] is a strong indication that reionization is completed at z ≈ 6. This is also compatible with the 3 year WMAP Compton depth measurement [8, 9]. If reionization is completed at z ≈ 6 and if it is a fast process occurring over a z 1 we should be able to identify the galaxies responsible for it [10]. Indeed, this was one of the primary motivations for the Hubble Ultra Deep Field (UDF, [11]). Focusing on galaxies found with the Lyman Break technique (LBT), the Great Observatory Origins Deep Survey (GOODS, [12]) and the UDF have provided us with sample of > 500 i-dropout galaxies likely to be at z ∼ 6 [13, 14]. Estimating their ionizing flux involves an extrapolation from the observed, non-ionizing, UV continuum to the Lyman continuum. Similarly, uncertainties on the ionizing photons M. Stiavelli () Space Telescope Science Institute, 3700 San Martin Dr., Baltimore, MD 21218, USA e-mail: [email protected] A. Moorwood (ed.), Science with the VLT in the ELT Era, Astrophysics and Space Science Proceedings, doi: 10.1007/978-1-4020-9190-2_49, © Springer Science + Business Media B.V. 2009

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escape fraction and on the gas clumping (driving recombinations) render uncertain the required amount of radiation needed for reionization. The observed flux is tantalizing close to what is needed but the uncertainties have led different groups to different conclusions depending on whether one assumes the presence of large numbers of faint dwarf galaxies [15], or a top heavy stellar mass function and low metallicity [16], or galaxy parameters similar to those found at lower redshift [17]. Whether or not the galaxies at z ∼ 6 are sufficient to reionize the Universe has implications on the evolution of the luminosity function (LF) of galaxies. Indeed, if the LF evolves very fast at z > 6, it is unlikely that galaxies at z > 6 will contribute much to reionization and those at z ∼ 6 have to do it by some combination of a steep, dwarf-rich LF and low metallicity. On the other hand, if there is rapid evolution of the LF at z > 6, galaxies at z > 6 will contribute to the ionizing background extending the duration of reionization and pushing the time of the formation of the first galaxies to higher redshift. In order to search for galaxies at z ≥ 7 we started a UDF followup study [18]. The preliminary result is that there is a deficit of high-z objects and this has been first reported by [19]. However, it is unclear whether the evidence of a rapid evolution of the LF is statistically significant once the systematic errors of the modeling and the impact of cosmic variance are properly included [20]. Thus, the jury is still out on whether the LF is really evolving rapidly. It is well possible that surveys carried out with the Wide Field Camera 3 installed in the Hubble Space Telescope during the upcoming servicing mission will clarify this issue. Ongoing surveys exploiting gravitational lensing might also provide us with very high redshift objects [21] but deriving a LF from them is going to be hard given the small effective volume and the related high cosmic variance. An alternative to the LBT is the search for Lyman α emitting sources by use of the narrow-band excess technique which has delivered the highest redshift spectroscopically confirmed galaxy known to date [22]. However, even with well identified candidates we are left with two open questions: how far can one trust dropout selections (or single line redshifts) in an essentially unprobed redshift region and how can one decrease the uncertainty in the ionizing photons output by improved physical modeling of these galaxies. Both questions require spectroscopy to be addressed.

3 The Need for Spectroscopy The majority of the galaxies we are interested in at z ≥ 6 are beyond the reach of spectroscopic study from the ground. This is illustrated in Fig. 1. In the left panel we show the z850 magnitude vs S/N plot for the UDF. Points in red are the i-dropout galaxies. As an illustration of the difficulty to obtain high S/N spectra at this faint levels, I report in Fig. 1 the integration time (in hours) with FORS2 on the Very Large Telescope (VLT) required to achieve a S/N = 3 per resolution element at 900 nm, which is a representative wavelength for the continuum below Lyman α at z < 6.4. Here we focus on the continuum because about 50% of Lyman break galaxies are expected not to have Lyman α in emission [23] and in any case studying

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Fig. 1 Left: S/N vs z850 magnitude for the UDF. Points in red represent i-dropout galaxies. Right: S/N vs magnitude for the NICP12 field of the UDF followup program. For both we give the integration times for spectroscopic followup (see text). A colour version of this figure is available at dx.doi.org/10.1007/978-1-4020-9190-2_49

the fainter lines needed for deriving the physical properties of these objects would require at least this level of S/N. Figure 1 shows that about 90% of the galaxies in the UDF and the vast majority of i-dropout galaxies would require 100 hours of integration or more with the VLT. For a few objects, single emission line redshifts can be obtained at fainter magnitudes but between the faintest objects at S/N = 5 and z850 = 29.5 and those for which one can typically obtain a useful spectrum we have a gap of 4.5 magnitudes and even for shallower (but wider) surveys like GOODS there is a gap of at least 2 magnitudes. A similar situation is present in the near-IR and is illustrated in the right panel of Fig. 1 where I plot the S/N vs J110 magnitude for the NICP12 field in the UDF area [18]. Here the spectroscopic benchmark is an integration with SINFONI on the VLT with S/N = 1 in the continuum at 1200 nm. I have adopted a lower S/N because when studying the continuum one could presumably rebin several resolution elements after having masked out the brightest OH lines. Using 100 hours as the practical limit, Fig. 1 shows that objects fainter than J110 = 25.2 are too faint for spectroscopic study and this implies a gap of about 3 magnitudes and includes 76% per cent of the galaxies in the NICP12 field. The installation of WFC3 on HST will further increase this gap. JWST will be able to obtain spectra of many galaxies using the multi-object spectrometer NIRSpec. However, the field-of-view (FOV) of NIRSpec is only about 10 square arcmin. This is well matched to the size of the UDF, but NIRSpec would require at least 30 separate integrations to cover GOODS. Morover, JWST will be able to obtain images with NIRCam deeper by a couple of magnitudes than the UDF thus preserving a gap between imaging and followup spectroscopy.

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4 Rarity of High-z Galaxies The first galaxies, almost by definition, would be rare. Clearly we do not yet know how rare, how faint, and at what redshift they will occur. Using a simple model, we have derived expected counts for high-z dropouts objects seen by NIRCam [20]. Based on that prediction and considering the FOV size of about 9.7 square arcmin for NIRCam one would expect about 10 objects per deep NIRCam field at z 10, about 1 per field at z 12 and about 0.1 per field at z 15. Moving from first galaxies to first stars also forces us to contemplate rare objects. Indeed it is likely that we will be able to detect Population III stars only when they produce ultra-bright pair-instability supernovae and such supernovae are probably extremely rare, e.g., 4 per square degree per year at z = 15 [25]. In order to detect them, one would need degree-scale surveys with sensitivity around AB = 26–27 in the near-IR. Clearly for these objects the limiting factor would not be the JWST sensitivity but its field of view.

5 How Can the ELT Help? JWST is optimized for IR imaging and low-medium resolution spectroscopy over a modest FOV. The need for identifying rare objects or employing higher spectroscopic resolving power require ground based telescopes of the 30+ m class like, e.g., the Extremely Large Telescope (ELT). To help close the gap between spectroscopy and imaging it would be very desireable to have a high-throughput multi-object spectrograph operating in the red and in the near-IR on the ELT with a field of view as large as practical and a resolving power between 3000 and 5000 so as to be able to reject the brightest OH lines and retain a low background between such lines. Alternatively one could work at lower resolution if efficient OH-suppression spectrographs become feasible. Full adaptive optics would not be very important for such an application as galaxies are marginally resolved at the resolution of HST and JWST and, unlike the case of stars, one would not gain a background reduction by having a sharper point spread function. To the contrary, in the detector limited case, perhaps achievable between the lines, one might lose S/N by over-resolving the galaxies. Thus, the two critical parameters for such a spectrograph would be throughput and multiplexing. Kinematical study of faint high-z galaxies might well require higher resolving power (up to 10 000) than afforded by JWST due to their low mass [24] and this could be addressed by the ELT. Similarly, absorption line studies of QSOs at z > 7 will need to be done from the ground at R ∼ 30 000. Helping JWST to identify Population III objects would require monitoring a large area (at least one square degree) down to AB = 26 or 27. This may or may not be possible on the ground depending on the feasibility and efficiency of OHsuppression imagers and possibly moderate field MCAO imagers (as supernovae are point-like and one would benefit from a sharper point spread function). The belief

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that pair-instability supernovae evolve slowly (e.g. 200 days to decrease by 2 magnitudes from the peak), combined with cosmological time dilation makes a yearly monitoring program acceptable and would render imaging programs requiring many nights of integration feasible.

References 1. 2. 3. 4. 5. 6. 7. 8. 9. 10. 11. 12. 13. 14. 15. 16. 17. 18. 19. 20. 21. 22. 23. 24. 25.

J.P. Gardner, J.C. Mather, M. Clampin et al., Space Sci. Rev. 123, 485 (2006) A.G. Riess, M. Livio, Astrophys. J. 648, 884 (2006) R. Barkana, A. Loeb, Phys. Rep. 349, 125 (2001) V. Bromm, R.B. Larson, Annu. Rev. Astron. Astrophys. 42, 79 (2004) M. Trenti, M. Stiavelli, Astrophys. J. 667, 38 (2007) J.H. Wise, T. Abel, submitted to Astrophys. J. (2007), see also arXiv:0710.3160 X. Fan, M.A. Strauss, R.H. Becker et al., Astron. J. 132, 117 (2006) D.N. Spergel, R. Bean, O. Doré, Astrophys. J. Suppl. 170, 377 (2007) M. Shull, A. Venkatesan, submitted to Astrophys. J. (2007), see also astro-ph/0702323 M. Stiavelli, M.S. Fall, N. Panagia, Astrophys. J. 600, 508 (2004) S.V.W. Beckwith, M. Stiavelli, A.M. Koekemoer et al., Astron. J. 132, 1729 (2006) M. Giavalisco, H.C. Ferguson, A.M. Koekemoer et al., Astrophys. J. 600, L93 (2004) R.J. Bouwens, G.D. Illingworth, J.P. Blakeslee, M. Franx, Astrophys. J. 653, 53 (2006) R.J. Bouwens, G.D. Illingworth, M. Franx, H. Ford, Astrophys. J. 670, 928 (2007) H. Yan, R.A. Windhorst, Astrophys. J. 600, L1 (2004) M. Stiavelli, M.S. Fall, N. Panagia, Astrophys. J. 610, L1 (2004) A.J. Bunker, E.R. Stanway, R.S. Elllis, R.G. McMahon, Mon. Not. R. Astron. Soc. 355, 374 (2004) P.A. Oesch, M. Stiavelli, C.M. Carollo et al., Astrophys. J., in press (2007), see also arXiv:0706.2653 R.J. Bouwens, G.D. Illingworth, Nature 443, 189 (2006) M. Trenti, M. Stiavelli, Astrophys. J., in press (2007), see also arXiv:0712.0398 D.P. Stark, R.S. Ellis, J. Richard et al., Astrophys. J. 663, 10 (2007) M. Iye, K. Ota, B. Kashikawa et al., Nature 443, 186 (2006) C.C. Steidel, K.L. Adelberger, A.E. Shapley et al., Astrophys. J. 532, 170 (2000) M. Stiavelli, in The Next Generation Space Telescope: Science Drivers and Technological Challenges. ESA SP, vol. 429, p. 71 S.M. Weinmann, S. J , Lilly. Astrophys. J. 624, 526 (2005)

Integral Field Spectroscopy of (U)LIRGs. From VLT to JWST L. Colina, S. Arribas, A. Bedregal, A. Monreal-Ibero, M. García-Marín, A. Alonso-Herrero and J. Alfonso

Some first results from the Integral Field Spectroscopy Survey of (U)LIRGs using VLT instruments VIMOS and SINFONI are presented. Detailed studies of the twodimensional ionization structure and kinematics of the stars and different gas phases are within reach with IFS techniques on 8 m class telescopes. The perspectives of extending these studies to high-z galaxies with future IFS instruments (NIRSpec and MIRI) onboard JWST are considered.

1 Integral Field Spectroscopy Survey of Low-z (U)LIRGs In the context of the high-z universe, the study of representative samples of low-z galaxies populations is of particular importance. That is the case of the luminous (LIRG, 1011 ≤ Lir < 1012 L ) and ultraluminous (ULIRG, 1012 ≤ Lir < 1013 L ) infrared galaxies. These objects could be the local counterparts of cosmologically important galaxy populations at z > 1 such as the submillimeter galaxies [3], and the majority of the infrared sources found in recent Spitzer surveys [6]. The IFS (U)LIRG Survey is a large program aimed at studying the internal ionization structure and kinematics of a representative sample of low-z (U)LIRGs, capturing their complex two dimensional structure on scales (kpc), relevant for studies of high-z galaxies with future facilities as the JWST. This program has initially been focused on ULIRGs [4] and is now being extended towards lower luminosities, i.e. the LIRG luminosity range. The IFS (U)LIRG Survey is using integral field spectroscopy facilities both in the south (VIMOS and SINFONI), and north (INTEGRAL and PMAS) hemispheres. The complete survey contains about 80 (U)LIRGs representing the different morphologies (Fig. 1; [1]).

L. Colina () DAMIR, Instituto de Estructura de la Materia (CSIC), Madrid, Spain e-mail: [email protected] A. Moorwood (ed.), Science with the VLT in the ELT Era, Astrophysics and Space Science Proceedings, doi: 10.1007/978-1-4020-9190-2_50, © Springer Science + Business Media B.V. 2009

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Fig. 1 Distribution of the objects of the sample in the luminosity—redshift and luminosity—morphological type planes. Isolated objects (0), pre-coalescence interacting pairs (1), and single post-coalescence systems with evidence of having suffered a merger (2) are indicated. Circles correspond to the VIMOS sample, crosses and full triangles to the INTEGRAL and PMAS samples, respectively. LIRGs (black symbols) and ULIRGs (red symbols) are indicated. A colour version of this figure is available at dx.doi.org/10.1007/978-1-4020-9190-2_50

2 VIMOS Optical IFS. First Results The observations use the integral field unit of VIMOS covering the 5250–7400 Å range with a spectral resolution of about 3400. The field of view covers 27 arcsec square (0.67 per fiber) with a total of about 1600 simultaneous spectra per pointing. A square four pointing dither pattern with a total integration of about 3000 seconds was used per galaxy.

2.1 IRASF 12115–4656: Neutral and Ionized Gas Kinematics The stars and ionized gas distributions present some clear differences (Fig. 2). The stellar distribution shows an elliptical shape with a well defined central nucleus which appears as a relatively weak Hα emitter. The ionized gas distribution is dominated by the presence of compact Hα knots. These knots at distances of 1 to 1.5 kpc are most likely HII regions located in the inner spiral arms of the galaxy. The neutral gas distribution (Fig. 3) is similar to that of the stars with a clear peak at the photometric center of the galaxy. To compare the neutral and ionized gas kinematics, the rotational velocity pattern of the neutral gas has been substracted from that of the ionized gas. The residuals map (Fig. 3) shows the morphology of a bipolar flow with velocities of up to ±100 km s−1 . This behavior can be explained as if the neutral gas was rotating systematically lower than the ionized gas by ∼ 70–90 km s−1 [1]. The neutral and ionized gas central velocity dispersions agree well, as expected if these values trace the mass in these regions. However, while the ionized gas V/σ values are high (> 3), the neutral gas shows significantly lower values (< 1). This result suggest that the

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Fig. 2 VIMOS results for IRASF 12115–4656 showing the stellar continuum, the Hα light distribution together with the ionization map ([NII]/Hα). The velocity field and velocity dispersion map as obtained from the Hα line profile are also shown. A colour version of this figure is available at dx.doi.org/10.1007/978-1-4020-9190-2_50

Fig. 3 Spatial distribution and kinematic properties of the neutral gas in IRASF 12115–4656 as traced by the NaI doublet (top panels). The equivalent width of the Hα line tracing the location of young star-forming regions is given for comparison. The residuals of the velocity field (v) and velocity dispersion (σ ) of the neutral gas after substraction of those measured for the ionized gas are shown. A colour version of this figure is available at dx.doi.org/10.1007/978-1-4020-9190-2_50

dynamical status of the two gas phases is different. One (ionized gas) traces the presence of a rotating star-forming disk while the other (neutral gas) could be tracing the overall dynamics of the stellar distribution [1].

3 SINFONI Near-IR IFS. First Results A subsample of the IFS (U)LIRG survey sample covering different morphologies and luminosities has been selected for near-infrared IFS studies. The combination of near-IR and optical IFS provides a full 2D multi-wavelength investigation of the

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Fig. 4 SINFONI results for NGC 5135. Top row shows the fluxes of four emission lines (logarithmic scale) each one mapping different galaxy components and/or physical phenomena in the central 1 kpc radius. The second row shows the preliminary velocity fields for the near-IR lines with respect to that of the K band nucleus. A colour version of this figure is available at dx.doi.org/10.1007/978-1-4020-9190-2_50

kinematics and physical properties of the stellar populations and gas phases in these galaxies.

3.1 NGC 5135: Young Stars and AGN The high-excitation coronal [SiVI] emission is centered in the nucleus tracing the presence of the Seyfert 2 nucleus (Fig. 4). The Brδ map traces the overall ionization field and therefore the presence of young massive star clusters. The [FeII] spatial distribution is in good agreement with the radio emission [7], and consistent with previous claims that the [FeII] emission traces supernovae. Warm molecular gas (H2 line map) shows a spatial structure similar to that of [FeII]. The relative surface brightness of H2 in the nucleus and extranuclear star-forming regions suggest different excitation mechanisms at work (hard UV radiation field from AGN vs. shocks from SNe). The ionized and molecular gas seem to share a similar velocity field given by the orientation of the spiral structure of the galaxy (Fig. 4). The [FeII] emission presents a localized velocity structure at its peak, likely associated with outflows from stellar winds and supernovae. The velocity structure of the [SiVI] emission shows departures from rotation of more than 100 km s−1 . Full account of the ionization and kinematical structure of NGC 5135 will be published elsewhere (Bedregal et al., in prep.).

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4 JWST IFS of High-z (≥ 6) Starbursts Some of the future instruments (NIRSpec and MIRI) onboard James Webb Space Telescope (JWST) will incorporate Integral Field Spectrographs ([5] for a review). NIRSpec will provide low (R ∼ 100) and intermediate (R ∼ 1000 to 2700) spectral resolution over the entire 0.6 to 5 µm range, with a spatial pixelation of 0.1 and a 3 arcsec square field of view. On the other hand, MIRI IFS will provide medium resolution spectroscopy (R ∼ 3000) over the full 5 to 28.3 µm range [9]. The MIRI IFU is subdivided into four image slicers covering concentric fields of view on the sky of about 3.5 arcsec square (4.9–7.7 µm) to 7 arcsec square (18–28 µm). The predicted NIRSpec spectral line sensitives (10σ detection on 10 000 seconds exposures) for unresolved, point sources are in the 0.25 to 1 × 10−20 W m−2 range (Jakobsen, private communication). The corresponding MIRI sensitives range from 1 to 6 × 10−20 W m−2 for the short (∼ 6.4 µm), and long (∼ 22.5 µm) wavelengths, respectively [8]. Low-z ULIRGs have observed Hα fluxes of 0.2 to 3 × 10−16 W m−2 integrated over areas of two to three kpc in radius [4]. At the redshift of 7 these galaxies will have apparent fluxes of 0.6 to 20 × 10−21 W m−2 in regions with sub-arcsec angular sizes, similar to those expected for intermediate-z Lyman break galaxies (LBGs; [2]). So, the two-dimensional ionization and kinematic structure of luminous high-z starbursts could be traced in detail with JWST, if they present physical properties similar to ULIRGs and LBGs.

References 1. 2. 3. 4. 5. 6. 7. 8. 9.

S. Arribas et al., Astron. Astrophys. (2008), in press D. Erb et al., Astrophys. J. 646, 107 (2006) D.T. Frayer et al., Astron. J. 127, 728 (2004) M. García-Marín, 2D structure and kinematics of a representative sample of low-z ULIRGs, PhD Thesis, Autonóma University, Madrid (2007) J.P. Gardner et al., Space Sci. Rev. 123, 485 (2006) P. Pérez-Gonzalez et al., Astrophys. J. 630, 82 (2005) J. Ulvestad, A. Wilson, Astrophys. J. 343, 659 (1989) B.M. Swinyard et al., Sensitivity estimates for MIRI on the JWST, in Optical, Infrared and Millimeter Space Telescopes, ed. by J.C. Mather. Proc. SPIE, vol. 5487 (2004), pp. 785–793 G.S. Wright et al., The JWST MIRI instrument concept, in Optical, Infrared and Millimeter Space Telescopes, ed. by J.C. Mather. Proc. SPIE, vol. 5487 (2004), pp. 653–663

Part VI

Second Generation VLT and VLTI Instrument Programme

VLT and VLTI Second Generation Instrument Overview and Resources Alan Moorwood

Abstract ESO’s program of second generation VLT and VLTI instrumentation comprises (i) 4 approved instruments already in development (X-Shooter, KMOS, MUSE and SPHERE), (ii) up to 3 second generation VLTI instruments whose Phase A studies were recently completed and (iii) tbd additional second generation instruments whose selection prompted this Workshop. The resources provisionally included for additional VLT instruments in ESO’s future planning are comparable to those allocated to those already approved.

1 Operational VLT and VLTI Instruments Figure 1 shows the instruments currently in operation on the VLT (ISAAC, FORS1 and 2, CRIRES, UVES, FLAMES + GIRAFFE, VIMOS, VISIR, NACO, SINFONI and HAWK-I) and VLTI (AMBER and MIDI). Together they provide extremely wide coverage of the observational ‘phase space’ defined by wavelength (UV— 25 µm), spectral resolution (up to 100.000 from UV to IR), spatial resolution (down to the diffraction limit in the IR plus higher with VLTI), field (up to 14 × 14 ) and multiplex capability (visible spectra of up to 1000 objects in one shot). Detailed information about their modes and performance characteristics can be found on ESO’s Web pages. In addition, one of the UT3 Nasmyth foci is reserved for visitor instruments and has been used recently by the ULTRACAM and DAZLE instruments. One general aspect maybe for discussion at this Workshop is that much of the success of the VLT so far is attributed by many to this diversity of observational capability provided by up to 14 instruments and individual use of the four unit telescopes. To date, only two UTs have been combined interferometrically, no instrument has been built for the incoherent combined focus and the only remnant of an early plan for multiple instrument copies are the two FORS’s which are no longer identical and have never been used in tandem. There are certainly also those who also believe that the limits of field and multiplex have not yet been reached!

A. Moorwood () ESO, Karl-Schwarzschild-Str. 2, 85748 Garching, Germany e-mail: [email protected] A. Moorwood (ed.), Science with the VLT in the ELT Era, Astrophysics and Space Science Proceedings, doi: 10.1007/978-1-4020-9190-2_51, © Springer Science + Business Media B.V. 2009

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Fig. 1 VLT/I instrumentation operational in 2007. A colour version of this figure is available at dx.doi.org/10.1007/978-1-4020-9190-2_51

2 2nd Generation VLT Instruments The approved 2nd generation instruments already in development are X-Shooter (UV-IR single object spectrograph), KMOS (I-K band MOS), MUSE (wide-field visible IFU spectrograph) and SPHERE (Extreme AO planet finder) which were selected by a process kicked—off by the first ESO Workshop on 2nd generation VLT instrumentation held in 2001 [2] and which should all have been installed by 2012. All of these instruments will be presented at this Workshop and further details can be found on ESO’s Instrumentation Web pages. In addition, by then UT4 should have been upgraded to become an Adaptive Optics Facility (AOF) with a deformable M2 and 4 laser guide stars which will provide wide—field ground layer AO correction for HAWK-I and ground-layer wide field plus higher Strehl ratio narrow field corrections for MUSE. To fully utilize it the Cassegrain focus of UT4 should also be equipped at some stage e.g. when SINFONI is de-commissioned, with an instrument capable of exploiting the AOF. Figures 2, 3 and 4 illustrate the expected evolution of the VLT instrumentation out to around 2020 assuming that no new instruments are approved. Beyond 2012, capabilities will begin to decrease due to technical obsolescence of 1st generation instruments. Assuming the nominal lifetimes of 10 years specified for VLT instruments all capability will have been lost by around 2022. Assuming individually estimated lifetimes for each instrument (up to around 20 yrs for FORS, UVES, X-Shooter) extends this but models still predict fewer than 4 instruments still operating in 2025. Because of the tendency to select more specialized

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Fig. 2 VLT/I instrumentation in 2008. A colour version of this figure is available at dx.doi.org/10.1007/978-1-4020-9190-2_51

2nd generation instruments so far the first scientific capabilities to be lost are likely to be those provided by the more work-horse instruments including thermal infrared imaging and spectroscopy, ‘wide-field’ imaging, high resolution spectroscopy and visible multi-object spectroscopy.

3 2nd Generation VLTI Instruments The selection process for second generation VLTI instruments started with an ESO Workshop in 2005 [1]) which generated about 9 concepts of which 3 (Gravity, Matisse and VSI) have finally been the subject of Phase A studies carried out by Consortia. Briefly, as these instruments will also be presented at this Workshop, Gravity is designed primarily for faint, high precision astrometry in the K band e.g. to exploit the region around the black hole at the center of our galaxy as a laboratory for the study of strong gravity; Matisse represents a natural evolution of MIDI but will cover the L, M as well as the N band and combine up to 4 telescopes for ‘imaging’ and spectroscopy; VSI would follow on naturally from AMBER for ‘imaging’ and spectroscopy in the J, H, K bands but combining up to six telescopes. A recommendation on how to proceed with these instruments is being sought from ESO’s Scientific and Technical Committee in October 2007. The contractual basis for selected instruments will be that Consortia will be expected to contribute most of the budget and manpower needed to develop the instruments, under contract to ESO,

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Fig. 3 VLT/I instrumentation 2012. A colour version of this figure is available at dx.doi.org/10.1007/978-1-4020-9190-2_51

in exchange for guaranteed observing time. ESO will contribute specific items e.g. detector systems plus any common VLTI infrastructure including, potentially, a 4–6 beam fringe tracker.

4 Additional 2nd Generation Instruments In order to ensure that the VLT remains at the scientific cutting edge for the foreseeable future both money and manpower for new instruments have been earmarked in ESO’s long term perspectives and will start to become available nominally in 2010. The amounts are comparable to those invested by ESO in the four 2nd generation VLT instruments already approved i.e., roughly 20 M€ for hardware and about 50 FTE’s of dedicated manpower. In addition, for the four approved instruments, the community has contributed an additional 10 M€ and several hundred FTE’s in return for guaranteed observing time. It is hoped that something comparable to this will prove possible again for the additional instruments. Of course this cannot be guaranteed, particularly as the ELT itself will also be competing for instrument resources. This Workshop has provided the first semi formal opportunity to present ideas for new instruments and to discuss scientific priorities taking into account the possible synergies with other major facilities including ALMA and JWST as well as, hopefully, a European ELT in around a decade from now. The results will be discussed with ESO’s Scientific and Technical Committee (STC) which will make

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Fig. 4 VLT/I instrumentation 2020. A colour version of this figure is available at dx.doi.org/10.1007/978-1-4020-9190-2_51

recommendations to Council to be converted into an updated instrumentation plan by ESO followed by launch of the necessary calls for Phase A studies for potential new instruments to be selected over the next few years and installed in roughly the 2015–2020 time frame.

References 1. A. Richichi, F. Delplancke, F. Paresce, A. Chelli (eds.), The Power of Optical/IR Interferometry: Recent Scientific Results and 2nd Generation Instrumentation. ESO Astrophysics Symposia (Springer, Berlin, 2008) 2. J. Bergeron, G. Monnet (eds.), Scientific Drivers for ESO Future VLT/VLTI Instrumentation, ESO Astrophysics Symposia (Springer, Berlin, 2002)

HAWK-I and Infrared Imaging on the VLT M. Casali, N. Ageorge, C. Alves de Oliveira, P. Biereichel, M. Casali, B. Delabre, R. Dorn, R. Esteves, G. Finger, D. Gojak, G. Huster, Y. Jung, F. Koch, M. Kiekebusch, M. Kissler-Patig, M. Le Louarn, J.-L. Lizon, L. Mehrgan, A. Moorwood, J. Pirard, E. Pozna, A. Silber, B. Sokar and J. Stegmeier

1 HAWK-I Specifications HAWK-I, the High Accuity Wide-field K-band Imager, will be a near-optimum camera for the VLT. The 7.5 arcminute square field is practically the largest IR field possible at Nasmyth while keeping reasonably uniform sensitivity in all bands. The mosaic of four 2k × 2k detectors which fills this field has a pixel scale of 0.1 arcsec/pixel, which is sufficiently small to adequately sample the best seeing at Paranal, even with future ground layer adaptive optics correction. An all-reflective optical design gives very high throughput. The end result is an imager with the best possible performance, limited predominantly by the telescope design and atmospheric seeing conditions. The HAWK-I focal plane consists of four Raytheon HgCdTe 2k × 2k detectors. Their high quantum efficiency combined with an all-reflective optical design give a total system throughput of 50 to 60% depending on the waveband. It is envisioned the instrument will be the first to be used with ground-layer correction at the Adaptive Optics Facility (a multi-laser facility using a deformable secondary mirror) which is currently under development at ESO. The HAWK-I pixel scale of 0.1 arcseconds was chosen to adequately sample the best images expected. The default filter complement consists of a broadband set of YJHK filters and a series of conventional (CH4 , H2 S(1), Brγ ) and cosmological (1061, 1187, 2090 nm) narrowband filters.

2 Commissioning and Performance Commissioning on UT4 began at the start of August 2007, with a second run at the end of September, followed by the start of science verification (SV) which continues to February 2008. The SV phase is designed to test and tune the instrument M. Casali () ESO, Karl-Schwarzschild-Str. 2, 85748 Garching, Germany e-mail: [email protected] A. Moorwood (ed.), Science with the VLT in the ELT Era, Astrophysics and Space Science Proceedings, doi: 10.1007/978-1-4020-9190-2_52, © Springer Science + Business Media B.V. 2009

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Fig. 1 HAWK-I mounted on the Nasmyth port of UT4

with real science projects selected following a review of proposals submitted by the community. The data from these projects will be made public.

2.1 Image Quality and General Performance Image quality was a primary concern during instrument integration and commissioning. Although the optical design has very good theoretical image quality, a number of off-axis and aspheric surfaces are present in the system, and proper optical alignment is important. HAWK-I was found to perform to specification, with excellent image quality across the field measured at less than 2 pixels at 80% encircled energy. The sharp images attainable are illustrated in Fig. 2 during a period of very good seeing on Parnal. The wide field and excellent optics should also allow very precise astrometry, and an analysis of the field shows relative astrometry accurate to 3 mas. In all other significant aspects HAWK-I meets the scientific requirements. • The instrument background is as expected in the different wavelength bands. • Flexure is always an important aspect for Nasmyth instruments, and after some initial problems and subsequent strengthening of the supports the flexure is now less than 2 pixels in a complete rotation, well within specification. HAWK-I will have substantial synergy with two other ESO projects. KMOS is a multi-IFU cryogenic spectrometer with a patrol field almost identical to the HAWK-

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Fig. 2 15 × 15 arcminute field in the LMC, taken during very good seeing conditions. FWHM is 0.35 arcsec in Y band. The inset image shows a magnified part of the field, and image scale

I field of view. Images taken with HAWK-I can therefore be used as inputs to the KMOS observing preparation tool. HAWK-I will also be the ideal follow-up imager for VISTA, the 4-m IR survey telescope due to be commissioned in 2008; the larger VISTA surveys are quite shallow so HAWK-I will be able to make much deeper and sharper images of any interesting objects which VISTA finds.

3 Widefield IR Imaging at ESO 3.1 Near-Term Capabilities By the end of 2008 the ESO community will have access to the best IR imaging facilities in the world. VISTA will be world class through its 16 2k × 2k detectors which make up the largest IR focal plane ever constructed for astronomy, while its 4-m aperture telescope is the largest ever dedicated to imaging surveys. On smaller scales, HAWK-I will be able to go much deeper due to the better image quality and larger aperture of the VLT. Together they will be a powerful combination in studies of both the high redshift and nearby universe.

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3.2 Possibilities for Enhancements There is great interest in IR survey imaging in the community, so that despite the new facilities available, there is always pressure to look at still further improvements. As one might anticipate, this is not a simple task. The HAWK-I throughput is already very high at 50–60%. VISTA will probably be similar, so that there is little gain to be achieved. Could the field of view be increased? HAWK-I already fills the unvignetted field at Nasmyth, so the only way a larger field could be achieved conventionally, would be to go to prime focus of the VLT. With a 1-m clear aperture corrector, and assuming a 37% filling factor as in the VISTA focal plane, one could achieve a single-shot field of 2.9 square degrees. In terms of survey speed this would be 19 time faster than VISTA. It would, however, require some 75 2k × 2k detectors and major engineering to the telescope topend, including replacing the active optics wavefront sensing. Such a project would be very major in scale, costing in the 20–40 Meuro range. Non-conventional approaches (e.g. forward-cassegrain as discussed elsewhere in these proceedings) may also be possible but are unlikely to match the prime focus field of view. If a big increase in widefield imaging on the VLT is probably not feasible, what general improvements could be achieved? One area where large gains are possible is in image quality. Although ground-layer AO will deliver reasonable improvements over a large field, real gains require the use of multi-conjugate adaptive optics (MCAO). The high strehl ratio possible over 1 arcminute fields with MCAO would result in large gains in imaging sensitivity. In background-limited operation, the integration time to reach a given depth scales as the psf FWHM2 , so great gains can be achieved by going from a seeing-limited to a high Strehl ratio regime.

4 Conclusions HAWK-I is an optimum IR imager for the VLT. It combines a fine pixel scale, excellent optical quality and throughput, and the largest possible unvignetted field at Nasmyth. Improvements in IR imaging on the VLT which will enable it to surpass VISTA in surveys cannot be easily achieved, requiring very large focal planes and a major re-engineering of the telescope structure.

X-Shooter: A Medium-resolution, Wide-Band Spectrograph for the VLT L. Kaper, S. D’Odorico, F. Hammer, R. Pallavicini, P. Kjaergaard Rasmussen, H. Dekker, P. Francois, P. Goldoni, I. Guinouard, P.J. Groot, J. Hjorth, M. Horrobin, R. Navarro, F. Royer, P. Santin, J. Vernet and F. Zerbi

X-shooter is the first second-generation instrument for the ESO Very Large Telescope, and will be installed in 2008. It is intended to become the most powerful optical & near-infrared medium-resolution spectrograph in the world, with a unique spectral coverage from 300 to 2500 nm in one shot. The X-shooter consortium members are from Denmark, France, Italy, The Netherlands and ESO.

1 X-Shooter: A Very Efficient Spectrograph The concept of X-shooter has been defined with one single main goal in mind: The highest possible throughput for a point source at a resolution which is just sky limited in about an hour of exposure over the broadest possible wavelength range, without compromising throughput at the atmospheric UV cutoff [2]. The moderate size of X-shooter is, as opposed to most existing or planned VLT instruments, compatible with implementation at the Cassegrain focus. The instrument design is based on multiple dichroics to split the light between the three spectrograph arms (Fig. 1). The central backbone supports three prism cross-dispersed echelle spectrographs (in double pass, optimized for the UV, visible and near-infrared wavelength range and based on the so-called 4C design [1]). The backbone contains the calibration and acquisition units, an IFU that can be inserted in the lightpath, the two dichroics that split the light into the three arms and relay optics to feed the entrance slits of the three spectrographs. The spectral performance and efficiency (better than 95% reflectance, resp. transmission) of the two dichroics (cross-over wavelength 550 and 1000 nm) turn out to be exceptionally good, especially when considering the enormous wavelength range covered by X-shooter. The standard slit measures 12 × 1 ; higher spectral resolution is obtained when using the 12 × 0.6 slit (Table 1). A wide slit (12 × 5 ) is available for flux calibration. A dedicated program is being executed with VLT/SINFONI to extend the L. Kaper () Astronomical Institute, University of Amsterdam, Amsterdam, The Netherlands e-mail: [email protected] A. Moorwood (ed.), Science with the VLT in the ELT Era, Astrophysics and Space Science Proceedings, doi: 10.1007/978-1-4020-9190-2_53, © Springer Science + Business Media B.V. 2009

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Fig. 1 Left: The X-shooter spectrograph depicted in the Cassegrain focus below the M1 mirror cell of the VLT. To reduce flexure, the center of gravity is located as close as possible to the telescope, and the stiff elements are light-weighted. The UV and VIS spectrograph arms are mounted to the side; the NIR arm is at the bottom of the instrument. Right: Schematic view of X-shooter [3]. A colour version of this figure is available at dx.doi.org/10.1007/978-1-4020-9190-2_53 Table 1 Predicted average efficiencies (at the blaze of the echelle orders and outside the dichroic cross-over range) per spectrograph arm. The technical specification requires an efficiency > 25%. Also the predicted spectral resolving power (R = λ/λ) is listed and compared to the technical specification Rspec for a 0.6 slit (excluding the detector LSF) Spectrograph

Spectral range (nm)

Average blaze efficiency

Rpred (0.6 slit)

Rspec (0.6 slit)

UVB

307–529

41.6%

8169

7600

VIS

558–966

35.6%

12 335

11 500

NIR

1040–2370

27.8%

7329

7000

calibration of 16 optical spectro-photometric standards to the near-infrared in order to perform flux calibration of X-shooter spectra aiming at an accuracy of better than 5%. The IFU has an entrance window of 4 × 1.8 and delivers three slices filling a 12 × 0.6 exit slit. The Acquisition and Guiding unit has a 1.5 × 1.5 field and includes a comprehensive filter set; atmospheric dispersion compensation is performed in the UV and VIS arms (though not when using the IFU). Given its location in the Cassegrain focus, X-shooter has a tight weight (less than 2.5 tonnes) and flexure budget. Measures have been taken to reduce the effects of flexure, e.g. the optical bench of the NIR arm is extreme light-weighted (machined from a block of 850 kg aluminium to a weight of 25 kg, but similar strength). Also, three active flexure correcting mirrors are added to the system. During target acquisition, a calibration exposure is obtained from which the flexure is measured at the current instrument position. This information is fed to the piezo mirrors making sure that the object is centered on the three spectrograph slits during the observation. The near-infrared arm is cooled by liquid nitrogen. Originally a closed-cycle cooler was planned to cool the instrument, but this could induce vibrations on the

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telescope platform prohibiting VLTI observations. The optical elements are cooled to 105 K, and the optimal operating temperature of the near-infrared 2k × 2k Rockwell Hawaii 2RG MBE detector is 81 K, just above the liquid N2 temperature (used area 1k × 2k). The UV detector is a 2k × 4k E2V CCD; the VIS detector a 2k × 4k MIT/LL CCD. Given the large wavelength coverage in one shot, a compromise has to be found between the contribution of the read-out-noise in the UV (requiring a long exposure time) and the sky background (variability) in the near-infrared (short exposure time). This compromise likely results in limiting the exposure time to about 20 minutes in “staring mode”. We are currently investigating whether nodding in the near-infrared (using the telescope) can be compensated for in the other two arms using the piezo mirrors. About 75% of the costs of the X-shooter hardware, as well as labor, is funded by the external members of the consortium. ESO is responsible for the detector systems, project management, and commissioning. More than 60 people are involved in the project at nine different institutes distributed over four ESO member states and at ESO. A complete list of the contributors to the project can be found in [3]. The overall cost of the project is 6.4 MEuro and the staff effort 69 FTEs. The consortium is compensated for the project investment with guaranteed time (156 nights over a period of three years). Even with the complex distribution of work over many different sites, the X-shooter project has advanced well on a relatively short timescale: ∼ 5 yr from the official kickoff to installation at the telescope. First light of the visual spectrograph was achieved in July 2007 (Fig. 2); the NIR arm had first light in December 2007. In January 2008 integration of the full

Fig. 2 Halogen flatfields and Ar line spectra in the VIS arm. The right panels show the first light images, the left panels display simulated images using the physical model. The first light Ar spectrum corresponds to the region enclosed by the white box in the simulated image. A colour version of this figure is available at dx.doi.org/10.1007/978-1-4020-9190-2_53

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instrument has started in ESO Garching. The final delivery (the near-infrared arm) is planned for March 2008. The system test phase will conclude with the so-called Preliminary Acceptance Europe (PAE) review currently planned in June 2008. First light at the telescope is expected in September 2008. Normal operations should commence on April 1, 2009.

2 Science with X-Shooter X-shooter will have a broad and varied usage ranging from nearby intrinsically faint stars to bright sources at the edge of the Universe. The unique wavelength coverage and unprecedented efficiency opens a new observing capacity in observational astronomy. At the intermediate resolution of X-shooter 80–90% of all spectral elements are unaffected by strong sky lines, so that one can obtain sky continuum limited observations in between the sky lines within a typical exposure time. Key science cases to be addressed with X-shooter concern the study of brown dwarfs, the progenitors of supernovae Type Ia, gamma-ray bursts, quasar absorption lines, and lensed high-z galaxies. The advantage of the large wavelength coverage is that e.g. the redshift of the target does not need to be known in advance (such as in the case of GRBs); also, the study of Lyman α in high-redshift galaxies will be possible in the redshift range 1.5 < z < 15. X-shooter will complement and benefit from other major facilities in observational astrophysics operational in the same period: survey instruments like VST/OmegaCAM and VISTA working in the same wavelength range, and observatories like ALMA, JWST and GLAST operational in other observing windows.

3 Expected Instrument Performance The performance of the instrument has been predicted on the basis of measured efficiencies of the telescope, optical elements, and detectors. Compared to the efficiencies predicted at the Final Design Review in June 2006, most delivered optical components (most importantly the dichroics and the gratings) are well above specification: an efficiency > 25% in the centers of all orders outside the dichroic crossover range (from the top of the telescope to the detector). Table 1 lists the predicted efficiency at the blaze of the echelle orders averaged over each of the three arms of the instrument; the detection quantum efficiency is well over 40% in most of the UVB range. Figure 3 shows the predicted limiting AB magnitudes (S/N = 10 per resolution element for a 1 h exposure, from the top of the atmosphere to the detector), calculated using a first version of the ETC, assuming that the sky background is due to the continuum in a region free of emission lines. In the near infrared the exposure is split in 3 exposures of 20 minutes and nodding is applied. The ETC uses the as-built values for optics and detector efficiency/noise, but still contains some assumptions that need to be verified during commissioning. The decrease in efficiency

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Fig. 3 Limiting AB magnitude of X-shooter per spectral bin at S/N = 10 in a one-hour exposure. The other parameters used in this prediction are: airmass 1.2, 0.8 seeing, 1 slit, 2× binning in the spectral direction. The first version of the ESO ETC was used to compute these values (http://www.eso.org/observing/etc). A colour version of this figure is available at dx.doi.org/10.1007/978-1-4020-9190-2_53

in the UV is due to atmospheric absorption; towards the red there is a decrease in CCD efficiency and the long-wavelength side of the near-infrared is limited to the rise of the thermal background. Also the predicted spectral resolution complies with the technical specifications.

4 Data Reduction Pipeline The X-shooter data reduction software is being developed as a state-of-the-art ESO archival pipeline. About 15% of the consortium’s budget is spent on pipeline development, with the aim to deliver spectra ready for scientific analysis. The pipeline contains many novel features which are not commonly found in ESO pipelines, such as: (i) full optimal extraction of data which is distorted in both X and Y directions. The optimal extraction will be able to automatically cope with arbitrary spatial profiles; (ii) The pipeline will perform end-to-end error propagation. The pipeline includes a physical model of the instrument (provided by ESO) that allows for perobservation calibration of data, rather than relying on daytime calibration data. This should deliver an absolute calibration accuracy of 0.1 pixel for every frame taken by the instrument; (iii) A single frame sky subtraction technique is being implemented, so that the near-IR arm is useful for the long staring observations required by faint targets at the shortest wavelength ranges of X-shooter.

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References 1. B. Delabre, H. Dekker, S. D’Odorico, E. Merkle, Astronomical spectrograph design with collimator compensation of camera chromatism (4C), in Raman Scattering, Luminescence and Spectroscopic Instrumentation in Technology, ed. by F. Adar, J.E. Griffiths, J.M. Lerner. Presented at the Society of Photo-Optical Instrumentation Engineers (SPIE) Conference. Proc. SPIE, vol. 1055 (1989) p. 340, http://adsabs.harvard.edu/abs/1989SPIE.1055..340D. Provided by the SAO/NASA Astrophysics Data System 2. S. D’Odorico, H. Dekker, R. Mazzoleni, J. Vernet, I. Guinouard, P. Groot, F. Hammer, P.K. Rasmussen, L. Kaper, R. Navarro, R. Pallavicini, C. Peroux, F.M. Zerbi, X-shooter UV- to K-band intermediate-resolution high-efficiency spectrograph for the VLT: Status report at the final design review, in Ground-based and Airborne Instrumentation for Astronomy, ed. by Ian S. McLean, Masanori Iye. Presented at the Society of Photo-Optical Instrumentation Engineers (SPIE) Conference. Proc. SPIE, vol. 6269 (2006) p. 626933, doi: 10.1117/12.672969 http://adsabs.harvard.edu/abs/2006SPIE.6269E..98D. Provided by the SAO/NASA Astrophysics Data System 3. J. Vernet, H. Dekker, S. D’Odorico, E. Pallavicini, P.K. Rasmussen, L. Kaper, F. Hammer, P. Groot, Coming soon on stage: X-shooter, Messenger 130, 5 (December 2007)

KMOS and KMOS++ Ray Sharples on behalf of the KMOS Consortium

1 Introduction KMOS is a unique near-infrared multi-object integral-field spectrometer which is one of the second-generation instruments currently under construction for the VLT. The instrument is being built by a consortium of UK and German institutes working in partnership with ESO and is currently in the final design phase with commissioning on the VLT scheduled to begin in 2010. In this paper we describe the baseline instrument concept derived from the KMOS science case and briefly discuss its potential as a first light pathfinder instrument on an Extremely Large Telescope.

2 Science Case and Functional Specification The main focus of cosmological studies at the start of the 21st century is the determination of the physical processes which drive galaxy formation and evolution. To achieve this goal requires a capability to map the variations in star formation histories, spatially resolved star-formation properties, merger rates and dynamical masses of well-defined samples of galaxies across a wide range of redshifts and environments. A few of the brightest examples e.g. [1] are now being observed using single integral field unit (IFU) spectrographs on 8-m telescopes but statistical surveys of these galaxy properties even at modest redshifts will require a multi-object approach. This is the capability which will be delivered to the VLT with KMOS and is one of the key science areas which will be pushed to new limits using an Extremely Large Telescope. For any instrument to address these fundamental questions about how galaxies evolve it should: (1) have a substantial multiplex capability, commensurate with the surface density of accessible targets; (2) have the ability to obtain more than just integrated or one-dimensional information; (3) be able to resolve the relatively small velocity differences observed in rotation curves, velocity dispersions, and in merging galaxy pairs; (4) have the ability to observe several targets concentrated in a small area of sky; (5) have the capability to observe high-redshift galaxies using the R. Sharples () Department of Physics, University of Durham, Durham DH1 3LE, UK e-mail: [email protected] A. Moorwood (ed.), Science with the VLT in the ELT Era, Astrophysics and Space Science Proceedings, doi: 10.1007/978-1-4020-9190-2_54, © Springer Science + Business Media B.V. 2009

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Table 1 Baseline capabilities for the KMOS final design Requirement

FDR design

Instrument throughput (inc detector)

IZ = 28%, YJ = 33%, H = 35%, K = 33%

Sensitivity (5σ , 8 hrs)

Z = 22.4, J = 22.0, H = 21.4, K = 19.4

Wavelength coverage

0.8 to 2.5 µm

Spectral resolution (IZ, YJ, H, K)

R = 3400, 3700, 4000, 3800

Number of IFUs

24

Extent of each IFU

2.8 × 2.8 arcsec

Spatial sampling

0.2 × 0.2 arcsec

Patrol field

7.2 arcmin diameter circle

Close packing of IFUs

> 3 within 1 sq. arcmin

Closest approach of IFUs

edge-to-edge separation of 6 arcsec

well-studied rest-frame optical diagnostic features used at low redshift. These general characteristics imply a near-infrared multi-object spectrograph using deployable integral field units (dIFUs). The specific choices in delivering these capabilities involves a complex trade of cost and scope which is reflected in the baseline capabilities for KMOS listed in Table 1.

3 Instrument Description KMOS will mount on the VLT Nasmyth rotator and will use the Nasmyth A&G facilities. The top-level requirements are: (i) to support spatially-resolved (3-D) spectroscopy; (ii) to allow multiplexed spectroscopic observations; (iii) to allow observations across the near infrared atmospheric windows (an extension down to 0.85 µm has now been incorporated). The baseline design employs 24 configurable arms that position fold mirrors at user-specified locations in the Nasmyth focal plane. The sub-fields thus selected are then magnified onto 24 advanced image slicer IFUs that partition each subfield into 14 identical slices, with 14 spatial pixels along each slice. Light from the IFUs is dispersed by three identical cryogenic grating spectrometers which generate 14 × 14 spectra, each with 1000 spectral resolution elements, for all of the 24 independent sub-fields. The spectrometers each employ a single 2k × 2k substrateremoved HgCdTe detector. The goal is to employ careful design choices and advances in technology to ensure that KMOS achieves a comparable sensitivity to the current generation of single-IFU infrared spectrometers. Figure 1 shows the optical path through the instrument which is partitioned into three layers (Pickoff module, IFU module and Spectrograph module). There is also a natural 3-fold axial symmetry to many of the opto-mechanical assemblies which has allowed a phased approach to manufacture, assembly, integration and test.

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Fig. 1 Optical raytrace through four pickoff arms, their associated IFUs and one of the spectrometers. Light exiting the pickoff arms is brought to an intermediate focus using a 3-element K-mirror, which aligns the edges of all 24 IFU fields on the sky so that they can be put into a compact sparse array configuration for blind surveys of contiguous areas on the sky. A colour version of this figure is available at dx.doi.org/10.1007/978-1-4020-9190-2_54

3.1 Pickoff Module One of the key KMOS elements is the pickoff module which relays the light from 24 selected regions distributed within the patrol field to an intermediate focus position at the entrance to the integral field unit module. The method adopted for selecting these subfields uses robotic pickoff arms whose pivot points are distributed in a circle around the periphery of the patrol field and which can be driven in radial and angular motions by two stepper motors which position the pickoff mirrors with a repeatable accuracy of < 0.2 arcsec. The arms patrol in one of two layers positioned either side of the Nasmyth focal plane. The changing path length within the arm is compensated via an optical trombone which uses the same lead screw, but with a different pitch, as for the main radial motion. The pickoff module also contains a central integrating sphere which relays the light from the external flatfield and wavelength calibration lamps into the pickoff arms, and a filter wheel which acts as a focus compensation device between the different bands. The cold stop for the instrument is at the base of the arm, after which the intermediate image is formed by a K-mirror assembly which also acts to orientate the pickoff fields so that their edges are parallel on the sky. This allows mapping of larger (1 × 1 arcmin) contiguous areas in a compact close-packed arm configuration using 16 exposures.

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3.2 Integral Field Unit Module The optical design of the IFU sub-systems is based on the Advanced Image Slicer concept [2] and draws heavily on experience developed in building the GNIRS integral-field unit for Gemini South [3]. In the current design the slicer mirrors are all spherical with the same radius of curvature, and so are the pupil mirrors. The slit mirrors are toroidal with the same radius of curvature in the spectral direction, but different radii of curvature in the spatial direction. This configuration was chosen because it is well adapted to the available methods of diamond machining. Each IFU sub-module produces a 256 mm long slit containing 112 separate slices from 8 subfields which feeds a single spectrograph.

3.3 Spectrograph Module The three identical spectrographs use a single off-axis mirror to collimate the incoming light, which is then dispersed via a reflection grating and refocussed using a 6element transmissive achromatic camera. The gratings are mounted on a 5-position turret which allows optimized gratings to be used for the individual IZ, YJ, H, K bands together with a lower resolution HK grating which covers the H & K bands simultaneously [4]. The spectroscopic filter bandpasses that define our IZ (0.80 µm– 1.08 µm) and YJ (1.02 µm–1.34 µm) bands overlap at the very dark 1.06 µm window which will provide a rich source of very high redshift (z = 7.7) Lyman-α emitters from narrow-band imaging surveys with HAWK-I and VISTA for subsequent KMOS follow-up. Each spectrograph contains a Teledyne 2048 × 2048 HgCdTe array mounted on a focus stage. These detectors will be fully substrate-removed to increase the quantum efficiency at wavelengths < 1 µm. All three spectrographs are mounted in a plane perpendicular to the Nasmyth rotation axis for maximum mechanical stability. The mechanical design of the whole KMOS system is shown in Fig. 2 which emphasizes the three-fold symmetry and the advantages from a mechanical perspective of positioning common components in a single plane.

4 KMOS on an ELT The scientific case for extending the KMOS multiple integral field capability to an ELT is very strong and such instruments have already been proposed for E-ELT (EAGLE; [5]) and TMT (IRMOS; [6]) including MOAO to deliver resolved spectroscopy at high spatial resolution. Although not able to address the scientific cases which underpin these instruments, we have done a limited study of the expected performance of KMOS on the baseline 42-m E-ELT and shown that it would provide an interesting pathfinder capability for integrated spectroscopy of extended sources over a field of view up to 1.4 arcmin. With minimal changes (replacement of the field

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Fig. 2 Cutaway view of the main KMOS cryostat showing the entrance window and the pickoff arm module at the front, the IFU module in the middle and the spectrograph module at the rear. The cryostat will be an aluminium/stainless steel hybrid to reduce weight. A colour version of this figure is available at dx.doi.org/10.1007/978-1-4020-9190-2_54

flattener) KMOS would deliver a multiplex capability of 24× with pickoff fields of 0.53 arcsec diameter at 40 mas sampling. In mosaicking mode the field covered is 12 × 8 arcsec; 5σ 8 hr sensitivities down to H = 23.6 in 0.4 arcsec seeing are predicted. The idea of using existing 8-m class instrument designs as a first-light ELT capability has also been adopted by the TMT IRMS (MOSFIRE) concept and is worthy of further study.

References 1. R. Genzel et al., Nature 442, 786 (2006) 2. R. Content, Proc. SPIE 2871, 1295 (1997) 3. M. Dubbeldam, R. Content, J.R. Allington-Smith, S. Pokrovsky, D.J. Robertson, Proc. SPIE 4008, 1181 (2000) 4. I.J. Lewis, J. Lynn, W. Lau, S. Yang, M. Wells, Proc. SPIE 5492, 1395 (2004) 5. J.-G. Cuby et al., These proceedings 6. Eikenberry et al., Proc. SPIE 6269, 188 (2006)

New Science Opportunities Offered by MUSE R. Bacon, S. Bauer, S. Brau-Nogué, P. Caillier, L. Capoani, M. Carollo, T. Contini, E. Daguisé, B. Delabre, S. Dreizler, J.P. Dubois, M. Dupieux, J. Dupin, E. Emsellem, P. Ferruit, M. Francois, M. Franx, G. Gallou, J. Gerssen, B. Guiderdoni, G. Hansali, D. Hofmann, A. Jarno, A. Kelz, C. Koehler, W. Kollatschny, J. Kosmalski, F. Laurent, S. Lilly, J. Lizon, M. Loupias, C. Monstein, J. Moultaka, H. Nicklas, L. Parés, L. Pasquini, A. Pecontal, R. Pello, C. Petit, A. Manescau, R. Reiss, A. Remillieux, E. Renault, M. Roth, J. Schaye, M. Steinmetz, S. Ströbele, R. Stuik, P. Weilbacher, L. Wisotzki and H. Wozniak

1 Introduction The Multi Unit Spectroscopic Explorer MUSE [4] is one of the second generation VLT instruments. MUSE is a wide-field optical integral field spectrograph operating in the visible wavelength range with improved spatial resolution. The MUSE Consortium consists of groups at Lyon (PI institute, CRAL), Gottingen (IAG), Potsdam (AIP), Leiden (NOVA), Toulouse (LATT), Zurich (ETH) and ESO. The project is currently in its final design phase. Manufacturing, assembly and integration will start after the Final Design Review which is foreseen for late 2008. The Preliminary acceptance in Europe is scheduled for mid 2011 and the instrument shall be in operation at Paranal in 2012.

2 Top Level Requirements Imagers and spectrographs are the most common tools of optical astronomers. In most cases, astronomical observations start with imaging surveys in order to find the interesting targets and then switch to spectrographic observations in order to study the physical and/or dynamical properties of the selected object. Thanks to the excellent throughput and large format of today’s detectors, large fractions of the sky can be surveyed in depth with imagers. The most limiting factor is the spectroscopic follow-up observations, which are time consuming and tend to have small multiplex capabilities. Pre-selection of sources is then mandatory. Usually the selection R. Bacon () Observatoire de Lyon, CRAL – Université de Lyon – Université Lyon I, 9 avenue Charles André, 69230 Saint-genis-Laval, France e-mail: [email protected] A. Moorwood (ed.), Science with the VLT in the ELT Era, Astrophysics and Space Science Proceedings, doi: 10.1007/978-1-4020-9190-2_55, © Springer Science + Business Media B.V. 2009

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criterion is based on a series of multi-color images and is intended to select the appropriate spectral characteristics of the population of the searched objects. This incurs a direct cost in telescope time since more than one exposure must be made at each sky location. As another disadvantage, the selection process is never 100% efficient, and thus a fraction of time of the follow-up spectroscopy is lost due to misidentifications. The major weakness of this approach, however, is probably not the relatively low efficiency of the method, but the a priori selection of targets. This pre-selection severely biases the spectrographic observations and limits considerably the discovery space. An alternative to the classical approach is to perform simultaneously imaging and spectroscopy. The idea is to merge into one instrument the best of the two capabilities: from the imaging world its field of view and high spatial resolution; and from the spectrograph’s world its high resolving power and large spectral range. Such an instrument will overcome the difficulty inherent to the classical method. Because there is no longer the need to pre-select the sources, one can even detect objects that would not have been found or pre-selected in the pre-imaging observations. In the most extreme case, such as object with very faint continuum but relatively bright emission lines, the objects can only be detected with this instrument, however not with direct imaging techniques. To achieve this goal MUSE has been designed to meet the following top level requirements: • 3D deep field capability: the instrument must be an integral field spectrograph with a field of view and a simultaneous spectral range large enough to allow source detection. The Wide Field Mode of MUSE has such a capability. • Improved spatial resolution: the spatial resolution must boost by a factor 2 the energy within one spaxel (spatial element) with respect to seeing limited observations. This is the main driver for the second generation adaptive optics system module for MUSE (GALACSI) which is part of the VLT Adaptive Optics Facility [3]. • Survey capabilities: the improved spatial resolution must be achieved for a large sky coverage (including galactic pole) and for a wide range of seeing conditions. Such capability requires the use of laser guide stars. • Very long integration time capability. For 3D deep fields we expect to sum up to 100 hours of integration on the same field. The instrument must then be very stable and the systematics must be kept minimum. • Efficient: The instrument must have a high throughput and a maximum open science shutter time. High quality coating and a controlled environment to prevent coating ageing and dust will then be used. Only day calibration in basic operation will be used. • HST/STIS-like spatial resolution with 2D spatial coverage and improved sensitivity: In this specific mode (the Narrow Field Mode), the instrument has a diffraction limited spatial resolution in the R band in a smaller field of view.

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Table 1 MUSE top level parameters Focus

Nasmyth B UT4

Deformable secondary mirror

1170 actuators

Laser guide stars

4 × 5–10 Watts

Instrument

Integral field spectrograph

Number of IFU units

24

Detectors

4k × 4k Deep depletion CCD

Simultaneous spectral range (nominal)

480–930 nm

Simultaneous spectral range (extended)

465–930 nm

Resolving power

1750 @ 465 nm–3750 @ 930 nm

Datacube size

1570 MB

Wide field mode Field of view

1 × 1 arcmin2

Spatial sampling

0.2 × 0.2 arcsec2

Spectra/exposure

90 000

Sky coverage in AO

70% @ galactic pole

Sky coverage in AO

99% @ galactic equator

AO energy gain wrt seeing

×2

Narrow field mode Field of view

7.5 × 7.5 arcsec2

Spatial sampling

25 × 25 milliarcsec2

Spectra/exposure

90 000

Spatial resolution

5–10% Strehl Ratio @ 650 nm

3 Instrument Overview and Performances The total number of information elements (Table 1) is given by the product of the number of spaxels (90 000) with the number of spectral pixels (4000), resulting in 360 million elements in the final datacube. Such a large number of pixels is not feasible with a single piece of optics and a single detector. MUSE is thus composed of 24 identical modules, each of which consists of an advanced slicer, a spectrograph and a detector vessel with a (4k)2 detector. A series of fore-optics and splitting and relay optics is used to derotate and split the square field of view into 24 sub-fields. These are placed on the Nasmyth platform between the VLT Nasmyth focal plane and the 24 IFU modules. 3D views of the complete instrument are shown in Fig. 1. The 24 IFUs are central to MUSE. They have been designed to achieve an excellent image quality (85% enclosed energy within a 15×30 µm2 in the detector plane), and make use of innovative slicer and spectrograph concepts. The slicer, which is described in Laurent et al. [2], is an evolution of the advanced slicer concept proposed by R. Content [1]. The compact spectrograph design achieves an excellent image quality over the large spectral bandwidth of MUSE. In this design, the tilt of

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Fig. 1 General view of MUSE at the VLT Nasmyth Platform. A colour version of this figure is available at dx.doi.org/10.1007/978-1-4020-9190-2_55

the detector compensates for the axial chromatism, which then does not need to be corrected optically. This is a cost effective solution, avoiding the use of expensive optical materials, e.g. CaF2. To maintain a high throughput (40% for the whole instrument) despite the relatively large number of required surfaces, attention is paid to use state-of-the art transmission and reflection coatings. Detectors are 4k × 4k 15 µm deep depletion devices with improved quantum efficiency in the red. Furthermore we will use new volume phase holographic gratings with a high efficiency over the large (one octave) spectral range.

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4 Wide Field Mode Science The most challenging scientific and technical application, and the most important driver for the instrument design, is the study of the progenitors of normal nearby galaxies out to redshifts z > 6. These systems are extremely faint and can only be found by their Lyα emission. MUSE will be able to detect these in large numbers (≈ 15 000) through a set of nested surveys of different area and depth. The deepest survey will require very long integration (80 hrs each field) to reach a limiting flux of 3.9 × 10−19 erg s−1 cm−2 , a factor of 100 times better than what is currently achieved with narrow band imaging. These surveys will simultaneously address the following science goals: • Study of intrinsically faint galaxies at high redshift, including determination of their luminosity function and clustering properties. • Detection of Lyα emission out to the epoch of reionization, study of the cosmic web, and determination of the nature of reionization. • Study of the physics of Lyman break galaxies, including their winds and feedback to the intergalactic medium. • Spatially resolved spectroscopy of luminous distant galaxies, including lensed objects. • Search of late-forming population III objects. • Study of active nuclei at intermediate and high redshifts. • Mapping of the growth of dark matter haloes. • Identification of very faint sources detected in other bands. • Serendipitous discovery of new classes of objects. Multi-wavelength coverage of the same fields by MUSE, ALMA, and JWST will provide nearly all the measurements needed to answer the key questions of galaxy formation. At lower redshifts, MUSE will provide exquisite two-dimensional maps of the kinematics and stellar populations of normal, starburst, interacting and active galaxies in all environments, probing sub-kiloparsec scales out to well beyond the Coma cluster. These will reveal the internal substructure, uncovering the fossil record of their formation, and probe the relationship between super massive black holes and their host galaxy. MUSE will enable massive spectroscopy of the resolved stellar populations in the nearest galaxies, outperforming current capabilities by factors of over 100. This will revolutionize our understanding of stellar populations, provide a key complement to GAIA studies of the Galaxy, and a preview of what will be possible with an ELT.

5 Narrow Field Mode Science Contrary to the Wide Field Mode, the Narrow Field mode science is dedicated to detailed study of single objects at very high spatial resolution. We give in the following a few examples.

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The study of super massive black holes: During galaxy mergers, super massive black holes sink to the bottom of the potential well, forming binary systems which ‘scour out’ lower-density cores in the central regions of the remnant. Such processes should leave detect able signatures in the environment of the SMBH. Likewise, accretion of mass onto super massive black holes should trigger activity and feedback to the local regions and beyond. However, observationally, very little is known about this environment, either in terms of stellar orbital structure or chemical enrichment history. Young stellar objects: The key contribution from MUSE will be both in spectral grasp (covering key diagnostics of density, temperature and ionization) and the ability to provide very high spatial resolution over a relatively large field of view. This will allow the physical processes involved in the formation and structure of the jets to be investigated in detail. Solar System: MUSE NFM would allow observation of various bodies within our Solar System at a spatial resolution approaching that of more costly space missions. Applications are: monitoring volcanic activity on the Galilean satellites, spectral monitoring of Titan’s atmosphere, global monitoring of the atmospheres of Uranus and Neptune, internal structure and composition of comets and mineralogical surface heterogeneities of asteroids.

6 The ELT instrument pathfinder By many aspects, MUSE is a precursor of future ELT instrumentations. For example manufacturing, integration and maintenance of a large number of identical, high performance optical systems at low cost and on reasonable time scale will be a critical aspect for most of ELT instruments. MUSE will also be the first AO-assisted IFU to address a key science case of future ELTs: massive spectroscopy of resolved stellar populations in nearby galaxies, employing crowded field 3D spectroscopy over a large field-of-view.

References 1. R. Content, in Optical Telescopes of Today and Tomorrow, ed. by Arne L. Ardeberg. Proc. SPIE, vol. 2871 (Springer, 1997), pp. 1295–1305 2. F. Laurent, E. Renault, R. Bacon, B. Delabre, J.-P. Dubois, F. Hénault, J. Kosmalski, in Optomechanical Technologies for Astronomy, ed. by Eli Atad-Ettedgui, Joseph Antebi, Dietrich Lemke. Proc. SPIE, vol. 6273 (2006), p. 62732B 3. S. Ströbele, R. Arsenault, R. Bacon, R. Biasi, D. Bonaccini-Calia, M. Downing, R.D. Conzelmann, B. Delabre, R. Donaldson, M. Duchateau, S. Esposito, E. Fedrigo, D. Gallieni, W.K.P. Hackenberg, N. Hubin, M. Kasper, M. Kissler-Patig, M. Le Louarn, R. McDermid, S. Oberti, J. Paufique, A. Riccardi, R. Stuik, E. Vernet, in Advances in Adaptive Optics II, ed. by Brent L. Ellerbroek, Domenico Bonaccini Calia. Proc. SPIE, vol. 6272 (2006), p. 62720B 4. MUSE public web site: http://muse.univ-lyon1.fr

SPHERE: A ‘Planet Finder’ Instrument for the VLT D. Mouillet, J.-L. Beuzit, M. Feldt, K. Dohlen, P. Puget, F. Wildi, A. Boccaletti, T. Henning, C. Moutou, H.M. Schmid, M. Turatto, S. Udry, F. Vakili, R. Waters, A. Baruffolo, J. Charton, R. Claudi, T. Fusco, R. Gratton, N. Hubin, M. Kasper, M. Langlois, J. Pragt, R. Roelfsema and M. Saisse

SPHERE (Spectro-Polarimetric High-contrast Exoplanet Research) is a second generation instrument for the VLT optimized for the very high-contrast imaging around bright stars [1]. The primary goal is the detection and characterization of new giant planets around a variety of nearby stars. Together with the observation of early planetary systems and disks, and in association with the results of other planet search techniques, SPHERE will be a primary contributor to get a complete picture of the variety of planetary systems and to better understand their mechanisms of formation and evolution. Such results will be obtained before even more ambitious projects for the direct imaging of planets either from the ground with ELTs or from space.

1 Scientific Context and Main Instrument Goals 1.1 Primary Astrophysical Issues For essentially a dozen years now, various observation techniques (mainly radial velocity measurements and transits) have revealed a large number of planets around other stars. The demonstrated existence of such common planetary systems and also some impressive differences with respect to our system did provide essential inputs to re-visit our complete understanding and modeling of this fundamental issue of planets formation and evolution. Among the very exciting results was the quite unexpected variety of discovered systems, including for instance the hot Jupiters (massive planets very close to the stars), planets on very eccentric orbits, the frequency of multiple systems or planets in binary systems. Exo-planets studies are now enriched with a solid basis of planet detections making possible the statistical analysis of the variety of possible planet formation scenarios. Moreover, the combined use of radial velocity measurements with (primary or secondary) transit D. Mouillet () Laboratoire d’Astrophysique de Grenoble, BP 53, 38041 Grenoble cedex 9, France e-mail: [email protected] A. Moorwood (ed.), Science with the VLT in the ELT Era, Astrophysics and Space Science Proceedings, doi: 10.1007/978-1-4020-9190-2_56, © Springer Science + Business Media B.V. 2009

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measurements allow further insights of the physics of the outer atmospheres or of the inner planetary structure in some particular cases. However, the techniques mentioned here have some important limitations and suffer from strong observational biases. Direct imaging of planets should provide the required complementary information to obtain a comprehensive view of the composition, structure and evolution of the exo-planetary systems. In particular, for a better statistical analysis of extra-solar systems, direct imaging should be complementary to current techniques in terms of: • detectable star-to-planet separation: current studies are limited to data up to a few astronomical units at most, i.e. at the very edge of the expected peak region of giant planet formation. On contrary, direct imaging is easier at large separation; it should reveal the composition of outer planetary systems, which is essential to understand their overall dynamical structure, just like in our own Solar System; • stellar types: here again the current techniques are strongly biased towards cool and quiet stars (most performing for G and K old stars); direct imaging should enlarge such a sample including massive stars and also, very importantly, the case of younger systems at earlier stages of dynamical evolution. Direct imaging is not only needed to complete our view of extra-solar systems in a statistical approach, it also has the potential to characterize the physics and composition of extra-solar planet atmospheres, through the NIR spectral properties of the intrinsic emission, the spectral and polarization properties of the reflected light in case of close-in planets, and possibly their variability along the orbit. The main challenge for such direct imaging of exoplanets consists in the very large contrast between the star and the planet at very small separations, typically inside the seeing halo. The achievement of astronomical adaptive optics in the 90’s allowed to really enter the domain of high contrast imaging from the ground and the high angular imaging instruments installed on 8-m telescopes, such as VLT/NACO, demonstrated another step forward in performance. Large observing surveys, including in particular the favorable cases of young nearby stellar associations, are underway with the successful first detections of brown dwarfs and of planetarymass companions. However, with typical contrast performance of m 10 at 0.5 , the current capabilities are only marginally opening the domain of planet detections, with the access to Jupiter-mass companions in the case of very young systems only (a few 106 yrs when the planet is still quite warm and not too faint) and at separations so large that the presence of planets is unlikely.

1.2 Main Instrument Goals As a second generation instrument on the VLT, dedicated to high contrast imaging, SPHERE will offer greatly enhanced capabilities (a gain of two orders of magnitudes in contrast with respect to existing instruments) to provide a clear view of

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Fig. 1 Contrast performance as a function of separation: comparison between SPHERE and existing instruments. Upper black curves corresponds to the published results, obtained on existing instruments in NIR. Lower red curves indicate the estimated level of performance for SPHERE in NIR, in the photon noise-limited regime for a J = 8 star in 1 hr (dashed line) and in the ultimate calibration limit in the case of the brightest stars (full line). All curves correspond to a 5-σ detection of a point-like companion, and are expressed in terms of magnitude difference in the considered narrow band spectral channels. A colour version of this figure is available at dx.doi.org/10.1007/978-1-4020-9190-2_56

the frequency of giant planets in wide orbits, searched for in a list of hundreds of potential stellar targets. Furthermore, for detected planets, SPHERE will provide information on NIR spectral properties of the atmosphere at low (R 50) or medium (R 500) resolution; a few planets shining by reflecting stellar light might be detected, if present around very nearby stars, using the SPHERE polarimetric visible channel (ZIMPOL). To achieve such astrophysical purposes, the main top level requirements of SPHERE include: • Very high contrast in NIR, in the stellar seeing halo: 10−6 to 10−8 at 0.5 , down to inner working angles of 0.1 . • Such contrast achievable in a typical observing time of 1 hr for a large number (hundreds to thousands) and a large variety (including various ages and various stellar types) of targets, corresponding to limiting magnitudes of V 10 and J 8. • Spectral information covering the main planetary features from 0.95 to 1.7 µm at low resolution for first detection with possible extension to 2.32 µm or at medium resolution for further characterization.

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• Operation plan compatible with a large survey of hundreds of stars: with small overheads, highly stable calibrations, and simultaneous covering of the complete 0.95–1.7 µm range with two complementary NIR instruments in parallel (an integral field spectrograph and a dual imager). • High contrast dual polarization imager (ZIMPOL) for the detection of the planet reflected light at short separations and high angular resolution (diffractionlimited) in the visible (0.6–0.9 µm). These instrumental capabilities will also allow to make great advances in related areas of study such as brown dwarfs and stellar and planetary formation processes (via imaging of inner disks at various evolutionary stages).

2 System Overview To fulfill these requirements, SPHERE is divided into 4 sub-systems [2]: • The Common Path and Infrastructure (CPI) supports the other sub-systems and provide a very stable and accurate coronagraphic optical beam, in visible and NIR, to the other imaging sub-systems. A critical part is the high order adaptive optics correction stage, controlled at a temporal frequency of 1.2 kHz. The servo loop ensures very low turbulence residuals of the corrected modes (< 3 mas jitter and < 60 nm rms for other modes) but also the pupil stability. A de-rotator makes possible to control either the field orientation or the telescope pupil orientation as seen by the instrument at the Nasmyth focus. CPI also includes both visible and NIR coronagraphs, including four-quadrant phase devices or apodized-pupil coronagraphs to provide both a high stellar extinction and a very small inner working angle down to one or a few λ/D. • The Integral Field Spectrograph (IFS) can be used either over the 0.95–1.35 µm spectral range with a resolution of 50 (simultaneously with the other NIR instrument in survey mode) or extended up to 1.7 µm. The number of independent spectral channels with essentially no differential wavefront errors will provide the deepest imaging performance over the central 1.7 field of view. • The Infra-Red Dual-beam Imaging and Spectroscopy (IRDIS) sub-system complementarily covers a larger field of view (11 ) and a larger spectral domain up to 2.32 µm. When used simultaneously with IFS in planet detection survey mode, it will probe the main planetary methane absorption feature at 1.6 µm, in simultaneous dual imaging. A set of dual imaging filters are defined to cover the whole spectral domain and various types of expected planetary spectral features. Characterization of detected companions is possible through slit spectroscopy (R = 50, 500); a dual polarization imaging mode is also proposed for the detection and characterization of reflected light (on dust disks for instance). • The Zurich Imaging Polarimeter (ZIMPOL) is a high-precision imaging polarimeter working in the visible (0.6–0.9 µm). Its principle is based on a fast polarization modulation (in the kHz range) with a ferro-electric retarder combined synchronously with charge transfer every second row of a modified CCD. Data

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processing extracts the images corresponding to two polarization states, obtained essentially simultaneously through the same optics and on the same detector pixels, leading to an extremely high differential polarization accuracy (better than 10−5 ). The CCD will cover a Nyquist-sampled field of 3 , with a possible offset up to 4 in radius. SPHERE is an instrument designed and built by a consortium consisting of LAOG, MPIA, LAM, LESIA, LUAN, INAF, Observatoire de Genève, ETH, NOVA, ONERA and ASTRON, in collaboration with ESO. The preliminary design phase has been closed in late 2007 and the first light is foreseen in 2011.

References 1. J.-L. Beuzit, M. Feldt, K. Dohlen et al., Messenger 125, 29 (2006) 2. F. Wildi, J.-L. Beuzit, M. Feldt et al., in The SPHERE Exoplanet Imager: Status Report at PDR, SPIE Conf., San Diego, 2007, in press

Milli-arcsecond Astrophysics with VSI, the VLTI Spectro-imager in the ELT Era F. Malbet, D. Buscher, G. Weigelt, P. Garcia, M. Gai, D. Lorenzetti, J. Surdej, J. Hron, R. Neuhäuser, P. Kern, L. Jocou, J.-P. Berger, O. Absil, U. Beckmann, L. Corcione, G. Duvert, M. Filho, P. Labeye, E. Le Coarer, G. Li Causi, J. Lima, K. Perraut, E. Tatulli, E. Thiébaut, J. Young, G. Zins, A. Amorim, B. Aringer, T. Beckert, M. Benisty, X. Bonfils, A. Chelli, O. Chesneau, A. Chiavassa, R. Corradi, M. de Becker, A. Delboulbé, G. Duchêne, T. Forveille, C. Haniff, E. Herwats, K.-H. Hofmann, J.-B. Le Bouquin, S. Ligori, D. Loreggia, A. Marconi, A. Moitinho, B. Nisini, P.-O. Petrucci, J. Rebordao, R. Speziali, L. Testi and F. Vitali

Abstract Nowadays, compact sources relatively warm like surfaces of nearby stars, circumstellar environments of stars from early stages to the most evolved ones and surroundings of active galactic nuclei can be investigated at milli-arcsecond scales only with the VLT in its interferometric mode. We propose a spectro-imager, named VSI (VLTI spectro-imager), which is capable to probe these sources both over spatial and spectral scales in the near-infrared domain. This instrument will provide information complementary to what is obtained at the same time with ALMA at different wavelengths and the extreme large telescopes.

At the beginning of the 21st century, infrared observations performed at the milliarcsecond scale are essential for many astrophysical investigations either to compare the same physical phenomena at different wavelengths (like sources already observed with the VLBI or soon to be observed by ALMA) or to get finer details on observations carried out with the Hubble Space Telescope (HST) or 10-m class telescopes equipped with adaptive optics. The astrophysical science cases at milliarcsecond scales which cover from planetary physics to extragalactic studies can only be studied using interferometric aperture synthesis imaging with several optical telescopes. In this respect, the Very Large Telescope (VLT) observatory of the European Southern Observatory (ESO) is a unique site world-wide with 4 × 8-m unit telescopes (UTs), 4 × 1.8-m auxiliary telescopes (ATs) and all the required infrastructure, in particular delay lines (DLs), to combine up to 6 telescopes. The VLT Interferometer (VLTI) infrastructure can be directly compared to the Plateau de Bure Interferometer (PdBI) which combines 6 × 15-m antenna over 500-m in the F. Malbet () CNRS, Laboratoire d’Astrophysique de Grenoble, Université J. Fourier, UMR 5571, BP 53, 38041 Grenoble cedex 9, France e-mail: [email protected] A. Moorwood (ed.), Science with the VLT in the ELT Era, Astrophysics and Space Science Proceedings, doi: 10.1007/978-1-4020-9190-2_57, © Springer Science + Business Media B.V. 2009

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millimeter-wave domain. The quality of the foreseen images can be directly compared to the images provided by the PdBI. However, the angular resolution of the VLTI is a few hundred times higher due to the observation at shorter wavelengths. The large apertures of the VLTI telescopes and the availability of fringe tracking allow sensitivity and spectral resolution to be added to the imaging capability of the VLTI. We present in this contribution an instrument called the VLTI spectro-imager (VSI) as a response [1] to the ESO call for phase A proposals for second generation VLTI instruments.

1 VSI Overview The VLTI Spectro Imager will provide the ESO community with spectrally-resolved near-infrared images at angular resolutions down to 1.1 milli-arcsecond and spectral resolutions up to R = 12 000. Targets as faint as K = 13 will be imaged without requiring a brighter nearby reference object; fainter targets can be accessed if a suitable off-axis reference is available. This unique combination of high-dynamic-range imaging at high angular resolution and high spectral resolution for a wide range of targets enables a scientific program which will serve a broad user community within ESO and at the same time provide the opportunity for breakthroughs in many areas at the forefront of astrophysics. A great advantage of VSI is that it will provide these new capabilities while using technologies which have extensively been tested in the past and while requiring little in terms of new infrastructure on the VLTI. At the same time, VSI will be capable to make maximum use of the new infrastructure as it becomes available. VSI provides the VLTI with an instrument capable of combining up to 8 telescopes, enabling rapid imaging through measurement of up to 28 visibilities in hundreds of wavelength channels within a few minutes. Operations with less than 8 telescopes is the scope of the first phases of VSI. Three development phases are foreseen: VSI4 combining 4 telescopes (UTs or ATs), VSI6 combining 6 telescopes (4 UTS + 2 ATs or 4 ATs + 2 UTS and eventually 6 ATs), and perhaps ultimately, in the long-run, VSI8 combining 8 telescopes (4 UTs + 4 ATs or eventually 8 ATs). The current studies were focused on a 4-telescope version with an upgrade to a 6-telescope one. The instrument contains its own fringe tracker and wavefront control in order to reduce the constraints on the VLTI infrastructure and maximize the scientific return.

2 Science Cases for VSI The high level specifications of the instrument are derived from science cases based on the capability to reconstruct for the milli-arcsecond-resolution images of a wide range of targets. These science cases are detailed below.

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• Formation of stars and planets. The early evolution of stars and the initial conditions for planet formation are determined by the interplay between accretion and outflow processes. Due the small spatial scales where these processes take place, very little is known about the actual physical and chemical mechanisms at work. Interferometric imaging at 1 milli-arcsecond spatial resolution will directly probe the regions responsible for the bulk of excess continuum emission from these objects, therefore constraining the currently highly degenerate models for the spectral energy distribution. In the emission lines a variety of processes will be probed, in particular outflow and accretion magnetospheres. The inner few AUs of planetary systems will also be studied, providing additional information on their formation and evolution processes, as well as on the physics of extrasolar planets. • Imaging stellar surfaces. Optical and near-infrared imaging instruments provide a powerful means to resolve stellar features of the generally patchy surfaces of stars throughout the Hertzsprung-Russell diagram. Optical/infrared interferometry has already proved its ability to derive surface structure parameters such as limb darkening or other atmospheric parameters. VSI, as an imaging device, is of strong interest to study various specific features such as vertical and horizontal temperature profiles and abundance inhomogeneities, and to detect their variability as the star rotates. This will provide important keys to address stellar activity processes, mass-loss events, magneto-hydrodynamic mechanisms, pulsation and stellar evolution. • Evolved stars, stellar remnants & stellar winds. HST and ground-based observations revealed that the geometry of young and evolved planetary nebulae and related objects (e.g., nebulae around symbiotic stars) show an incredible variety of elliptical, bi-polar, multi-polar, point-symmetrical, and highly collimated (including jets) structures. The proposed mechanisms explaining the observed geometries (disks, magnetohydrodynamics collimation and binarity) are within the grasp of interferometric imaging at 1 mas resolution. Extreme cases of evolved stars are stellar black holes. In microquasars, the stellar black-hole accretes mass from a donor. The interest of these systems lies in the small spatial scales and high multi-wavelength variability. Milliarcsecond imaging in the near-infrared will allow disentangling between dust and jet synchrotron emission, comparison of the observed morphology with radio maps and correlation of the morphology with the variable X-ray spectral states. • Active galactic nuclei & supermassive black holes. AGN consist of complex systems composed of different interacting parts powered by accretion onto the central supermassive black hole. The imaging capability will permit the study of the geometry and dust composition of the obscuring torus and the testing of radiative transfer models. Milli-arcsecond resolution imaging will allow us to probe the collimation at the base of the jet and the energy distribution of the emitted radiation. Supermassive black hole masses in nearby (active) galaxies can be measured and it will be possible to detect general relativistic effects for the stellar orbits closer to the galactic center black hole. The wavelength-dependent differentialphase variation of broad emission lines will provide strong constraints on the size

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and geometry of the Broad Line Region (BLR). It will then be possible to establish a secure size-luminosity relation for the BLR, a fundamental ingredient to measure supermassive black hole masses at high redshift. We have shown that this astrophysical program [2, 3] could provide the premises for a legacy program at the VLTI.

3 Instrument Concept The phase A study has led to an instrument concept consisting of: • Integrated optics multi-way beam combiners providing high-stability visibility and closure-phase measurements on multiple baselines; • A cooled spectrograph providing resolutions between R = 100 and R = 12 000 over the J, H, or K bands; • An integrated high-sensitivity switchable H/K fringe tracker capable of real-time cophasing or coherencing of the beams from faint or resolved sources; • Hardware and software to enable the instrument to be aligned, calibrated and operated with minimum staff overhead. These features act in synergy to provide a scientific capability which is a step beyond existing instruments. Compared to the single closure phase measured by AMBER, the 3 independent closure phases available by VSI4, the 10 independent closure phases measured by VSI6 and the 21 independent closure phases measured by VSI8 will make true interferometric imaging, as opposed to simply measuring visibilities, a routine process at the VLTI. The capability to cophase on targets up to K = 10 will allow long integrations at high spectral resolutions for large classes of previously inaccessible targets. A system analysis of VSI has allowed the high level specifications of the system to be defined, the external constraints to be clarified and the functional analysis to be performed. The system design [4] features 4 main assemblies: the science instrument (SI), the fringe tracker (FT), the common path (CP) and the calibration and alignment tools (CAT). The global implementation is presented in Fig. 1. The optics design of the science instrument features beam combination using single mode fibers, an integrated optics chip and 4 spectral resolutions through a cooled spectrograph. The common path includes low-order adaptive optics (with the current knowledge reduced to only tip-tilt corrections). VSI also features an internal fringe tracker. These servo-loop systems relax the constraints on the VLTI interfaces by allowing for servo optical path length differences and optimize the fiber injection of the input beams to the required level. An internal optical switchyard allows the operator to choose the best configuration of the VLTI co-phasing scheme in order to perform phase bootstrapping for the longest baseline on over-resolved objects. Three infrared science detectors are implemented in the instrument, one for the Science Instrument, one for the fringe tracker, and one for the tip-tilt sensor. The instrument features 3 cryogenic vessels.

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Fig. 1 General implementation of the VSI instrument. A colour version of this figure is available at dx.doi.org/10.1007/978-1-4020-9190-2_57

An important part of the instrument is the control system which includes several servo-loop controls and management of the observing software. The science software manages both data processing and image reconstruction since one of the products of VSI will be a reconstructed image like for the millimeter-wave interferometers. The instrument development includes a plan for assembly, integration and tests in Europe and in Paranal. An instrument preliminary analysis report [5] discusses several important issues such as the comparison between the integrated optics and bulk optics solutions, the standard 4- and 6-telescope VLTI array for imaging, the proposed implementation of M12 mirrors to achieve these configurations with VSI4 and VSI6, implication of using an heterogeneous array and analysis of the thermal background. The needs for future VLTI infrastructure can be summarized [6] in an increasing order of completeness as: (1) Interferometry Supervisor Software upgrade from 4telescope version to a 6-telescope version; (2) AT5 and AT6, 2 additional ATs. On a longer term, 8T combination at the VLTI could be foreseen but this is not a VSI priority.

4 VSI Project Management For VSI4, the management plan [7] identifies a total cost of 3986 k Euros for hardware and a manpower of 87 FTEs over 4 years before the commissioning begins. Since the instrument is designed from the beginning for maximum VLTI capacity, the VSI6 version would cost only 385 k Euros and 6 FTEs in addition to the VSI4 version.

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Acknowledgement The VSI phase A study has benefited from a contract by ESO, and support from JRA4 of OPTICON and from CNRS/INSU.

References 1. F. Malbet, P. Kern, J.-P. Berger et al., in doc. VSI-PRO-001, issue 1.0, in VSI Technical Proposal in Response to ESO Call for Phase-A Proposals for 2nd Generation VLTI Instruments (2006) 2. P. Garcia, O. Absil, F. Baron et al., in doc. VSI-PRO-002, issue 1.0, in Science Cases for the VSI Instrument in Response to ESO Call for Phase-A Proposals for 2nd Generation VLTI Instruments (2006) 3. P. Garcia, J.-P. Berger, F. Malbet et al., in doc. VLT-SPE-VSI-15870-4335, issue 1.0, in VSI Phase A Document Package (2007). Science cases 4. L. Jocou, P. Kern, J.-P. Berger et al., in doc. VLT-SPE-VSI-15870-4339, issue 1.0, in VSI Phase A Document Package (2007). System design 5. J.-P. Berger, F. Malbet, P. Kern et al., in doc. VLT-TRE-VSI-15870-4341, issue 1.0, in VSI Phase A Document Package (2007). Preliminary instrument analysis report 6. F. Malbet, P. Kern, D. Buscher et al., in doc. VLT-SPE-VSI-15870-4335, issue 1.0, in VSI Phase A Document Package (2007). Execcutive summary 7. P. Kern, F. Malbet, in doc. VLT-PLA-VSI-15870-4338, issue 1.0, in VSI Phase A Document Package (2007). Management plan

Prospects for Near-infrared Characterisation of Hot Jupiters with the VLTI Spectro-imager (VSI) S. Renard, O. Absil, J.-P. Berger, X. Bonfils, T. Forveille and F. Malbet

1 Introduction Since the discovery of the first exoplanet around 51 Pegasi, the study of planetary systems receives an increasing attention, with the development and test of more and more detection techniques. Among the direct detection techniques, interferometry is one of the most promising for the near future. It already provides the required angular resolution, but the dynamic range needs to be improved. The detection and characterisation of extrasolar planets is one of the main science cases of the 2nd generation VLTI Spectro-Imager instrument (VSI). The goal of this work is to study the feasibility of obtaining near-infrared spectra of bright extrasolar giant planets (EGP) with VSI.

2 Method To determine the feasibility of EGP spectroscopy with VSI, we simulate interferometric observations of several EGPs that have been discovered by radial velocity surveys. Our detection method is based solely on the measurement of the closure phase with the four possible triplets of UTs. 1. Theoretical observed phase: The formula used to simulate closure phase measurements is based on the bispectrum (Bij k ): CPij k = arg Bij k where Bij k = Vij Vj k Vki using the visibility model of a high contrast binary with a flux ration 1 rλ : Vij = 1+r [V exp(−iπ λ1 (uij x + vij y)) + rλ Vp exp(iπ λ1 (uij x + vij y))]. λ 2. Simulated observations: The observations are computed with the four 8-meter telescopes of the VLT for three successive nights, four hours of observation per night and one exposure of ten minutes per hour. We take into account fundamental noises (detector and photon noises) and we use a low spectral resolution (R = 100). S. Renard () Laboratoire d’Astrophysique de l’Observatoire de Grenoble (LAOG), Grenoble, France e-mail: [email protected] A. Moorwood (ed.), Science with the VLT in the ELT Era, Astrophysics and Space Science Proceedings, doi: 10.1007/978-1-4020-9190-2_58, © Springer Science + Business Media B.V. 2009

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Fig. 1 Results for the synthetic planetary spectra from Burrows. From left to right: τ Boo, 51 Pegasi and HD209458; first line: K-Band, second line: H-Band. In black: theoretical curve; in colour: reconstructed spectra from the simulated observations. A colour version of this figure is available at dx.doi.org/10.1007/978-1-4020-9190-2_58

3. Fitting planetary parameters to the data: Standard χ 2 fitting methods are used to derive the planetary radius and its unknown orbital parameters (inclination and position angle) from the simulated data. The capability of VSI to constrain the chemical composition and atmospheric physics of EGPs is evaluated by using synthetic planetary spectra from Burrows’ website (http://zenith.as.arizona.edu/~burrows/sbh/sbh.html) (Fig. 1). We use a wavelength-dependant radius in the fitting procedure in order to recover planetary spectra.

3 Discussion and Conclusions The quality of the reconstructed spectra shows that such observations would strongly constrain the planetary temperature and albedo, the energy redistribution mechanisms, as well as the chemical composition of their atmospheres. We note that the simulated observations are more successful and constraining if the star-planet system is close to the observer (= 20 pc). For such targets, the VLTI angular resolution matches well the star-planet separation and the planet is bright enough to provide a good SNR. It is thus recommended to choose the closer systems as the first targets of VSI.

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Systematic errors, not included in our simulations, could be a serious limitation to these performance estimations. The use of integrated optics is however expected to provide the required instrumental stability (around 10−4 ) to enable the first thorough characterisation of extrasolar planetary spectra in the near-infrared.

MATISSE B. Lopez, S. Lagarde, S. Wolf, W. Jaffe, G. Weigelt, P. Antonelli, P. Abraham, J.-Ch. Augereau, U. Beckman, J. Behrend, N. Berruyer, Y. Bresson, O. Chesneau, J.M. Clausse, C. Connot, W.C. Danchi, M. Delbo, K. Demyk, A. Domiciano, M. Dugué, A. Glazenborg, U. Graser, H. Hanenburg, Th. Henning, M. Heininger, K.-H. Hofmann, Y. Hugues, S. Jankov, S. Kraus, W. Laun, Ch. Leinert, H. Linz, A. Matter, Ph. Mathias, K. Meisenheimer, J.-L. Menut, F. Millour, L. Mosoni, U. Neumann, A. Niedzielski, E. Nussbaum, R. Petrov, Th. Ratzka, S. Robbe-Dubois, A. Roussel, D. Schertl, F.-X. Schmider, B. Stecklum, E. Thiebaut, F. Vakili, K. Wagner, L.B.F.M. Waters, O. Absil, J. Hron, N. Nardetto, J. Olofsson, B. Valat, M. Vannier, B. Goldman, S. Hönig and W.D. Cotton

Abstract MATISSE is foreseen as a mid-infrared spectro-interferometer combining the beams of up to four UTs/ATs of the Very Large Telescope Interferometer (VLTI). MATISSE will measure closure phase relations and thus offer an efficient capability for image reconstruction in the L, M and N bands of the mid-infrared domain.

1 Instrument Characteristics and Sensitivity The instrument characteristics derive from the goals defined in the ‘Statement of Work and Technical Specifications for the Phase A Study of the MATISSE Instrument’ (VLT-SPE-ESO-15860-0001), from the ‘MATISSE Phase A Science Case’ study (VLT-TRE-MAT-15860-4325), and, from the feasibility study and choices developed by the ‘MATISSE Phase A System Design’ (VLT-TRE-MAT-15860-4327). MATISSE has the following characteristics: • 4 Telescope beam combiner. 3T or 2T configurations possible. • Spectral Coverage: – Sensitivity optimized for L and N bands. – Other: M band. – Simultaneous observations in L and N bands. • Spectral Resolution: – L (and M): Low = 30, Medium = 300–500, High = 750–1500. – N : Low = 30, High = 300. • Other: Field rotation module; 2D mode for acquisition; Polarization filters for the L band; Calibration devices (including Beam Commuter). B. Lopez () Observatoire de la Côte d’Azur, Nice, France e-mail: [email protected] A. Moorwood (ed.), Science with the VLT in the ELT Era, Astrophysics and Space Science Proceedings, doi: 10.1007/978-1-4020-9190-2_59, © Springer Science + Business Media B.V. 2009

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Table 1 MATISSE sensitivity to the coherent flux and corresponding magnitudes

Without fringe sensing or tracking

N coherent flux limit

L coherent flux limit

UT

AT

UT

AT

0.4 Jy

8 Jy

0.03 Jy

0.6 Jy

N =5

N = 1.7

L = 9.9

L = 6.7

0.04 Jy

1 Jy

0.03 Jy

0.6 Jy

N = 7.2

N =4

L = 9.9

L = 6.7

With external on-axis

0.02 Jy

0.4 Jy

0.02 Jy

0.4 Jy

K band fringe tracking

N = 8.2

N =5

L = 10.4

L = 7.1

With external off-axis

0.02 Jy

0.4 Jy

0.001 Jy

0.02 Jy

K band fringe tracking

N = 8.2

N =5

L = 13.6

L = 10.4

With internal L band fringe sensing

The performance goal in term of sensitivity is shown in Table 1.

2 Summary: Science Cases In this section a summary of the potential science programs offered by MATISSE is given (see also Wolf et al. in these proceedings). The related key issues are listed for the primary science cases, which are, the ‘Star and Planet Formation’ and the ‘Active Galactic Nuclei’.

2.1 Star and Planet Formation 1. Low-mass star and planet formation: (a) Complex disk structures on large (∼ 100 AU) and small scale (∼ 1 AU); Transitional objects: Status of inner disk clearing. (b) Mineralogy of proto-planetary disks; Evidence for dust grain growth and sedimentation. (c) Characteristic structures in disks: Evidence for the presence of giant protoplanets. (d) The binary mode of star formation: Circumbinary and circumstellar disks; Disk alignment and early evolution of binary systems. (e) Nature of outbursting YSOs: Structure of young accretion disks. 2. Late stage of planet formation—Debris disks: (a) The outcome of planetesimal collisions and exo-comets evaporation: Dust grain properties and disk geometry. (b) Complex spatial disk structure—direct indicators for the presence of planets. (c) Characterization of Darwin/TPF targets.

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3. Massive star formation: (a) Spatial distribution of the gas (carbon monoxide and hydrogen) and dust (silicates/graphite and CO ice) in the typically complex and distant highmass star-forming regions. (b) Link between low and high-mass star formation: Search and characterization of accretion disks around young massive (proto)stars.

2.2 Active Galactic Nuclei Hydro-dynamical models of the central gas and dust distribution in AGN show a dense inner disk (supported by angular momentum) and an outer filamentary structure—the torus. 1. Can we establish the existence of the dense inner disks? Are the disks present in both Seyfert 1 and 2 galaxies? 2. Can we find direct evidence that tori are clumpy or detect filamentary structures. Outflow phenomena (supersonic winds, jets) are connected with most kinds of AGN activity. 3. To which extend is the torus structure regulated by the outflows? 4. What fraction of the dust emission from within the inner few parsecs of an AGN is emitted by the torus, what by dust entrained in the outflows?

2.3 Secondary Science Cases The secondary science cases include: the ‘Evolved Stars’, the ‘Solar System Minor Bodies’ and the ‘Extrasolar Planets’.

3 Optical Layout and Functions The Field Rotation Compensator optimizes the field orientation in regard to the slits of the spectrograph and compensates the field and the polarization rotation existing in the focal laboratory. The Beam Commuting Device is used as a calibration tool for instrumental effects (including chromatic and detector effects). The beam combination mode is “multi-axial”. The fringes are sampled by 4 pixels and the number of fringes per PSF should be 3. It leads to have 12 pixels per λ/D in the spatial direction. In order to limit the size of the spectrum, an image magnification (and a pupil reduction) in the spatial direction is necessary by a factor 4. This magnification is made thanks to cylindrical mirrors called Anamorphic Optics. A spectral separation between the L (and M) band and the N band is necessary because the materials of optics and the detectors are different. This separation is

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Fig. 1 Optical layout of MATISSE. A colour version of this figure is available at dx.doi.org/10.1007/978-1-4020-9190-2_59

made by a set of 4 dichrocs (Spectral Separator). The L (and M) band is transmitted; the N band is reflected to the flat mirror. Two independent sets of Delay Lines (one for the L band an other for the N band) made of cat’s eyes transfer the pupil at the cold pupil mask level of the spectrograph and allow a fine OPD adjustment. At the entrance of the cold optics, the number of beam is 8 (4 for the L band and 4 for the N band). The distance between the beams is 112.5 mm. The beam size is 18 × 4.5 mm (due to the anamorphic optics). The Shutter selects the chosen beams (permitting observations with 2, 3 or 4 telescopes). Pupil masks reduce the thermal background (in particular the thermal background emitted by the optical mounts outside the cryostat) and avoids parasitic light. Pinholes or slits are used as Spatial Filter. A study has shown that a pinhole (with a size between λ/D and 1.5 λ/D) is almost equivalent to an optical fiber according to the performance of MACAO. In the dispersion direction, the pinholes or the slits are needed for the spectral resolution and reduce also the thermal background. Focusing mirrors produce an intermediate image at the spatial filter level. After the spatial filter, collimating mirrors collimates the beams. The set of focusing and collimating mirrors is called Re-Imager. To avoid crosstalk between the beams, which would produce parasitic fringes due to the diffraction, 4 sets of focusing optics, pinholes and colli-

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mating optics per spectral band are necessary instead of a common one. Cat’s eyes (called Pupil Transfer Optics) relocate the pupil plan to the image positioning device level in order to limit the size of the cold optics, avoid vignetting and hence increase the field of view of the instrument. The Pair Creator module made of beam splitters and mirrors separates the 4 beams in 6 pairs of beam. The OPD adjustment of 3 pairs can be made by the warm Delay Lines. The OPD adjustment of the 3 other pairs has to be made by delay lines (OPD Corrector) located after the Pair Creator. The π Phase Shifter is made of beam splitters. This module produces the π phase difference between the two outputs of the beam splitter (like in MIDI for example). A shift between the beams is introduced and dispersed fringes are produced on the detector (“multi-axial combination”) thanks to the Camera Optics. The value of this separation is 3 times the pupil size in order to obtain 6 fringes per 2 λ/D. In case of observations with 4 telescopes, 12 images are produced (2 images per pair and 6 pairs of beam). 12 tilted flat mirrors (called Image Positioning Device) position all these images at different locations on the detector. Spectral Filters or/and Polarizers can be inserted in order to reduce the spectral band (for 2 dimensional image observation during acquisition for example), to avoid spectral order overlap (with the use of grism for example) or to select a polarization direction. The Dispersive Optics module is made of prisms (low spectral resolution R = 30) and grisms (higher spectral resolution R > 300). A lens can be added also to be able to visualize the pupil during alignment. The Camera Optics combine the beams on the detectors. The foreseen Detectors are the Hawaii 2RG from Rockwell for the L band and Aquarius from Raytheon for the N band.

MATISSE Science Cases S. Wolf, B. Lopez, W. Jaffe, G. Weigelt, J.-Ch. Augereau, N. Berruyer, O. Chesneau, W.C. Danchi, M. Delbo, K. Demyk, A. Domiciano, Th. Henning, K.-H. Hofmann, S. Kraus, Ch. Leinert, H. Linz, Ph. Mathias, K. Meisenheimer, J.-L. Menut, F. Millour, L. Mosoni, A. Niedzielski, R. Petrov, Th. Ratzka, B. Stecklum, E. Thiebaut, F. Vakili, L.B.F.M. Waters, O. Absil, J. Hron, S. Lagarde, A. Matter, N. Nardetto, J. Olofsson, B. Valat, M. Vannier and MATISSE Science team

Abstract MATISSE is foreseen as a mid-infrared spectro-interferometric instrument combining the beams of up to four UTs/ATs of the Very Large Telescope Interferometer (VLTI). MATISSE will measure closure phase relations and thus offer an efficient capability for image reconstruction. In addition to this, MATISSE will open 2 new observing windows at the VLTI: the L and M band in addition to the N band. Furthermore, the instrument will offer the possibility to perform simultaneous observations in separate bands. MATISSE will also provide several spectroscopic modes. In summary, MATISSE can be seen as a successor of MIDI by providing imaging capabilities in the mid-infrared domain (for a more detailed description of MATISSE see Lopez et al., these proceedings).

Limitation of Current Instrumentation—The Need for MATISSE The European Southern Observatory (ESO) has recently opened a new era of astronomy for Europe. Mid-infrared interferometry with the Mid-Infrared Interferometric Instrument for the VLTI, MIDI, operating since the end of the year 2002, allows spatially resolved observations of emission regions of hot dust in circumstellar disks, AGB stars, winds of hot stars, and the tori of AGNs with a resolution of 10–20 mas. It has been proven to be very successful in interferometric spectroscopic observations since it allows for comparing the chemical composition of dust on very different spatial scales. However, the investigation of small-scale spatial structures in general, and the quantitative analysis of spectroscopic observations in particular, are strongly limited due to the small number of visibility points measured in a reasonable amount of time (over one or a few nights) and due to the lack of phase information. Furthermore, the comparison of observed visibilities with model visibility is currently the main strategy for the interpretation of MIDI observations. The selection of realistic models and the removal of possible ambiguities in the model fitting and resulting interpretations require the reconstruction of model-independent images. S. Wolf () Max-Planck-Institut für Astronomie, Koenigstuhl 17, 69117 Heidelberg, Germany e-mail: [email protected] url: http://www.obs-nice.fr/matisse A. Moorwood (ed.), Science with the VLT in the ELT Era, Astrophysics and Space Science Proceedings, doi: 10.1007/978-1-4020-9190-2_60, © Springer Science + Business Media B.V. 2009

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Fig. 1 MATISSE performance. A colour version of this figure is available at dx.doi.org/10.1007/ 978-1-4020-9190-2_60

In addition, the classical analysis usually relies on two-dimensional models with rotation symmetry. In contrast, MATISSE will for the very first time allow image reconstruction of the small-scale regions traced with MIDI and thus finally allow an investigation of these structures based on an unprecedented level of constraints. In addition to the very large spectral coverage, our studies underline the need for a higher spectral resolution. MATISSE will allow one to perform interferometric spectroscopy with three different spectroscopic resolutions in the range of R ∼ 30–1500, providing the basis for a fundamental analysis of the composition of gases and dust grains in various environments. MATISSE—A unique Instrument In most astrophysical domains which require a multi-wavelength approach, MATISSE will be a perfect complement of forthcoming high angular resolution facilities such as the Atacama Large Millimeter Array (ALMA). MATISSE covers the mid-infrared spectral domain, between the near-infrared domain, for which many interferometric facilities are developed, and (sub)millimeter wavelengths at which ALMA will operate. With the extended wavelength coverage from the L to the N band, MATISSE will not only allow one to trace different spatial regions of the targeted objects, but also different physical processes and thus provide insights into previously unexplored areas, such as the investigation of the distribution of volatiles in addition to that of the dust.

GRAVITY: Microarcsecond Astrometry and Deep Interferometric Imaging with the VLT F. Eisenhauer, G. Perrin, W. Brandner, C. Straubmeier, A. Böhm, H. Baumeister, F. Cassaing, Y. Clénet, K. Dodds-Eden, A. Eckart, E. Gendron, R. Genzel, S. Gillessen, A. Gräter, C. Gueriau, N. Hamaus, X. Haubois, M. Haug, T. Henning, S. Hippler, R. Hofmann, F. Hormuth, K. Houairi, S. Kellner, P. Kervella, R. Klein, J. Kolmeder, W. Laun, P. Léna, R. Lenzen, M. Marteaud, V. Naranjo, U. Neumann, T. Paumard, S. Rabien, J.R. Ramos, J.M. Reess, R.-R. Rohloff, D. Rouan, G. Rousset, B. Ruyet, A. Sevin, M. Thiel, J. Ziegleder and D. Ziegler

We present the AO assisted, near-infrared VLTI instrument GRAVITY for precision narrow-angle astrometry and interferometric phase referenced imaging of faint objects. With its two fibers per telescope beam, its internal wavefront sensors and fringe tracker, and a novel metrology concept, GRAVITY will not only push the sensitivity far beyond what is offered today, but will also advance the astrometric accuracy for UTs to 10 µas. GRAVITY is designed to work with four telescopes, thus providing phase referenced imaging and astrometry for 6 baselines simultaneously. Its unique capabilities and sensitivity will open a new window for the observation of a wide range of objects, and—amongst others—will allow the study of motions within a few times the event horizon size of the Galactic Center black hole.

1 Fundamental Measurements in Astrophysics GRAVITY, an interferometric imager with 10 µas astrometric capability, coupled with spectroscopic and polarization modes and optimized to exploit the exquisite sensitivity of the 4 × 4 VLTI system, will revolutionize dynamical measurements of celestial sources interacting through gravity. It will carry out the ultimate test of determining whether or not the Galactic Centre harbors a 4 × 106 M black hole. It has the potential to directly measure the space-time metric around this black hole, and thus may be able to test General Relativity in the presently unexplored strong curvature limit. GRAVITY will also be able to unambiguously detect and measure the mass of black holes in massive star clusters throughout the Milky Way and in F. Eisenhauer () Max-Planck-Intitut für extraterrestrische Physik, 85748 Garching, Germany e-mail: [email protected] A. Moorwood (ed.), Science with the VLT in the ELT Era, Astrophysics and Space Science Proceedings, doi: 10.1007/978-1-4020-9190-2_61, © Springer Science + Business Media B.V. 2009

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Fig. 1 Key experiments with GRAVITY: with a precision of 10 µas, GRAVITY will see objects move throughout the Milky Way and even in distant galaxies: the horizontal axis denotes the maximum distance from the Earth to which a particular project can be carried out with the capabilities of GRAVITY. The vertical axis represents the time required for such measurements, separated into three rough categories of one season, three and ten years. A colour version of this figure is available at dx.doi.org/10.1007/978-1-4020-9190-2_61

many AGN. It will make unique measurements on gas jets in YSOs and AGN. It will explore binary stars, exoplanet systems and young stellar disks. Because of its superb sensitivity GRAVITY will excel in milliarcsecond phase-referenced imaging of faint objects of any kind. Because of its outstanding astrometric capabilities, it will detect motions throughout the local Universe and perhaps beyond. Because of its spectroscopic and polarimetric capabilities it is capable of detecting gas motions and magnetic field structures on sub-milliarcsecond scales.

2 Instrument Design and Performance The VLTI, with its four 8 m telescopes and a collecting area of 200 m2 , is the only interferometer to allow direct imaging at high sensitivity and image quality. GRAVITY will for the first time utilize the unique 2 field of view of the VLTI, providing simultaneous interferometry of two objects with four telescopes. This is the key for narrow angle astrometry with a precision of 10 µas at large telescopes. Figure 2 shows an overview of the instrument concept. The application of phase referenced imaging—instead of closure phases—is a major advantage in terms of model-independence and fiducial quality of interferometric maps with a sparse array such as the VLTI. The second major new element of GRAVITY is the use of infrared wavefront sensors to open a new window for interferometry. In addition to broad band (K) imaging and astrometry, GRAVITY also features modest resolution spectroscopy and polarization analysis capabilities. The following Table 1 gives an overview of the expected performance of GRAVITY. The baseline for the infrared

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Table 1 Expected performance of GRAVITY Adaptive optics on K = 7 star

36% Strehl

Fringe tracking on K = 10 star

270 nm rms OPD on science channel

Astrometry on K = 10 primary

10 µas in 5 minutes

and K = 15 secondary Interferometric imaging

S/N Visibility = 10

on K = 16 in 100 s Size and position measurements

K ≥ 19 in 6 hours

wavefront sensor is a Shack Hartmann system. It will be located in the VLTI laboratory, thus also correcting for tunnel seeing in the VLTI optical train. The wavefront correction will be applied to the MACAO deformable mirrors located at the UT Coude focus. The interferometric beam combiner is based on fiber-fed integrated optics. The instrument is equipped with polarization-control, differential delay lines, and fast tip/tilt and fringe tracking actuators. GRAVITY will have two beam combiners for two objects. The first is optimized for phase referencing/fringe tracking at a high frame rate. The second, the science beam combiner, is optimized for long integrations. Its spectrometer provides moderate spectral resolution (R ≈ 500), and a Wollaston prism for polarization analysis. GRAVITY will have all its components enclosed in a vacuum cryostat for optimum stability and background suppression. The GRAVITY metrology is optimized for astrometric accuracy. Laser light is backpropagated from the GRAVITY beam combiners up to the telescope secondary mirrors, producing fringe patterns, which carry the differential optical path information. These fringe patterns are observed in scattered light through cameras mounted on the telescopes. Other than classical laser metrologies, the GRAVITY metrology measures the full beam and covers the entire optics train except the primary mirror, therefore reducing all systematic errors from non-common optical paths to a minimum.

3 Outlook GRAVITY can be built and installed at the VLTI in less than 5 years. In its baseline implementation GRAVITY will run on up to four adaptive optics operated UTs with on- and off-axis wavefront sensing—both in the optical with MACAO and in the infrared with the GRAVITY wavefront sensor—and GRAVITY internal fringe tracking within the 2 field of view. When used with ATs, GRAVITY will run with tip/tilt correction only, without high-order adaptive optics. Several upgrade paths are foreseen, for example external fringe tracking on stars outside the 2 VLTI field of view. GRAVTIY would then provide micro-arcsecond astrometry between very faint objects (both too faint to be used for fringe tracking) within the 2 field of view.

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Fig. 2 GRAVITY concept and some of its subsystems: infrared wavefront sensor (bottom left), beam combiner instrument (bottom right), fiber coupler, polarization and differential OPD control (top right), and fringes of the laser metrology on the VLT secondary mirror obtained during the prototype testing (top left). A colour version of this figure is available at dx.doi.org/10.1007/978-1-4020-9190-2_61

GRAVITY

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References 1. F. Eisenhauer et al., in The Power of Optical/IR Interferometry, ed. by A. Richichi, F. Delplancke, A. Chelli, F. Paresce. ESO Astrophysics Symposia (Springer, 2008), p. 431 2. T. Paumard et al., in The Power of Optical/IR Interferometry. ed. by A. Richichi, F. Delplancke, A. Chelli, F. Paresce, ESO Astrophysics Symposia (Springer, 2008), p. 313

Part VII

New Instrument Concepts and VLT/I Operating Modes

Smart Focal Plane Technologies for VLT Instruments C.R. Cunningham and C.J. Evans

1 Introduction As we move towards the era of ELTs, it is timely to think about the future role of the 8-m class telescopes. Under the OPTICON program novel technologies have been developed that are intended for use in multi-object and integral-field spectrographs. To date, these have been targeted at instrument concepts for the European ELT, but there are also significant possibilities for their inclusion in new VLT instruments, ensuring the continued success and productivity of these unique telescopes.

2 ‘Smart Focal Planes’ A smart focal plane is a system that maximises the use of the telescope focal plane for science observations. Examples range from wide-field, multi-object spectrometers and multiple integral field units (IFUs), to future systems such as roboticallyfed, photonic spectrometers. Under the Framework Six OPTICON program1 , a wide range of novel systems has been developed by a team from eight European countries [1], providing a toolkit for instrument designers including: • Replicated image slicers to enable more economic production of multiple IFUs [2]. • ‘Starbugs’—miniature robots that carry fibres or pick-off mirrors to any given place in the focal plane, potentially at cryogenic temperatures [3]. • ‘Starpicker’—an alternative cryogenic, robotic positioner to place pick-off mirrors on a potentially curved, dual focal-plane that is tumbled into its observing position [4]. • Active correction mirrors used to correct for aberrations caused by the movement of a pick-off mirror across the focal plane [5]. • Micro-mirror arrays fabricated in silicon that can be used to form multi-object slitlets with high densities of objects [6]. C.R. Cunningham () Royal Observatory Edinburgh, UK Astronomy Technology Centre, Blackford Hill, Edinburgh, EH9 3HJ, UK e-mail: [email protected] 1 OPTICON has received research funding from the European Community’s Sixth Framework Program under contract number RII3-CT-001566. A. Moorwood (ed.), Science with the VLT in the ELT Era, Astrophysics and Space Science Proceedings, doi: 10.1007/978-1-4020-9190-2_62, © Springer Science + Business Media B.V. 2009

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3 Wide-field Options for the VLTs In the era of ELTs, an obvious role for an 8-m telescope could be for wide-field imaging surveys, in the same manner that CFHT and UKIRT have been very productive in the 8-m era. However, the VLTs were never designed for wide-field operation and there are severe optomechanical and operational constraints on providing a wide-field facility at any of the VLT focal stations. It will be very difficult and expensive to extend the VLT focal plane to beyond one degree [7]. Indeed, widefield imaging in the near-IR is not competitive unless a field of at least one degree can be accessed; at visible wavelengths even this would not be competitive with Pan-STARRS and LSST. A number of presentations at this workshop have made a clear case for spectroscopy of a large number of sources over a field of a degree of more, e.g. unprecedented galaxy redshift surveys to map large-scale structure, or wide-area stellar surveys to trace the mass-assembly history of the Milky Way. Smart focal plane technologies can now deal with the physical size, and possible curvature, of such a field-of-view, even if cooled to cryogenic temperatures. However, the technical and cost problems relating to modification of the telescope itself are severe. The classical option for wide-field astronomy is at prime focus, but the VLTs have significant mass, space and systems limitations. For instance, the secondary mirrors are part of the active telescope system, and carry out real time tip-tilt and focus correction at up to 50 Hz. Any prime focus system would need to provide similar functionality. With the notable exception of Subaru, decisions taken in the early design stages of the 8-m class telescopes have made prime focus options difficult. In principle, Gemini has the option for a replaceable top-end to allow conversion to a wide-field mode, but this has never been implemented; replacing the top end with a carbonfibre structure has been proposed to allow an increased mass at the prime focus [8], but using such a solution on one of the VLTs would be very expensive. It would also be necessary to provide correcting optics—a non-trivial challenge for an 8-m telescope. For instance, the Subaru corrector for the 30 FMOS field (optimised from 0.9 to 1.8 µm) requires three elements of up to 600 mm diameter in BSM51Y glass [9]. Other options, such as the forward-Cassegrain arrangement used to convert UKIRT to a wide-field mode for WFCAM [10], would likely be even more challenging for the VLT.

4 Cladistics & Natural Selection It can be argued that technology selection for astronomical instruments follows a similar process to natural selection in nature. Selection pressures have analogies in terms of resources, competition and even predation—although it is hard to see where sexual selection takes a role! Many ideas follow an evolutionary path with ideas splitting-off from an original concept, through an almost cladistic mapping, as illustrated by the example of wide-field spectroscopy (Fig. 1), where most systems

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Fig. 1 Cladistic map for multi-object spectroscopy. Adapted from [12]. A colour version of this figure is available at dx.doi.org/10.1007/978-1-4020-9190-2_62

can be traced back to Medusa, the first fibre plug-plate system at Steward Observatory [11]. It is noteworthy how evolution has led to significant novelty in an isolated continent like Australia. Similarly, the team at the Anglo-Australian Observatory (AAO) have been responsible for many innovations! A recent example is the Echidna positioner developed for FMOS, that deals with dense packing of targets within a physically small field, in which traditional pick-and-place devices are not feasible [9]. A similar design to the Echidna positioner has been proposed for Gemini/SubaruWFMOS, while an alternative concept uses small field lenses to select sub-fields which are then fed by miniature pick-offs to a fixed grid of fibres. A further concept proposed at this workshop is an update of the traditional fibre plug-plate, using robotics to assemble replaceable fibre modules [13].

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5 Multi-slit Spectroscopy The science case for slit spectroscopy remains strong, particularly in cases where both a high multiplex and good throughput are required for point-like sources with known positions. New technologies are now becoming available to replace slit masks with configurable slit-arrays, that can also operate at cryogenic temperatures. The slitlet mechanism developed for the European Space Agency by CSEM under the OPTICON smart focal planes program is now being incorporated into the Keck MOSFIRE instrument [14]. Moreover, devices such as micro-mirror arrays [6] and the shutter arrays in the JWST NIRSPEC instrument [15] offer potential arrays of up to 10,000 slitlets. These could be used in concepts such as the proposed VLTMegaMOS [16].

6 IFU Spectroscopy In many cases, such as attempts to disentangle observations of merging galaxies, we require three-dimensional imaging spectroscopy. In this instance, image slicers offer significant advantages over slits. Precision manufacturing techniques such as diamond machining, stacked-element glass slicers and electroformed replication could provide a larger number of channels, or larger field-of-view, than the capabilities of KMOS or MUSE. When combined with new OPTICON pick-off technologies (e.g. Starbugs and Starpicker) instruments can be forseen that combine a relatively wide patrol field, with high spatial resolution (i.e. fed by future AO systems).

7 Miniature Spectrometers In the longer-term, new technologies from the photonics industry could result in integrated, miniature spectrometers. For instance, array waveguide devices could be used if the problems of efficient light-coupling and bandwidth can be solved [17]. There are also integrated devices under development for fluorescence spectroscopy that are based on bulk optics and holographic gratings, yielding spectral resolutions of up to 400 at visible wavelengths from a device that can fit within a 1 cm cube [18]. These concepts may be useful in a highly-multiplexed instrument, in which a swarm of miniature spectrometers are self-propelled around the focal plane by autonomous robotic carriers. The ultimate miniature imaging-spectrometer would be an energy sensitive detector. Unfortunately, current devices such as superconducting tunnel junction arrays, or transition edge superconductors, do not offer large pixel-densities, nor useful spectral resolution. However, it is possible that such components could revolutionise spectrometer design in the future, possibly in conjunction with photonic feed devices.

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8 Conclusions In summary, there are considerable technical challenges in developing a wide-field spectrograph for the VLT, particularly when one considers the available resources within ESO and of its partners; a wide-field prime focus instrument would also lead to operational problems for VLTI. Future smart focal plane and photonics technologies may enable a lighter, more compact, prime focus (or forward-Cassegrain) instrument, but current technology readiness levels suggest that a compromise would be to exploit the maximum field available at the Nasmyth foci, e.g. Super-Giraffe [19]. Meanwhile, technology development for ELT instrument concepts such as EAGLE, and through the next Framework Program, should be supported to help break the existing paradigm.

References 1. 2. 3. 4. 5. 6. 7. 8. 9. 10. 11. 12. 13. 14. 15. 16. 17. 18. 19.

C.R. Cunningham et al., Proc. SPIE 5904, 281 (2005) J. Schmoll et al., New Astron. Rev. 50, 263 (2006) R. Haynes et al., Proc. SPIE 6273, 56 (2006) P.R. Hastings et al., Proc. SPIE 6273, 91 (2006) F. Madec et al., Proc. SPIE 6273, 58 (2006) F. Zamkotsian et al., Proc. SPIE 6273, 52 (2006) S.P. Todd, in Science with the VLT in the ELT Era, ed. by Alan Moorwood. Astrophysics and Space Science Proceedings (Springer, Netherlands, 2009), p. 379 Miziarski et al., Proc. SPIE 6273, 95 (2006) P.R. Gillingham et al., Proc. SPIE 4841, 985 (2003) M. Casali et al., Astron. Astrophys. 467, 777 (2007) J.M. Hill et al., Astrophys. J. 242, L69 (1980) G. Smith et al., Proc. SPIE 5495, 35 (2004) I. Parry, in Science with the VLT in the ELT Era, ed. by Alan Moorwood. Astrophysics and Space Science Proceedings (Springer, Netherlands, 2009), p. 417 I.S. McLean, S.M. Adkins, Proc. SPIE 6269, 1 (2006) M.J. Li et al., Proc. SPIE 5650, 9 (2005) O. Le Fèvre et al., in Science with the VLT in the ELT Era, ed. by Alan Moorwood. Astrophysics and Space Science Proceedings (Springer, Netherlands, 2009), p. 163 A. Horton et al., in Science with the VLT in the ELT Era, ed. by Alan Moorwood. Astrophysics and Space Science Proceedings (Springer, Netherlands, 2009), p. 423 I. Avrutsky et al., Appl. Opt. 45, 7811 (2006) M. Lehnert et al., in Science with the VLT in the ELT Era, ed. by Alan Moorwood. Astrophysics and Space Science Proceedings (Springer, Netherlands, 2009), p. 431

Applications of Digital Micromirror Devices to Astronomical Instrumentation M. Robberto

Abstract MEMS devices are among the major technological breakthroughs of the last two decades. Besides finding widespread use in high-tech and consumer market electronics, MEMS enable new types of astronomical instruments. I concentrate on Digital Micromirror Devices, which have been already adopted in astronomy and can enable scientific investigations that would otherwise remain beyond our technical capabilities.

1 MEMS and DMDs Micro-Electro-Mechanical-Systems (MEMS) are systems where mechanical elements are integrated with their control electronics on a common silicon substrate through microfabrication technology. Whereas the control electronics is fabricated using standard integrated circuit processes, the mechanical components are fabricated using “micromachining” processes that selectively etch away parts of the silicon wafer or add new structural layers. Digital Micromirror Devices (DMDs) represent a particular application of MEMS. Invented at Texas Instrument (TI) by L. Hornbeck, a DMD is a chip having on its surface several hundred thousand microscopic mirrors arranged in a rectangular pattern. Each mirror can be rotated ±12◦ to an ON/OFF state and can oscillate between these two positions at several KHz. Uniformly illuminated, the ON/OFF duty cycle of the mirrors generates a gray scale image that can be projected on a screen (DLP projectors). Located at the focal plane of a telescope, a DMD can be used as an optical switch, reflecting the light of selected sources e.g. into a multi-object spectrograph.

2 Instruments Based on DMDs The use of DMDs for astronomical research was first pioneered for the NIRSPEC spectrograph on board JWST [1, 2]. To ensure multi-object capabilities to NIRSPEC, NASA funded both DMD and Micro-Shutter Arrays (MSAs) development. In MSAs the target selection occurs in transmission rather than in reflection [3]. In 2002 MSAs were eventually selected due to their capability of operating at M. Robberto () Space Telescope Science Institute, 3700 S. Martin Dr., Baltimore, MD 21218, USA e-mail: [email protected] A. Moorwood (ed.), Science with the VLT in the ELT Era, Astrophysics and Space Science Proceedings, doi: 10.1007/978-1-4020-9190-2_63, © Springer Science + Business Media B.V. 2009

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cryogenic temperatures and to the optical constrains posed by a diffraction limited 8-m class telescope. JWST MSAs come in much smaller size than DMDs (171 × 365 pixels vs. 2048 × 1080 for the largest TI DMD) and are still prone to cosmetic defects, whereas other figures of merit (e.g. contrast) appear comparable to those achieved by DMDs. IRMOS [4] has been developed as a pathfinder for NIRSPEC/JWST. NIRMOS is a near-IR spectrograph based on a 848 × 600 DMD operated at about −45 °C. It offers multi-object capability (∼ 20 sources) with resolving power R = 300, 1000, 3000 in the Y, J, H, and K band. IRMOS is currently offered at the KPNO 4-m telescope [8] and is producing scientific results, e.g. [5]. RITMOS [6] fully exploits the capabilities of DMDs. In RITMOS the telescope beam impacts the DMD perpendicularly and the light reflected in the ON/OFF states is analyzed by a spectrograph and an imager, respectively. In practice, RITMOS provides multi-object spectroscopy (DMD-ON) and imaging (DMD-OFF) in parallel. RITMOS, which operates at visible wavelenghts, is also based on a 848 × 600 TI DMD and operates at the Mees 24 telescope. The SPACE mission [7], which is now considered by ESA for an European Dark Energy Mission, exploits DMDs. In the proposed configuration, SPACE performs R ∼ 400 spectroscopy of > 500 million galaxies, 6000 at a time, down to AB ∼ 23 between 0.6 and 1.8 µm, producing the largest three-dimensional evolutionary map of the Universe and observing baryonic acoustic oscillation patterns between 5 to 10 billion years ago. Versatility is one of the main advantages of DMDs. A DMD bases instrument can do imaging, multi-object spectroscopy and even integral field spectroscopy over a large (arcmins) field of view exploiting Hadamard transforms, a technique recently validated with IRMOS [8]. It is easy to imagine other applications for DMDs in astronomy. For example, a multi-object OH suppression spectrograph can be built by projecting the spectra on a DMD first. Once the OH lines are automatically identified, they can be removed by switching off the appropriate mirrors, with the reflected spectrum recombined at lower resolution for higher signal to noise.

References 1. 2. 3. 4. 5. 6. 7. 8.

Moseley et al., ASP Conf. 207, 262 (2000) MacKenty, Stiavelli, ASP Conf. 195, 443 (2000) Moseley et al., SPIE 5487, 645 (2004) MacKenty et al., SPIE 4841, 953 (2003) Figer et al., Astrophys. J. 643, 1166 (2006) Meyer et al., SPIE 5492, 200 (2004) Robberto, Cimatti, arXiv:0710.3970v1 [astro-ph] (2007) MacKenty et al., SPIE 6269, 37 (2006)

FORS in the Era of Second Generation VLT Instrumentation Kieran O’Brien

The FORSes remain among the most productive and highest impact instruments at the VLT. This has been in part due to the continuous upgrading of the instruments. However, the arrival of X-shooter in 2008 is the biggest challenge yet for the FORS project, with only one Cassegrain focus available for FORS operations. I will present the current plans for the continued operation of the FORSes, which ensure that as little as possible of the current functionality is lost. In addition, I will look towards the future to identify some possible directions for the FORS project by incorporating new technologies, including the use of novel detectors and adaptive optics facilities.

1 Current Status The strengths of the FORS project were detailed by Appenzeller [1]. They are (1) MOS with high quality and instantly adjustable slits, (2) uniform high image quality from 330–1100 nm, (3) very low stray-light and ghosts, (4) interference filter imaging without field effects, and (5) accurate (spectro-)polarimetry. A number of additional points have been added during the lifetime of FORS, namely (1) a Mask eXchange Unit (MXU) allowing higher multiplex observations (FORS2), (2) a red sensitive, low fringe MIT detector (FORS2), (3) a blue sensitive E2V detector (FORS1), and (4) high throughput Volume Phased Holographic (VPH) grisms, for higher spectral resolution and throughput. The FORSes are the most productive instruments in the VLT instrumentation suite. According to a recent study [2], data collected with FORS1 & 2 have led to 525 and 342 papers respectively, which have been cited more than 23 000 times since their publication. An analysis of the papers used in this study reveals the main uses for FORS1 at present are imaging and low resolution (R ∼ 300) spectroscopy of single objects. Whereas for FORS2, we find that most of the time is being spent performing low-resolution (R ∼ 300–600), multi-object spectroscopy (with the MXU) using the grisms with red central wavelengths. Following the recent arrival of the blue sensitive CCD on FORS1 (April 2007), we expect the 2 instruments to become more specialized based on the wavelength of interest. K. O’Brien () European Southern Observatory, Alonso de Cordova 3107, Santiago 19, Chile e-mail: [email protected] A. Moorwood (ed.), Science with the VLT in the ELT Era, Astrophysics and Space Science Proceedings, doi: 10.1007/978-1-4020-9190-2_64, © Springer Science + Business Media B.V. 2009

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2 What Happens Next? As we enter the era of the 2nd generation instruments it is important to look at the capabilities of the instrument suite as a whole to determine in which direction to take the project. FORS1 is the only instrument with a polarimetric capability (apart from EFOSC2 at a Silla), whilst FORS2 is still the preferred instrument for MOS observations beyond 600 nm, due to the low fringe amplitude of the MIT CCD. For these reasons, the FORS Instrument Operation Team (IOT) have recommended to move the polarization optics to FORS2 from P82 and to retire FORS1. We have also recommended to initially offer the new blue-sensitive E2V CCD in Visitor Mode. In subsequent periods this might change to either a time-share basis or on demand, depending on the level of interest from the community. Looking further into the future, one possibility for the retired FORS1 is that it could become a demonstrator for ELT instrument technologies, which would also be instruments in their own right. Two such possibilities are: • GLAO-FORS: A Ground Layer Adaptive Optics assisted multi-object spectrograph, combined with a very red-sensitive chip. The science goals would be to study galaxy clusters in the range z = 1.4–1.8, where significant evolution is expected to be taking place. This would be possible due to the increased S/N from doubling the encircled energy using GLAO, combined with greater red sensitivity and very low fringing. We anticipate an image quality of ∼ 0.3–0.4 , leading to a resolving power of 4000 over a 4 field of view. • ULTRA-FORS: FORS1 with an electron-multiplying CCD (or EMCCD), as was currently demonstrated with ULTRASPEC [3]. This would effectively reduce the read-out noise to zero and allow photon-counting spectroscopy. Amongst the science goals would be to measure the mass of compact objects in Interacting Binaries, the accurate determination of times of occultations and planetary transits to search for minor bodies and the rapid identification of GRB optical counterparts. Acknowledgements I wish to thank the other members of the FORS IOT for their continued work with the FORS project and for help in filtering some of my more outlandish ideas and to Uta Grothkopf for help in compiling the statistics for FORS.

References 1. I. Appenzeller et al., Messenger 94, 1 (1998) 2. U. Grothkopf et al., Messenger 128, 67 (2007) 3. V. Dhillon et al., Messenger 127, 41 (2007)

Wide Field Options on the VLT Stephen Todd

1 Introduction Over the past decade, as the 8–10 m class telescopes have entered service, many of the existing 2–4 m class telescopes have been adapted for wide field imaging or spectroscopy surveys. Increasing the field of view of one of the unit telescopes might allow the VLT to complement the capabilities of the E-ELT and TMT in a similar way. The size of the field currently available on the VLT is shown in Fig. 1.

2 Wide Field Imaging For wide field imaging we want the point spread function to cover 2 pixels on the detector. For an instrument on the VLT, the entrance pupil will be approximately 8 m in diameter, so the focal ratio needed to achieve a given plate scale is f/# = 7.2 × 10−3 × (plate scale in mm/degree).

(1)

If we assume 0.5 arcsec seeing and 18 micron pixels—typical for a near-infrared detector—then this would require an f/1.8 camera. A slower focal ratio than this will over-sample the PSF, requiring more detectors to cover the same area of sky in a single exposure. In addition to the increased cost and complexity, the larger number of pixels also increases the contribution of the detector read-noise.

S. Todd () Royal Observatory, UK Astronomy Technology Centre, Blackford Hill, Edinburgh, EH9 3HJ, United Kingdom e-mail: [email protected] A. Moorwood (ed.), Science with the VLT in the ELT Era, Astrophysics and Space Science Proceedings, doi: 10.1007/978-1-4020-9190-2_65, © Springer Science + Business Media B.V. 2009

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Field diameter at Cassegrain Unvignetted 2.68 arcmin Vignetting < 2.5% 15 arcmin Field diameter at Nasmyth Unvignetted 7.1 arcmin Vignetting < 2.5% 22 arcmin

Fig. 1 Vignetting at the Cassegrain and Nasmyth foci as a function of radial field position

Mapping speed The main figure of merit for a wide field imaging system for a given wave-band is the mapping speed: M = AΩ,

(2)

where A is the collecting area of the telescope, Ω is the solid angle on the sky observed in a single exposure, and is the efficiency of the system (this includes the optical transmission of the telescope and instrument, detector efficiency and operational overheads). We will assume that will not vary much from one system to another, and consider only the value of AΩ. We can compare the potential mapping speed of the VLT with a number of existing or proposed wide field imaging facilities, as shown in Table 1. Table 1 Comparison of the potential mapping speed of the VLT with a number of other facilities A/m2

Ω/deg2

Vista IR [2]

11.4

0.6

VLT 0.5° field, 50% filling

49.3

0.1

4.8

VLT 1° field, 50% filling

49.3

0.4

19.2

VST [3]

4.5

1

MegaCam (CFHT) [4]

9.59

1.0

9.6

Suprime-Cam (Subaru) [5]

51.6

0.26

13.2

PAN-STARRS [6]

4 × 1.8

7.0

46.0

LSST [7]

33

9.6

319.0

VLT 0.5° field, 100% filling

49.3

0.2

9.6

VLT 1° field, 100% filling

49.3

0.8

38.4

Facility

AΩ

Near-infrared 6.8

Visible 4.5

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In the near-infrared, a 0.5° diameter field on the VLT would have a lower mapping speed than Vista IR. Even increasing the field of view to 1°, which would be extremely challenging, will only achieve a speed increase of a factor of ∼ 3. We can also consider the number of detectors which would be needed. If we assume 2k × 2k detectors with 18 micron pixels (as used in Vista IR), then with a focal ratio of f/1.8 we could cover 50% of a 0.5° field with ∼ 5 detectors, or 50% of a 1° field with ∼ 18 detectors. If the focal ratio is increased to f/3 then these numbers are increased to 13 and 50 respectively. It seems likely that the next generation of available near-IR detectors will probably be 4k × 4k with 15 micron pixels, and allowing a filling factor of up to ∼ 80%. The increased filling factor would allow a 0.5° field instrument on the VLT to be similar in mapping speed to Vista IR. While it could conceivably be possible to exceed the mapping speed of Vista IR it seems unlikely that a convincing case can be made for building such an instrument. In the time taken to build such an instrument Vista IR should have collected many years of data. We can make similar comparisons at visible wavelengths. A mapping speed comparable to, or maybe even slightly greater than MegaCam (CFHT) or Suprime-Cam (Subaru) would be possible with a 0.5° field. It would not be plausible to match the mapping speed of PAN-STARRS or LSST.

2.1 Wide Field Spectroscopy For wide field spectroscopy the requirements on the focal plane are slightly different. Multi-object spectroscopy requires some sort of pick-off system, such as pickoff-arms, or positionable fibres. For any of these systems we need a focal plane of a manageable size. Working with a focal plane larger than ∼ 1 m is likely to be difficult. Some sort of field correction optics would probably be needed before the pickoff system to flatten the field, correct the off-axis aberrations and image the pupil to the correct position, possibly producing a telecentric focal plane. These field correction optics would probably be refractive, so these also limit the maximum diameter of the field to ∼ 1 m. It is already possible to access a field of 22 arcmin diameter at the Nasmyth focus, so to justify major modifications of the telescope we would need to achieve a field approaching 1°, which would have a diameter of 1 m at f/7.2, or 250 mm at f/1.8.

3 Optical Configurations From the discussion above we can see that for wide field imaging we need to aim for a very fast focal ratio—ideally around f/1.8. For spectroscopy focal ratios of around f/1.8–f/7 would be required.

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3.1 Prime Focus The obvious solution would be to place an instrument at prime focus. The VLT primary mirror is f/2—close to our ideal focal ratio either for imaging or spectroscopy. We know from the experience of Subaru, which has a very similar primary mirror, that this is a feasible method of putting a wide field instrument on a large telescope. However, unlike Subaru, the VLT was never designed to accommodate a prime focus instrument. It is likely the construction of such an instrument would be limited by the mechanical strength of the telescope and the space available within the telescope enclosure.

3.2 Moving the Secondary Mirror The focal length of the telescope can be reduced by changing the curvature of the secondary mirror. If the curvature alone is changed then the focal plane is formed above the primary mirror, inside the telescope structure. By also changing the position of the secondary we can move the focal plane. The system shown in Fig. 2(a) produces an f/7 focal plane at the position of the existing Nasmyth focus. To do this, the secondary mirror is moved about 2 m closer to the primary. To produce an unvignetted 1° field the secondary has to be increased to 2 m in diameter—approximately double the size of the existing secondary—and the tertiary mirror would also need to be enlarged.

3.3 Placing an Instrument above the Primary Instead of moving the focal plane to the Nasmyth focus, we can place the instrument, or part of it, above the primary mirror, in a ‘Forward Cassegrain’ position. This configuration was used to convert the 3.8 m UK Infrared Telescope into an imaging survey telescope with WFCAM [1]. The concepts described below appear to be feasible optically. Further work would be required to examine the mechanical engineering challenges of installing such an instrument in the VLT. The layout in Fig. 2(b) shows an f/7 focal plane formed above the primary mirror. The secondary mirror has been moved 200 mm closer to the primary, and is only slightly larger than the existing secondary. The instrument is then positioned in place of the tertiary mirror tower. This could either be an entire instrument, or a field corrector and pickoff system which then feeds light via fibres to a multi-object spectrometer which might be at the existing Cassegrain or Nasmyth positions. Another possibility would be to reimage the focal plane which is formed above the primary mirror to an instrument in a more convenient location. In Fig. 2(c) the f/7 focal plane is reimaged to an f/4 image plane in the position of the existing Cassegrain focus. Installation of optics above the primary mirror is still necessary,

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Fig. 2 A number of optical configurations which could be used to increase the field of view of the VLT, as described in 3.2 and 3.3. A colour version of this figure is available at dx.doi.org/10.1007/978-1-4020-9190-2_65

but this concept avoids the need to support the whole of a large, heavy instrument in this position. The concept closest to the WFCAM design used on UKIRT is to replace the secondary mirror with one producing a much faster focal ratio, and then re-image this focal plane into an instrument located entirely above the primary mirror, as shown in Fig. 2(d).

4 Summary The VLT is not a competitive platform for wide field imaging. On the time scales that we are considering, there will be large amounts of data available from VISTA and PAN-STARRS. Achieving large speed increases over these systems would require a large telescope specifically designed for this purpose, such as LSST. Extending the field of view for wide field spectroscopy may be more attractive. Achieving a significant increase in the field of view would require substantial modifications to one of the unit telescopes, probably removing it from the VLTi in the process. Further work would be required on the science case to quantify the scientific gains, mechanical constraints and costs of such a modification.

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References 1. 2. 3. 4. 5. 6. 7.

D.M. Henry, M.M. Casali, D. Montgomery et al., SPIE 4841, 63–71 (2003) J.P. Emerson, W. Sutherland, SPIE 4836, 35–42 (2002) O. Iwert, D. Baade et al., SPIE 6276, 62760A (2006) O. Boulade, X. Charlot et al., SPIE 4841, 72–81 (2003) S. Miyazaki et al., Publ. Astron. Soc. Jpn. 54, 833 (2002) N. Kaiser, SPIE 5489, 11–22 (2004) K. Gilmore, S. Kahn, M. Nordby et al., SPIE 6269, 62690C (2006)

S. Todd

A Few Degrees Very Wide Field of View Camera for VLT as a Finder for ELT Roberto Ragazzoni, Jacopo Farinato, Emiliano Diolaiti, Giorgia Gentile, Carmelo Arcidiacono, Renato Falomo and Emanuele Giallongo

The quest for wide field imaging, with some added multi-object spectroscopic capability, immediately evokes the Prime Focus option. On the existing 8 m class telescopes the Subaru [1] and the LBT [2] owns these kind of instruments, while for Gemini some has been planned although none has been actually built. About VLT, because of its interferometric vocation, little or no efforts have been made in order to provide such an option from the baseline. On the other hand it is well known that, given a certain optical quality, the field of view and the size of the largest optical element of the corrector grows together in an almost linear fashion. Even with options different from the Prime Focus one it is interesting to point out how such a sort of “rule of thumb” still hold with a certain degree of precision. Both the LBT and the Subaru Prime Focus requires a largest optical element (the front lens of the Prime Focus cage) of the order of 0.6–0.8 m, and offer a field of view slightly smaller than half a degree in diameter. The LSST [3], that is a three mirror telescope with a sort of optical corrector, indeed very similar to a Prime Focus one, does not escape from such a relationship. Its first element after the primary mirror, namely the secondary one, in fact, with a diameter of 3.4 m well match the Field of View of 2.5 degrees [4] that such a facility will encompass. Also, the trapped Cassegrain wide field imager foreseen in the past for LBT, do not escape from such law too. Making a Prime Focus is, definitively, a risky game, and the overall weight of the Prime Focus cage is, roughly speaking, about one order of magnitude larger than the weight of the largest lens itself. It is definitively illusory to conceive a conventional Prime Focus imager for VLT with a Field of View significantly larger than half a degree with a frontal lens lighter than several hundred of Kgs and with an overall weight of the order of a few metric tons. Such a facility will have to be positioned away from the current secondary mirror and, even neglecting the constraints imposed by the envelope of the building to avoid interference, this will surely poses some challenges, not necessarily impossible to overcome, to the telescope structure. Moreover it would be very demanding R. Ragazzoni () INAF – Osservatorio Astronomico di Padova, Vicolo dell’Osservatorio 5, 35122 Padova, Italy e-mail: [email protected] A. Moorwood (ed.), Science with the VLT in the ELT Era, Astrophysics and Space Science Proceedings, doi: 10.1007/978-1-4020-9190-2_66, © Springer Science + Business Media B.V. 2009

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Fig. 1 A conceptual design of the camera idea: a correctors array posed on the VLT focal plane. A colour version of this figure is available at dx.doi.org/10.1007/978-1-4020-9190-2_66

Fig. 2 A collimator composed by 6 elements is posed on the VLT F/4 focal plane. In correspondence of the pupil image there is a corrector plate that convolve light on a CCD. Because of radial distribution of aberration it is possible to split the whole field of view in micro-region within the correction remains roughly the same in a range fixed by the image quality we want to achieve

to imagine this as an option preserving other functionalities of the telescope itself, in particular interferometry, unless with major mechanical workload in the daylight for turning the telescope from a Prime Focus option to the conventional one. Finally, this can maybe achieve a Field of View of the order of one degree, and still this would require a Prime Focus with a front lens rivaling the largest refractive elements ever made in the world (still the famous 1.5 m lens of the Paris Exhibition of 1900 and then lost keep the record, leaving some hopes that in a little more than a century something larger could be one day built on the purpose). We recently introduced a completely different concept for a very wide field imager. Here we refer, pushing a little to the edge of the feasibility, to a full field of view of 3 degree in diameter. This represent about 36 times the field covered by existing Prime Focus and would surely be in direct competition with LSST. The concept rely on a large number of small adiacent cameras mounted on a low focal ratio Cassegrain focus, namely an about F/4 one (Fig. 1 and Fig. 2).

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Fig. 3 In figure it can be seen the upper part of one of the VLT Unit Telescope. At the top center one can see the secondary mirror as it is now, while just below one can see a design in scale of the secondary mirror needed to make the UT an F/4 Cassegrain. A colour version of this figure is available at dx.doi.org/10.1007/978-1-4020-9190-2_66

This of course requires the change of the secondary mirror with a larger one (Fig. 3), but positioned closer to the primary mirror, leaving no problem for any possible interference with the existing VLT building and ameliorating the mass and momentum stress to the telescope structure. Also, the existence nevertheless of a Cassegrain foci, although with a different focal ratio and in a position that can be placed with some degree of freedom in some preferred place, shed some hope that, with the proper retrofitting of the VLTI arm related to the Unit Telescope chosen for the conversion to Wide Field imager and spectrograph, the interferometric option can be fully retained. This should be the subject of a specific study, but the tiny Field of View of any interferometric channel makes such a chance a realistic one. There are different ways this camera can have spectroscopic capability. The one envisaged in the original concept involves a certain number of fiber positioner to be placed in alternative to the detector. Note that these will be in the, mechanically comfortable, condition to requires to move some fibers in an area smaller than the one the mechanical device can occupy. From preliminary discussions on such an option for LBT we found that off-the-shelf components could be, in principle, used for such an option. Alternatively a further intermediate focal plane and pupil plane could be introduced to allow for some multi-slit and grism directly in the camera design. Depending upon the required quality of the intermediate focal plane the golden rule of optical design (requiring more elements for more “tasks” to be achieved from the optical device) could be to some extent dampened.

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Surely such an instrument poses a challenge in term of data reduction and data handling. While the second is common to any instrument like the LSST one (who is going to take advantage of Google technology) the first would be rather particular in this camera. Individual distortion and mapping on the sky of the several cameras would be taken into account, with the remarkable problem that a small overlap in the border between adjacent cameras is to be considered, given the (poor) optical quality of the image at the pure Cassegrain F/4 focal plane. On the other hand this will gives a virtually no dead zones in the covered Field of View. In the spectroscopic option, also operations would require some challenging and demanding approach as the number of slits to be positioned simultaneously to make the instrument effective would becomes huge. On the other hand the instrument is inherently suitable to massive parallelization and one can imagine a dedicated computing unit for a certain small amount of contiguous imaging elements. As mass production of both optical elements, detectors, and electronics are to be envisaged for this instrument, it is difficult to trace a detailed cost profile using existing instrumentation as a baseline, but the cost will probably be somewhere in the 10 to 20 million Euros range. Finally, keeping in mind how boring has been the second Prime Focus of LBT, I can imagine a change in the attitude has to happen to handle, test, qualify and adjust the hundreds of camera that one requires for such a wide field camera. A larger involvement of the industry and/or automated procedures for aligning and tuning the optics depending upon the off-axis position of the single cameras will be likely to be required.

References 1. 2. 3. 4.

Y. Komiyama et al., Proc. SPIE 5492, 525–532 (2004) R. Ragazzoni et al., Proc. SPIE 5492, 507–512 (2004) J.R.P. Angel et al., Astron. Astrophys. Suppl. Ser. 33, 1462 (2001) L.G. Seppala et al., Proc. SPIE 4836, 111–118 (2002)

Science with a 16 m VLT: The Case for Variability of Fundamental Constants Paolo Molaro

Abstract Only astronomical observations can effectively probe in space-time the variability of the physical dimensionless constants such as the fine structure constant α and proton-to-electron mass ratio, μ, which are related to fundamental forces of nature. Several theories beyond the Standard Model (SM) allow fundamental constants to vary, but they cannot make quantitative predictions so that only laboratory experiments and astronomical observations can show if this is the case or set the allowed bounds. At the moment of writing there are claims for a variability of both α and μ at 5 and 4 σ of C.L., respectively, although for α they are contrasted by null results. The observations are challenging and a new spectrograph such as ESPRESSO at the combined incoherent focus of 4 VLT units (a potential 16 m equivalent telescope) will allow for a significant improvement in the precision measurement clearing up the controversy. If the variations will be confirmed, the implications are far reaching, revealing new physics beyond the SM and pointing a direction for GUTs theories. A most exciting possibility is that a variation of α is induced by quintessence through its coupling with the electromagnetic field. If this is the case an accurate measurement of the variability could provide a way for reconstructing the equation of state of Dark Energy [P. Avelino, C.J.A.P. Martins, N.J. Nunes, K.A. Olive in Phys. Rev. D 74, 083508 (2006)].

1 Introduction The Standard Model (SM) of particle physics needs 26 dimensionless physical constants for the description of the natural world [19], of these few are directly related to the strength of fundamental forces. Among them the fine structure constant (α = e2 /(c)) and the proton-to-electron mass ratio, (μ = mp /me ) are of particular interest for us since they can be measured accurately by astronomical observations of intervening absorption systems towards distant QSOs. The fine structure constant α is related to the strength of the electromagnetic force; me is related to the vacuum expectation value of the Higgs field, namely the scale of the weak nuclear force, and mp is related to the ΛQCD or the strong nuclear force, therefore μ is related to the ratio between the strong and weak nuclear forces. P. Molaro () INAF-OAT, Trieste Via G.B. Tiepolo 11, Trieste 34143, Italy e-mail: [email protected] A. Moorwood (ed.), Science with the VLT in the ELT Era, Astrophysics and Space Science Proceedings, doi: 10.1007/978-1-4020-9190-2_67, © Springer Science + Business Media B.V. 2009

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A whatever small variability of these constants will produce a violation of the Weak Equivalence Principle (WEP) and would have far reaching implications revealing new physics beyond the Standard Model. Laboratory measurements with cooled atomic clocks failed to detect variations at the 17th decimal place [17]. Astronomy is providing some evidence for both α and μ variations, although the evidence for α has been contrasted by other groups. Several space-based missions as ACES, μSCOPE, STEP will soon improve existing laboratory bounds for WEP up to 6 orders of magnitude, and they should find violations if present claims of variability are correct under simple linear extrapolation. It is thus desirable that the astronomical community will be able to clear up the case before these accurate experiments will fly, but only astronomical observations can probe WEP non-locally.

1.1 Why Constants Should Vary? Strings and multidimensional theories predict variable constants since the constants are defined in the whole multidimensional space and vary as extra dimensions are varying. The coupling between a scalar field with the electromagnetic field gives also varying constants. The required cosmological constant value is so small that a quintessence is a likely candidate for Dark Energy. Thus varying constants could provide insights into the nature of dark energy and provide evidence for scalar fields [4, 10]. A precise detection of the variability of a constant could be used for the reconstruction of the quintessence potential and of the equation of state of Dark Energy [1, 14]. If one constant is varying, then all the gauge and Yukawa couplings are also expected to vary. There is a relation between the variation of α and μ, but it depends on the context the unification is realized in. Thus, simultaneous measurements of the variability of α and μ at similar redshift will be a key discriminant of the several GUTs models. Theoretical preferences are for a relative change between the μ and α variations of ≤ 50, but larger values are also possible, implying that the strongcoupling constant is running faster than α and therefore δμ should be found to be larger then δα.

2 The Observations Observations of the Werner and Lyman series of the molecular hydrogen in Damped Lyα galaxies (DLA) can be used to bound μ variations. The electron-vibrorotational transitions have different dependence from the reduced mass and can be used to constrain a variability of μ. UVES observations of the DLA at zabs = 3.0 towards QSO 0347-383 [5], but see [6], and of the DLA towards QSO 0405-443 have provided δμ = (24 ± 6) ppm, when the two systems are combined together [16]. The handful of systems investigated for this purpose reflects the difficulties of the measurement. There are few DLA showing H2 and the restframe H2 lines are

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at ≈ 1000 Å, falling in the Lyman forest and requiring a zabs ≥ 2 to be redshifted into the optical window. H2 systems are extensively searched at the moment so that probably new observations will be available in the near future to verify these first findings. Fine structure variability can be probed in the early universe through the primordial nucleosynthesis or through the CMB power spectrum but at the level of a few percent. The most effective way has been achieved through the analysis of metal lines of intervening absorption systems observed in the spectra of distant QSOs. The energy levels of high mass nucleus are subject to relativistic corrections which are sensitive to the mass number. These have been calculated for the most frequently observed resonance lines and constitute the popular Many-Multiplet method. Murphy and collaborators [11] by comparing the redshift of several lines in a sample of 143 systems in the redshift interval 0.2 < zabs < 4.2 found evidence for Δα/α = (−5.7 ± 1.1) ppm. However, this evidence has been contrasted by two other groups which did not find evidence for variability at the level claimed. Chand et al. found an average value of (−0.6 ± 0.6) ppm in a sample of 23 systems, while Levshakov and collaborators found (−0.12 ± 1.79) ppm and (5.66 ± 2.67) ppm in two systems at z = 1.15 and 1.84, respectively, and by using lines of Fe II only [7, 15]. What is the best methodology is currently under debate [9, 12, 13, 18].

2.1 Would You Like an ESPRESSO? These observations are challenging the instrumental performances of UVES-VLT or HIRES-Keck telescopes. Measuring the variability of μ or α implies the measurement of a tiny variation of the position of one or few lines with respect to other reference lines. It is not much different than revealing exoplanets, but with the limitations that only few lines can be used and QSO are much fainter than stellar sources. The precision in the measure of a line position increases with the spectrograph resolving power till the intrinsic broadening of the metal lines is resolved, the signal-to-noise and with the decreasing of the pixel size (Δλ3/2 , see [2] for a precise relation). The ESPRESSO spectrograph described by L. Pasquini at this conference, both in the Super-HARPS or Super-UVES modes, holds the promise for one order of magnitude improvement compared to what presently achieved. Accuracies of few 10 m s−1 are reachable for single lines with relatively short exposures even for faint sources. An error of 30 m s−1 corresponds to an error of 1 ppm for α; such an accuracy will be enough to resolve the present controversy and establish in a definitive way whether α or μ are varying as claimed. A follow up for the ELT has been discussed in [8].

3 Constants and Dark Energy Avelino and collaborators [1] have shown that the measurement of the behavior of variations in α and μ with redshift can be used to infer the evolution of the scalar

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Fig. 1 MC data set based on the scalar potential given in the text producing a α/α = −5 ppm at z = 3. Error bars are of 1 ppm for α and μ as expected with ESPRESSO@4VLT. Reconstruction of the equation of state and its error band. Dashed line represents the assumed DE and the solid line the reconstruction’s best fit. Shaded regions show the 1 and 2 CL of the reconstruction. A colour version of this figure is available at dx.doi.org/10.1007/978-1-4020-9190-2_67

field and of the equation of state of the DE, not very differently from the reconstruction of the potential from the motion of a particle. N. Nunes kindly adapted their detailed analysis to a realistic set of observations which can be performed with ESPRESSO@4VLT. It is assumed that it has been possible to measure α and μ for a sample of 200 and 50 systems respectively and with an equal, for simplicity, accuracy of 0.5 ppm. In the example case the scalar potential is taken as V (φ) = V0 (exp(10kφ + exp(0.1kφ))), which is one of the simplest possible potential accounting for the accelerated expansion. Figure 1 shows the Monte Carlo redshift distribution of the data with this scalar potential assuming that the variation of α is −5 ppm at z = 3 and that the two constants are mutually linked by a fix ratio of −6, as it is suggested by some of the observations. In Fig. 1, right panel, the red dotted line shows the assumed behavior of the w(z) while the black continuous line shows its recovering through a fitting of the simulated data points with a polynomial of order m = 3 (cf. [1] for details). The shaded regions show the 1 and 2 CL of the reconstruction, when both α and μ measurements have been considered. We emphasize that only few observations would clearly show if w(z) is an evolving function of z.

4 Conclusions Variability of physical constants is an important issue for physics and only astronomy can probe this possibility for α and μ in the full space-time. Present observations provide hints of variation for both constants but those for α have been contrasted by other investigations. The ESPRESSO spectrograph presently conceived for the incoherent combined focus of the 4VLT would improve present accuracy by a significant factor and therefore clarify the case. A confirmation of the variability

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would have far reaching implications revealing new physics beyond the SM, showing the right path for GUTs and possibly providing insights into the nature of Dark Energy. If no variability is found, then the new more stringent bounds will be usefully combined with local space experiments for WEP violation. Overall, this seems to be a great opportunity for the astronomical community and I hope that ESO will take advantage of it by considering the construction of the new high precision spectrograph at the incoherent combined focus of the 4 VLT units, a ≈ 16 m equivalent telescope. Acknowledgements It is a pleasure to thank N. Nunes for his simulations for ESPRESSO, all the ESPRESSO collaboration, and in particular S. Levshakov and M. Murphy.

References 1. 2. 3. 4. 5. 6. 7. 8. 9. 10. 11. 12. 13. 14. 15. 16. 17. 18. 19.

P. Avelino, C.J.A.P. Martins, N.J. Nunes, K.A. Olive, Phys. Rev. D 74, 083508 (2006) R.C. Bohlin, E.B. Jenkins, L. Spitzer Jr. et al., Astrophys. J. S 51, 277 (1983) H. Chand, R. Srianand, P. Petitjean, B. Aracil, Astron. Astrophys. 417, 853 (2004) Y. Fujii, arXiv:0709.2211 [astro-ph] (2007) A. Ivanchik, P. Petitjean, Varshalovich et al., Astron. Astrophys. 440, 45 (2005) S.A. Levshakov, M. Dessauges-Zavadsky, S. D’Odorico, P. Molaro, Mon. Not. R. Astron. Soc. 333, 373 (2002) S.A. Levshakov, P. Molaro, S. Lopez et al., Astron. Astrophys. 466, 1077 (2007) P. Molaro, M. Murphy, S.A. Levshakov, astro-ph/0601264 (2006) P. Molaro, D. Reimers, I.I. Agafonova, S.A. Levshakov, in ACFC, ed. by S.G. Karshenboim, E. Peik, 0712.4380 [astro-ph] (2007) C.J. Martins, A.P., astro-ph/0610665 (2007) M.T. Murphy, V.V. Flambaum, J.K. Webb et al., Lecture Notes Phys. 648, 131 (2004) M.T. Murphy, J.K. Webb, V.V. Flambaum, 0708.3677 [astro-ph] (2007) M.T. Murphy, J.K. Webb, V.V. Flambaum, astro-ph/0612407 (2007) D. Parkinson, B. Bassett, J.D. Barrow, Phys. Lett. B 578, 235 (2003) R. Quast, D. Reimers, S.A. Levshakov, Astron. Astrophys. 415, L7 (2004) E. Reinhold, R. Buning et al., Phys. Rev. Lett. 96, 151101 (2006) T. Rosenband et al., Sciencexpress (March 2008) R. Srianand, H. Chand, P. Petitjean, B. Aracil, arXiv:0711.1742 (2007) M. Tegmark, A. Aguirre, M. Rees et al., Phys. Rev. D 73, 023505 (2006)

ESPRESSO: A High Resolution Spectrograph for the Combined Coudé Focus of the VLT Luca Pasquini, A. Manescau, G. Avila, B. Delabre, H. Dekker, J. Liske, S. D’Odorico, F. Pepe, M. Dessauges, C. Lovis, D. Megevand, D. Queloz, S. Udry, S. Cristiani, P. Bonifacio, P. Dimarcantonio, V. D’Odorico, P. Molaro, E. Vanzella, M. Viel, M. Haehnelt, B. Carswell, M. Murphy, R. Garcia-Lopez, J.M. Herreros, J. Perez, M.R. Zapatero, R. Rebolo, G. Israelian, E. Martin, F. Zerbi, P. Spanò, S. Levshakov, N. Santos and S. Zucker

1 The Drivers In the frame of the call for proposal for the E-ELT instrumentation studies, the CODEX team carried out the feasibility study for a spectrograph for extremely stable Doppler measurements. The CODEX project and its scientific goals are described in [6]. During the development of this study, the CODEX team recognized that a CODEX-like instrument would be of high scientific interest also on the VLT. The ESPRESSO concept was born. The contribution by J. Liske in this volume highlights the direct links between the two instruments (see also [4]). The ESPRESSO concept evolves from previous positives experiences at ESO, combining the stability of HARPS [3] with the efficiency of UVES and FEROS [1, 2]. In summary, it is a high-efficiency, high-resolution, fiber-fed spectrograph of high mechanical and thermal stability using, if necessary, the simultaneous reference technique. The first purpose of ESPRESSO is to be a competitive, innovative high-resolution spectrograph to fully exploit the VLT potentiality and to allow new science. ESPRESSO has indeed many very interesting applications, and several have been addressed in this conference by different speakers. The quest for enhanced radial velocity capabilities at the VLT for exo-planet search has been emphasized by Renzini, Bono, Queloz and Udry. P. Molaro discussed the relevance of investigating the variability of physical constants, and V. D’ Odorico the results which can be obtained by studying the chemistry of the Intergalactic Medium. Finally, the detailed chemical analysis of stars will greatly benefit from ESPRESSO, as highlighted, for instance in the presentations by Bonifacio. Additional interesting topics are widely discussed in the proceedings of the “Precision Spectroscopy in Astrophysics” conference [7]. L. Pasquini () ESO, Garching, Germany e-mail: [email protected] A. Moorwood (ed.), Science with the VLT in the ELT Era, Astrophysics and Space Science Proceedings, doi: 10.1007/978-1-4020-9190-2_68, © Springer Science + Business Media B.V. 2009

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The second purpose of ESPRESSO is to gather fundamental experience for CODEX. We finally find extremely exciting the possibility of using a 16 m equivalent telescope, in advance the E-ELT will be fully available to the community.

2 From Requirements to Design The first challenge is to obtain the highest stability, while preserving an excellent efficiency. High spectrograph optomechanical stability is obtained through a controlled environment in vacuum and avoiding movable components. One critical item is the light input system, which must scramble the signal to ensure that the variability at the fiber input does not degrade the stability of the spectrograph, still keeping an excellent transmission. The requirement that ESPRESSO is kept in vacuum and thermally controlled implies containing the instrument volume and the optics size. Obtaining the results from an integrated, system perspective: in order to obtain the demanding ESPRESSO performances, the whole chain must work, from object acquisition, to the data reduction and analysis. While a definitive answer will eventually come only from the use of the instrument, two main system tools have been adopted: first tool is the capitalization of the HARPS experience and its extension to the ESPRESSO requirements. Second tool is the extensive use of simulations generated to quantify the calibration requirements, the main requirements, the subsystem requirements. A number of critical items have been identified, and they are addressed with dedicated R&D efforts: 1. Efficiency: Improvements in efficiency with respect to HARPS include a shorter fiber length, a more efficient scrambling system and a two arm spectrograph design, with the use of two VPHs as crossdispersers. 2. Scrambling: This aspect differs between the 1-UT case and the 4-UTs case. In the 1-UT case the problem is analogue to the one of HARPS, but with 50 times more stringent requirements. In the 4-UT case there are 4 independent pupils; the scrambling shall therefore happen after combining the light from the sub-pupils, and to this scope 3 different systems are tested. 3. CCD temperature control: HARPS shows a correlation between detector temperature and radial velocity variations. A copy of the HARPS cryostat is studied, and design changes are made to improve this aspect. The construction of a prototype is planned. 4. Calibration System: A novel calibration system based on a laser frequency comb has been proposed and its feasibility is studied through a contract with MPQ. The final prototype is expected in 3 years from now (cf. Manescau et al., these proceedings). 5. Slanted VPHG: The slanted fringes VPHG is a potential innovative feature since, to the best of our knowledge, such a VPHG has never used in combination with a spectrograph, as crossdisperser and beam compressor. The use of slanted VPHs

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is not mandatory for ESPRESSO, but seems unavoidable for CODEX. We have therefore opted to proceed to the prototyping.

3 The Design Following HARPS, ESPRESSO is designed with a dual fiber system, where the second fiber can be used either to record the sky or to monitor the spectrograph shifts by recording a simultaneous calibration source. It is actually one of the ESPRESSO aims to gather sufficient information for deciding if a simultaneous calibration fiber is required in CODEX, or if an operations scheme with interloped calibrations is acceptable for that instrument. ESPRESSO will sit in the Combined Coudé room, fed by a Coudé train which brings the light from the B Nasmyth platform to the room, where it is collected by the instrument acquisition and guide system into a fiber, which feeds the spectrograph. The spectrograph itself is contained in a vacuum tank enclosed in a thermally controlled room. No movable nor motorized functions are present inside the tank.

3.1 The Coudé Train (see Avila et al.) Even if in the VLT provision for the Coudé optics is made, the design and hence the components of the train require to be developed. The ducts distance from each telescope to the combined focus is different for each UT; even if the concept for the Coudé train is the same for all UTs, each one will be different. The Coudé train has a FoV of 5 arcsec radius and it will be coated for an excellent response in the 300– 750 nm wavelength range. Given the large distance traveled by the light in the duct, some induced seeing is expected at the fiber entrance, which will be compensated by a stabilization system in the fiber head . This component is critical, because the HARPS experience has shown as, even with a good scrambler, a movement of the source of 0.5 arcsec may induce a shift of up to a few m/sec.

3.2 Calibration Unit (see Manescau et al.) The relevance of a novel, precise, predictable stable calibration system cannot be over-stressed and the characteristics of an ideal calibration source are described, for instance in [5]. The wavelength calibration unit for the spectrograph is based on Laser Frequency Comb system. Provision for the use of Th-Ar lamps and other lamps for flat field will be made.

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Fig. 1 The optical layout of ESPRESSO. The light, after injection form the fibers, passes through a pupil anamorphoser and the pupil is split. The two half-pupils are projected onto the echelle. The dichroic separates the blue and red arms, which are crossdispersed by VPH gratings. A colour version of this figure is available at dx.doi.org/10.1007/978-1-4020-9190-2_68

3.3 Spectrograph Optics (see Spanò et al.) The basic concept for the spectrograph is a cross dispersed echelle with two arms. The echelle grating size is 1700 × 200 mm, consisting of a 4 × 1 mosaic of 408 × 200 mm grating segments, or two UVES gratings. To limit the size of the echelle, pupil slicing is applied. The grating operates in near Littrow configuration. The optical design makes use of anamorphism and pupil slicing. In this way a compact design is obtained, and a 20 cm optical beam and a 20 × 160 cm echelle provides the resolving power of a un-sliced 40 cm beam spectrograph. The spectrograph optical design is given in Fig. 1.

3.4 Mechanics & Vacuum Vessel (see Zapatero and Osorio et al.) The spectrograph mechanics shall maintain the correct alignment and the configuration has been chosen to allow highest stability providing easy access for installation and maintenance. The whole spectrograph, including the detector head, is mounted on an optical bench and installed within a evacuated vacuum vessel with its temperature precisely controlled (few mK). The vacuum tank will be hosted in a thermally controlled room in the Coudé lab.

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Table 1 Characteristics of the ESPRESSO design Characteristic

Standard 1 UT

Faint object 4 UT

High efficiency 4 UT

Wavelength range

350–780 nm

350–780 nm

350–780 nm

Resolving power

160 000

40 000

80 000

Sampling (average)

4 pixels

16 pixels

8 pixels

Spatial pixels

24

24

48

Simultaneous calibration

YES

NO

YES

Sky subtraction

YES (either SimCal)

YES

YES (either SimCal)

3.5 Instrument Control and Software (see Megevand et al.) The instrument control hardware consist in a number of LCUs and the different controllers to control the instrument functions. The Data Reduction Software will transform raw frames into clean, extracted, flat fielded, wavelength calibrated spectra. The main difference between the ESPRESSO DRS and the standard ESO pipelines is that it shall deliver the best science quality data on-line. Data Analysis SW is that part of the science SW which is non-common to the various scientific domains and requires specific tools and focus. The Data-Analysis SW is preferably automatic. For all tools the input is an extracted, wavelength-calibrated spectrum delivered by the DRS, and its output are scientific observables. We distinguish following domains for which different data-analysis tools/functions are required: • • • •

Analysis of non-stellar spectra, low SNR science. Correlation, RVs, bi-sector analysis. General (stellar) spectroscopy. Package of other RV extraction methods.

4 Performances ESPRESSO is proposed to have three operating modes, one with 1-UT and two with 4-UTs. The summary of the characteristics is given in Table 1.

References 1. 2. 3. 4. 5. 6. 7.

H. Dekker, S. D’Odorico et al., SPIE 4008, 534 (2000) A. Kaufer, L. Pasquini, SPIE 3355, 844 (1998) M. Mayor, F. Pepe et al., Messenger 114, 20 (2003) J. Liske, A. Grazian, E. Vanzella et al., 2008MNRAS. tmp..460L, in press M. Murphy, T. Udem et al., Mon. Not. R. Astron. Soc. 379, 1407 (2007) L. Pasquini, S. Cristiani et al., Messenger 122, 10 (2005) N. Santos, L. Pasquini, A. Correia, M. Romaniello, Precision Spectroscopy in Astrophysics (Springer, Berlin, 2007)

Feeding Optics for the ESPRESSO Spectrograph G. Avila, P. Dimarcantonio and F. Zerbi

1 Introduction and Optical Design ESPRESSO will be located in the Coudé Combined Laboratory (CCL). The instrument will be fed by one or four UT Telescope beams through the Coudé B path of the UT units. The CCL is communicated with 4 underground ducts arranged in radial distribution. The Coudé Train is composed by prisms and a rely lens. The VLTI Coudé Train is made only with mirrors (Fig. 1). In the Coudé room, an off-axis parabola acts as a collimator and since the distances between the UT’s and the CCL are slightly different (between 60 and 70 m) a ‘Field Mirror-Prism’ in the Coudé Room projects the pupil to a fix position in the CCL. Finally the telescope beams will be launched into optical fibres in the CCL by means of an F/2.3 objective. This optics group includes the ADC and the ‘Stabilization Opto-Mechanical System’ (Fig. 1). The advantage to use prisms to bend beams instead of mirrors is efficiency. The beams arrive perpendicular to the faces of the prisms, so high efficiency AR coatings may be used. With mirrors the incidence angles are very high (70 degrees), so the reflectivity drops substantially. On the other hand, an optical train made with aspherical mirrors like the ones in the VLTI is much more expensive that the proposed prism solution. However, the prisms require high degree of glass homogeneity.

2 Observing Modes and Image Stabilization In the CCL the individual telescope beams will feed optical fibres which will be arranged in the slit of ESPRESSO spectrograph. A Fibre Selector (Fig. 2) will be used to set a fibre configuration according to the required observing modes. The Fibre Selector will allow: observations with any of the VLT UT units (High Resolution mode, using a single fibre) and simultaneous observation with the 4 UT units. In the latter case, a light pipe (a glass plate of 100 by 500 microns and few centimeter long) may be used to scramble the light of the four fibres coming from the telescopes. The fibres can be arranged on a line on the light pipe or in a square G. Avila () ESO, Karl-Schwarzschild-Str. 2, 85748 Garching, Germany e-mail: [email protected] A. Moorwood (ed.), Science with the VLT in the ELT Era, Astrophysics and Space Science Proceedings, doi: 10.1007/978-1-4020-9190-2_69, © Springer Science + Business Media B.V. 2009

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Fig. 1 Optical layout for the visible Coudé path and optical elements of the visible Coudé train. A colour version of this figure is available at dx.doi.org/10.1007/978-1-4020-9190-2_69

Fig. 2 Fibre selector and beam stabilizator. A colour version of this figure is available at dx.doi.org/10.1007/978-1-4020-9190-2_69

configuration to obtain the following additional modes: (a) low resolution mode but high RV accuracy (fibres arranged along the dispersion) and (b) medium resolution mode and high RV accuracy (fibres compacted in a square distribution).

New Design Approach for a Very-High Resolution Spectrograph for the VLT Combined Focus Paolo Spanò and Hans Dekker

To achieve very-high spectral resolution (R > 100 000) with large telescopes (D > 8 m) new solutions have been investigated, like anamorphic pupil slicing, elliptical pupils, slanted VPH gratings, starting from the initial design of CODEX, as proposed by B. Delabre. Generally, for a given slit width, the spectral resolution of a spectrograph is proportional to the length of the grating and inversely proportional to the telescope diameter. For ESPRESSO, as in CODEX, a trick has been used: i.e. to slice the pupil, allowing to reduce the grating length by a factor two. The price to pay is a detector area that is larger by the same amount. An anamorphic cross-disperser is used to create a fast beam in the cross-dispersion direction and so reduce the required detector area.

1 The Optical Design Light from fibers will feed the spectrograph entrance slit that can accept F/2.5 beams, giving a spectral resolution in excess of 150 000 when fed by one VLT unit telescope. The optical layout (see Fig. 1) has evolved starting from initial CODEX design, and includes pupil slicing, elliptical pupil onto the echelle, anamorphic cross-disperser to recover a circular pupil entering the camera and reduce effective slit height. Light from fibers is processed by the anamorphic pupil slicer unit, near the center of the instrument, creating two images of the fiber core and a 40 × 20 cm rectangular pupil projected onto the echelle. A three-mirrors anastigmat collimator, used in double pass, collimates the light on a linear mosaic of four R4 echelles, each 220 × 420 mm in size and having a total length of 1.7 m, where it is dispersed. A transfer collimator creates a collimated beam, where a dichroic is placed, splitting light into two arms, optimized for different spectral ranges. VPH grisms act as cross-dispersers. Two cameras focus light onto 8K × 8K detectors. P. Spanò () Oss. Astr. Brera, INAF, Via E. Bianchi 46, 23805 Merate, Italy e-mail: [email protected] A. Moorwood (ed.), Science with the VLT in the ELT Era, Astrophysics and Space Science Proceedings, doi: 10.1007/978-1-4020-9190-2_70, © Springer Science + Business Media B.V. 2009

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Fig. 1 ESPRESSO overall optical layout and that one of the anamorphic pupil slicer unit (inset). A colour version of this figure is available at dx.doi.org/10.1007/978-1-4020-9190-2_70

2 The Anamorphic Pupil Slicer Unit This unit makes two different functions: it enlarges a circular pupil into an elliptical one (ratio 4 : 1) and then, it slices this elongated pupil into two sub-pupils, overlapping them onto the grating. The first function is operated by two cylindrical parabolic mirrors, acting along perpendicular directions. Such mirrors have no aberration on-axis, giving the highest image quality and efficiency. Then an off-axis parabolic slicer mirror creates two different images near the field slicer mirror. This all-mirror device has no chromatic aberrations and no ghosts.

3 The VPH Grism Cross-disperser Another key element of the ESPRESSO optical design is the cross-disperser. Indeed, it is required to be very efficient, with higher enough dispersion,and to act as an anamorphoser to recover the strong elongation of the pupil, giving a more circular pupil onto the camera entrance aperture, and to decrease the projected slit height onto the detector to increase overall wavelength coverage. Volume Phase Holographic (VPH) gratings have been selected, being the most efficient known dispersers. To obtain the required anamorphic magnification, a prism can be put ahead of the grating. Another promising possibility, under study and prototyping, is given by “slanted-VPH”, where both VPH efficiency and anamorphism can be obtained by only one device.

ESPRESSO Optomechanics J. Pérez, H. Dekker, R.J. García López, J.M. Herreros, R. López, F. Pepe, J.L. Rasilla, P. Spanò and M.R. Zapatero Osorio

ESPRESSO is a new generation, super stable, high-resolution optical spectrograph, conceived to obtain the best performance in Doppler shift measurements. A concept study of the instrument is being carried out by a Consortium of scientific institutions formed by ESO, OG, IAC, INAF and Institute of Astronomy in Cambridge, all of them leaded by ESO (see Pasquini et al., these proceedings). Within this framework, the optomechanical, mechanical, and thermal conceptual design of ESPRESSO is being developed by engineers at the IAC. This contribution presents an overview of the current optomechanical conceptual design of this spectrograph.

1 General Mechanical Concept ESPRESSO is to be located inside the Coudé Combined Laboratory (CCL) at the VLT Observatory (Chile), which imposes some internal, spatial as well as transport and handling restrictions on our mechanical concept of the instrument (Fig. 1). Apart from them, for the mechanical study, we have taken into account the instrumental specifications, especially stringent in terms of structural and thermal stabilities, and issues related to the integration and management of the optical elements. The main subsystems are described next. Optomechanics The optomechanics include all the assemblies to locate and support optical elements in a tight way. The design relies on simple and robust solutions to keep optical elements aligned in a repeatable way without any configuration change or intervention during the instrument lifetime. Optical bench The optical bench (wire visualization in Fig. 2) is especially designed to trace out the optical layout by supporting and locating the optomechanics. It is conceived to provide the required thermal equilibrium (maximum temperature difference within the spectrograph ≤ 0.1 K) and structural stability (maximum flexure of the optical bench ≤ 0.01 mm). Vacuum vessel The spectrograph will be enclosed in a vacuum vessel as shown in Fig. 2. Vacuum conditions (below 10−4 mbar) are required to improve the thermal J. Pérez () Instituto de Astrofísica de Canarias, Vía Láctea, E38205 La Laguna, Tenerife, Spain e-mail: [email protected] A. Moorwood (ed.), Science with the VLT in the ELT Era, Astrophysics and Space Science Proceedings, doi: 10.1007/978-1-4020-9190-2_71, © Springer Science + Business Media B.V. 2009

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Fig. 1 ESPRESSO inside the CCL. A colour version of this figure is available at dx.doi.org/10.1007/978-1-4020-9190-2_71

Fig. 2 (Left) Side view of the optomechanics concept. (Right) Mechanical concept. A colour version of this figure is available at dx.doi.org/10.1007/978-1-4020-9190-2_71

stability and guarantee an stable optical path during the instrument lifetime. The vessel is pumped-out through a vacuum system located outside the thermal enclosure. Thermal enclosure The environment around the vacuum vessel, limited by a thermal enclosure, is to be thermally controlled at a constant reference temperature, through a thermal system, as another source of improvement of the thermal stability. The thermal system comprises both an air conditioning system and a warm-up system. The air conditioning system allows the thermal enclosure and the vacuum vessel to be blown with a continuous dry air flow, whereas the warm-up system allows the support structure legs to be thermally controlled. Some critical thermal stability specifications have to do with the maximum change of the spectrograph temperature, requiring to be better than 2 mK during 24 hours, 10 mK during 1 month and 30 mK during its lifetime.

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Support structure It is intended to support the vacuum vessel and all the elements rigidly attached to it, protecting them from any external vibration sources. Acknowledgements We thank the HARPS engineering team for providing us with plenty of valuable information.

ESPRESSO Science Software D. Mégevand, V. D’Odorico and C. Lovis

Abstract We present concepts for the software that will be used by scientists to prepare and launch their observations with ESPRESSO, then reduce and analyse their data. The ESPRESSO science software is mainly divided in three parts, the dataflow system (DFS) aiming mainly at the preparation of the observations, the data reduction software (DRS) generating reduced products of scientific quality common to all science cases and the data analysis software (DAS) processing the reduced data to obtain scientific results specific to each science case.

1 Introduction ESPRESSO is an instrument that will inherit positive concepts and experience from preceding instruments, combining the high efficiency of UVES with the high stability from HARPS, allowing some yet unattained results in specific science cases and fully exploiting the VLT potentialities. The software concepts will rely on the successful experience gained with HARPS, both at the night preparation level and at the online full reduction level.

2 Data Flow System The Observation Preparation Software will inherit capabilities from the HARPS Short Time Scheduler (STS). The targets scheduled or foreseen for one night are entered in an exposure list. The tool will offer the following features: exposure list editor, user catalogue interface, extended time calculator, optimising scheduler, OB exportation to a repository facility and world wide access. Archiving will follow the standard VLT archiving for the raw and the reduced frames, possibly also for the data analysis products. An appropriate tool will help the visiting observer or the night assistant to record the raw, reduced and data analysis product frames on a removable and transportable media (CD, DVD, HDD or any new medium available at the installation time). All raw data will be reprocessed from the ESO data archive at regular intervals, to ensure homogeneity on the entire set of data, independently from the date of the observations, allowing the data reduction algorithms to be continuously enhanced. D. Mégevand () Observatoire Astronomique de l’Université de Genève, ch. des Maillettes 51, 1290 Sauverny, Switzerland e-mail: [email protected] A. Moorwood (ed.), Science with the VLT in the ELT Era, Astrophysics and Space Science Proceedings, doi: 10.1007/978-1-4020-9190-2_72, © Springer Science + Business Media B.V. 2009

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3 Data Reduction Software The ESPRESSO Data Reduction Software will be based on the same concept as the HARPS DRS. Its goal is to provide online reduced data products of scientific quality that can be directly used as inputs for the main ESPRESSO science cases. At its top level, the DRS is a set of recipes performing specific tasks. Each observing mode (object-only, object-sky, object-reference and 4-UT) will have its own science recipe. The reduction of science raw frames will follow the following steps: bias and dark subtraction, bad pixels correction, background subtraction, order extraction with cosmic rays rejection, flat-fielding, wavelength calibration, merging and rebinning of the spectral orders, sky subtraction (if applicable), instrumental drift correction (if applicable) and flux calibration. The two main final products will be the extracted 2D (order, pixel) and 1D spectra and their associated error matrices. All relevant information accumulated during the reduction, such as wavelength calibration, cosmic hits maps or resolution maps, will be stored in the keywords of the FITS headers or in separate files.

4 Data Analysis Software The Data Analysis Software will provide a set of specific tools to the ESPRESSO users, useful to obtain the observables they need to accomplish their scientific research. The techniques used will include, among others, automatic spectrum co-addition, automatic continuum determination, automatic identification of metal lines, fit of lines, correlation and bisector computation. The critical points that need further investigation are the automatisation of the continuum level determination and the line fitting in crowded spectral regions (e.g. the Ly-α forest in QSO spectra). The final products of the Data Analysis Software will be normalised, co-added spectra; (preliminary) list of redshifts of metal systems; redshift, column density and Doppler parameter of absorption lines; equivalent width of absorption lines; radial velocities; bisectors and chromospheric activity indicators.

High Resolution Wavelength Calibration: Advancements with the Laser Frequency Comb Development A. Manescau, C. Araujo-Hauck, L. Pasquini, M.T. Murphy, Th. Udem, T.W. Hänsch, R. Holzwarth, A. Sizmann, H. Dekker and S. D’Odorico

Wavelength calibration in high resolution spectroscopy is a key issue for the delivery of high quality science data. ESO and MPQ have agreed to develop a wavelength calibration unit demonstrator based on a Laser Frequency Comb, which in comparison with current standards will significantly improve calibration data. The status and next milestones are summarized.

1 Introduction High resolution, high precision spectroscopy needs very stable and accurate wavelength standards [1]. To provide the ESO community with such standards for the current and future generation of high resolution spectrographs, optical frequency combs seem to be the adequate technology [2, 3] which will overcome limitations of current wavelength calibration standards [4]. ESO and MPQ have agreed to develop a demonstrator for a novel calibration system based on a laser frequency comb. This program has started already in 2007 and aims at exploring a prototype in the next three years.

2 The Laser Frequency Comb A self referenced laser frequency comb consists of thousands of equally-spaced frequencies over a bandwidth of several hundred THz (or several hundred nm) [2]. It is based on the properties of femtosecond mode-locked lasers. The shorter the laser pulses, the broader the range of frequencies in the comb. A. Manescau () European Southern Observatory, Karl-Schwarzchild-Str. 2, 85748 Garching, Germany e-mail: [email protected] A. Moorwood (ed.), Science with the VLT in the ELT Era, Astrophysics and Space Science Proceedings, doi: 10.1007/978-1-4020-9190-2_73, © Springer Science + Business Media B.V. 2009

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Fig. 1 Mode filter cavity setup at MPQ laboratories. On top left is the source (an Er-fiber laser) and the Hänsch-Couillard setup for locking the cavity to the laser. The mode filter cavity itself is shown at the bottom. On the top right side a photodiode is used for transmission measurement. A colour version of this figure is available at dx.doi.org/10.1007/978-1-4020-9190-2_73

The resulting modes of the frequency comb have their origin in the repetitive pulse train of the mode-locked laser. The mode spacing, which is constant in frequency space, is given by the pulse repetition frequency and resides in the radio frequency domain. The repetition frequency can readily be synchronized with a precise radio frequency reference such as an atomic clock. These clocks provide by far the most precise measurements of time and frequency currently available, the most reliably determined quantities in physics. Frequency combs therefore satisfy three requirements of the perfect calibration source which other methods do not: uniform line-spacing, long-term stability and reproducibility and known of each line frequency by first principles.

3 Status and Perspectives Our study shows that for a spectrograph like CODEX at R = 1.0 × 105 the optimal comb frequency is about 15 GHz [3]. The laboratory development started with a 3.75 GHz filter mode cavity, which is considered as an important step toward the 15 GHz cavity, as it is less demanding in terms of performance and will provide the necessary experience for the development. In Fig. 1 it is shown the first laboratory setup to measure and optimize performances of the filter mode cavity.

References 1. N. Santos, L. Pasquini, M. Romaniello (eds.), Precision Spectroscopy in Astrophysics, in press

Laser Frequency Comb Development 2. Th. Udem, R. Holzwarth, T.W. Hänsch, Nature 416, 233 (2002) 3. M.T. Murphy et al., Mon. Not. R. Astron. Soc. 380, 839 (2007) 4. C. Araujo-Hauck et al., in Proc. 2007 ESO Instruments Calibration Workshop, in press

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Precision Radial Velocities in the Infrared Hugh R.A. Jones, John Rayer, Larry Ramsey, Bill Dent, Andy Longmore, Bill Vacca, Mike Liu, Adrian Webster, Alex Wolscznan and John Barnes

Over 250 extra-solar planets have been discovered to date using a variety of techniques. The majority have been discovered at optical wavelengths from the Doppler shift of F, G and K stars induced by orbiting planets. We have constructed models simulating likely planets around M dwarfs and demonstrated the ability to recover their radial velocity signals in the infrared. We have conducted experiments in the infrared with a brass-board instrument to explore real-world issues. We are thus confident that a stabilised radial velocity spectrometer with a single-shot 1 and 1.7 microns coverage at a resolution of around 70 k can achieve an instrumental radial velocity error of 0.5 m/s. This enables the efficient measurement of radial velocities for M, L and T spectral classes. We have modelled the radial velocity information in low-mass star spectra and checked our ability to recover this signal in the face of the telluric contamination in the infrared. Including instrumental error, telluric contamination and photon noise we predict a total radial velocity error of less than 2 m/s on a typical M6V star at 10 pc. We use these results as an input to a simulated 5-year survey of nearby M stars envisaged for Gemini Observatory. Based on a conservative scaling of optical results, such a survey has sensitivity to detect several terrestrial mass planets in the habitable zone around nearby stars. It can test theoretical planet formation models, which predict an over-abundance of terrestrial-mass planets around low-mass stars. Improvements in the efficiency and sampling of searches at optical wavelengths promise long-term precisions of 0.5 m/s and 5 MEarth detections around solartype stars. While this may be the limit for CCD-based surveys of solar type stars until larger telescopes become available it is nonetheless feasible to survey lower mass primaries to achieve a corresponding smaller mass limit. Thus the lowest mass Doppler signals have been found around M dwarfs (e.g., GJ581b 5 MEarth sin i [1]) and detections down to a few MEarth detections should be feasible around mid-type M dwarfs. This is the Precision Radial Velocity Spectrometer (PRVS) approach: to search around lower-mass primary stars since the radial velocity signal will be larger for lighter primary stars (http://www.roe.ac.uk/ukatc/projects/prvs). Based on the spectral information in real and synthetic spectra of M dwarfs, PRVS is designed to measure the peak of their energy distribution. We find that for low rotation M dwarfs a resolution of around 70 000 is optimum. This leads to us H.R.A. Jones () Centre of Astrophysics Research, University of Hertfordshire, Hatfield, UK e-mail: [email protected] A. Moorwood (ed.), Science with the VLT in the ELT Era, Astrophysics and Space Science Proceedings, doi: 10.1007/978-1-4020-9190-2_74, © Springer Science + Business Media B.V. 2009

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Fig. 1 Radial velocity amplitude as a function of host mass or mean habitable zone distance comparing PRVS with optical radial velocity surveys. A colour version of this figure is available at dx.doi.org/10.1007/978-1-4020-9190-2_74 Fig. 2 Number of exoplanets known per spectral type based on exoplanet.eu shown with a think line. The thick line shows the breakdown of stars within 10 pc (still incomplete for M dwarfs) from RECONS (http://www.chara.gsu.edu/ RECONS/TOP100.htm)

using an R4 31.6 lines/mm echelle grating illuminated in pseudo-Littrow mode. During the PRVS design study, Ramsey et al. [2] built a simple version which has obtained RMS precisions of better than 10 m/s for integrations of a minute over several hours on the Sun in the Y band giving us confidence that many of our basic design assumptions are reasonable.

References 1. S. Udry et al., Astron. Astrophys. 469, 43L (2007) 2. L. Ramsey et al., Bull. Am. Astron. Soc. 86, 1016 (2006)

Very Large Spectroscopic Surveys with the VLT I.R. Parry

Very large spectroscopic surveys (several million spectra) are required to advance our understanding of Dark Energy (via baryonic wiggles) and the detailed history of our Local Group of galaxies (via Galactic Archaeology). In this paper I make a preliminary exploration of how this might be done by putting a wide field, optical, prime-focus fibre-fed spectroscopic facility on one of the VLT’s UTs.

1 Introduction The Gemini Aspen process identified the need for a wide field (1–2 deg. diam.), optical, multi-object (few thousand at a time) spectroscopic survey facility to study Dark Energy (DE) and Galactic Archaeology (GA). A feasibility study was completed in March 2005 and following that, two competing design studies are currently being carried out for a system which will go at prime focus on the Japanese Subaru Telescope. This proposed system is called WFMOS (Wide Field Multi-Object Spectrograph). Four potential surveys were identified. (1) low-z DE survey, (2) high-z DE survey, (3) low-resolution GA survey, and (4) high resolution GA survey. These require a total of about 7 million spectra and about 600 nights of observing. Although the Gemini design study process is well underway it is still worth considering doing a similar project on the VLT for several important reasons. The GA case is better served from the southern hemisphere which gives access to the galactic bulge and the Magellanic Clouds. It is very likely that most future large investments in ground based facilities (i.e. ELTs) will be in the South and this instrument will be an important path-finder for ELT follow up studies. At the time of writing, no European astronomers have access to Gemini. The facility could be used for science other than DE and GA (basically it is about 50 times faster than the current FLAMES instrument). On the VLT one of the UTs could be 80%–90% dedicated to this project (the remainder going to VLTI) and the surveys could be done in about 2 years—this cannot happen at Subaru where the fraction of time given to WFMOS would be much less. WFMOS is a design study and it is not 100% certain that it will be fully funded and continue in to the construction phase. Below I present a very preliminary discussion of the technical, operational and cost implications of putting something like WFMOS on the VLT. I.R. Parry () Institute of Astronomy, Madingley Road, Cambridge, CB3 0HA, United Kingdom e-mail: [email protected] A. Moorwood (ed.), Science with the VLT in the ELT Era, Astrophysics and Space Science Proceedings, doi: 10.1007/978-1-4020-9190-2_75, © Springer Science + Business Media B.V. 2009

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2 Instrumental Requirements The science case for the dark energy and galactic archaeology surveys demands ∼ 2400 (or more) simultaneous spectra at optical wavelengths for discrete objects in a 1.5 deg. FOV on an 8 m telescope with up to 20 field configurations per night. The large FOV, the large wavelength coverage and the high spectral resolution are best addressed by a fibre-fed multi-object spectroscopy (MOS) approach with one fibre per object. The field change time therefore has to be fast, < 5 minutes from the end of one exposure to the start of another one on a new field. Up to 2400 × 20 = 48 000 fibres have to be placed in a 24 hour cycle (for one of the DE surveys). Each fibre tip ideally has to be placed accurately with 5 degrees of freedom: X, Y for the object position, Z for focus and tip and tilt for pupil aiming. The Galactic Archaeology high resolution survey drives the spectrograph design. At a spectral resolution R ∼ 20 000, in the wavelength range from 480 nm to 680 nm, a spectral resolution element (SRE) is 0.029 nm. One spectrum has 6897 SREs, which is 245 760 detector pixels for a single spectrum (assuming an f/1.8 camera and allowing for gaps between the spectra). Given the surface density of targets, 800 simultaneous spectra are needed which requires at least 197 million detector pixels. A 16k × 12k detector format has 201 million pixels. These numbers also assume that the f-ratio into the fibres is about f/2.3, from the prime focus corrector (PFC) and the f-ratio of the collimator is about f/2.0 to allow for focal ratio degradation. A fibre diameter of about 92 microns (1 arcsec) is assumed and the 201 Mega-pixel detector format is achieved with 24 chips each with 2048 × 4096, 15 micron pixels. A large beam diameter is required to achieve the high spectral resolution of R ∼ 20 000. The spectrograph(s) will therefore be very large and will have to be located on one or both Nasmyth platforms. The fibre positioning system is a major challenge. Magnetically held fibres (OzPoz/2df/Autofib type systems) cannot cope with the large numbers required and neither can fishermen round the pond systems (like MX and KMOS). Systems which have a micro-positioner per fibre (e.g. like the Echidna system on FMOS/Subaru) are being investigated as part of the WFMOS design study and can feasibly deliver the ∼ 2400 fibres required in the FOV which is about 500 mm in diameter. Plugplate systems have the greatest potential in terms of the number of simultaneous on-sky fibres because they have the minimal amount of hardware associated with each fibre tip. A very preliminary PFC design is shown in Fig. 1. This design delivers excellent image quality (point source spots smaller than 0.25 arcsec) over the whole FOV. However, the mass of the glass alone in this design is 965 kg. Furthermore, it needs an atmospheric dispersion corrector. More work is required to optimise the design in terms of it’s mass, telecentric properties, and number of aspheric surfaces. It may well be possible to increase the FOV above 1.5 degrees but this will have significant mass and cost implications. Although this design is not optimal it does give a good estimate of the size and location of the focal plane and the location of the PFC optics.

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Fig. 1 Preliminary VLT PFC design with a 1.5 degree unvignetted FOV. A colour version of this figure is available at dx.doi.org/10.1007/978-1-4020-9190-2_75

Fig. 2 The M2 unit on UT4 showing some key dimensions and masses. A colour version of this figure is available at dx.doi.org/10.1007/978-1-4020-9190-2_75

Figure 2 shows the existing M2 unit (secondary mirror unit) on UT4 with approximate dimensions. This shows that there is very little space for a fibre positioning system given the potential for collisions with the existing enclosure infrastructure. If the mass of the existing M2 unit (∼ 1700 kg) cannot be exceeded then it will also be hard to fit a PFC, an acquisition and guidance system (including the optics to control the figure of the primary mirror) and a fibre positioning system at prime and stay within the mass budget.

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3 A Plug-Plate Solution to the Fibre Positioning Problem In this scheme the spectrographs and most of the length of the optical fibre bundle are permanently mounted on the telescope. The long fibre bundle is terminated at the top with an efficient, multi-way fibre-optic connector. The plugplate itself is a sheet of suitable material with accurately drilled holes into which the fibres are inserted. The plugplate modules are configured off the telescope. With exposures of ∼ 30 minutes, up to 20 pre-configured modules have to be ready for use at the start of a night. A special machine interchanges the modules on the telescope during the night: for each observation a pre-configured fibre module is mounted at the focus and connected to the fibre bundle. In each module the fibres are short (∼ 500–1000 mm) and the output ends are arranged in a fixed pattern (a fibre-optic connector). At the end of an exposure the telescope moves in altitude and azimuth so that the prime focus is at the access platform where a specialised module interchanger can access the fibre module. The permanently installed part of the fibre connector is backed off from the connector half on the plugplate module. The used module is then removed and the next one is put on. Next, the new module is referenced to the telescope focal plane. The permanently installed part of the fibre connector then mates with the other half of the connector on the module. The telescope slews to the next field. This entire process is the same for all modules. The fibre-modules are prepared in a dedicated space, probably at the base camp. The plates are drilled on CNC machines. The facility has enough machines to accurately drill up to 48 000 holes in an 8 hour shift. (e.g. 20 machines are needed if the time to drill one hole is 12 s). The plates are non-metallic to reduce costs and reduce fitting tolerances. The fibres are removed from used plates and inserted into new plates by either robots or people. The facility has enough fibre-pluggers to reconfigure 48 000 fibres in an 8 hour shift. For example, 20 fibre-pluggers are needed if the time to move one fibre is 12 s. A special quality control machine is used to check and clean each module after it has been prepared. Each module is bar-coded so that the system can keep track of everything. The modules have to be transported between the telescope and the preparation facility each day. The choice between people or robots comes down to cost. The cost of employing people is straightforward to calculate. For 12 s per fibre and a total of 7 million spectra we have 23 333 hours of work. Allowing for rest periods the total work time for the plugging team could be 30 000 hours. At a pay rate of 60 euros per hour (equivalent to a cost to ESO of about 100 000 euros per year per person allowing for weekends and leave) the fibre plugging team costs a total of 1.8 M euros to complete the surveys. The cost of using robots is much harder to quantify. The use of plugplates has many advantages. They have the least possible positional constraints: there are no magnets or prisms or crossing fibres. There are no positioning mechanisms attached to the fibres. This is fundamental for scientific versatility! Operation during the night is extremely simple and so there is a very low risk of telescope time being lost compared to systems which reconfigure the fibres on the telescope. The DE and GA surveys have quite different target density requirements (800 per field for the high resolution GA survey cf. up to ∼ 7000 for the high

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redshift DE survey: both of these utilise all the detector pixels). Plugplates handle this difference much better than other positioners. The interface with the spectrographs is simplified (no need for fibre selection mechanisms at the slit end). All 5 axes for fibre placement (X, Y, Z, tip, tilt) can be controlled if required. The modules can be very small and light and so they help with the tight mass and space budgets on the VLT. By having extra fibre-pluggers, extra fibre-modules and extra drilling machines, a routine maintenance program will ensure that production always continues at the required pace. When fibres break, initially this only impacts on “sky” fibres so the surveys are not affected. Eventually, once too many fibres are broken in a module it is taken out of operation and repaired. Similarly, by having spare fibre-pluggers and drilling machines there are always enough operational to meet the survey needs.

4 Conclusions and Discussion The DE and GA science cases are extremely attractive. Furthermore, there are other exciting projects (for example, scaled up versions of projects currently being done with FLAMES) which the facility envisaged here would be capable of doing. Optically, a prime focus corrector design with a 1.5 degree FOV is entirely feasible. However, the space and mass constraints are tight. An investigation into what drives the mass constraint would be useful—it may be possible to put more than 1.7 tonnes at prime. Also, the space limitations could be relaxed significantly by modifying some parts of the enclosure infrastructure or by using anti-collision interlocks. The spectrographs need 201 million pixels but are feasible and would probably fit on a single Nasmyth platform. The total cost will be substantial (for comparison Gemini have decreed that WFMOS has to cost less than US$45M). A proper study is required to establish the cost of a similar facility on the VLT. If this facility is implemented as a non-interchangeable, dedicated facility the VLT loses 3 instrument stations and more significantly the option of using VLTI with all 4 UTs for 2–3 years. On the other hand if the facility is designed to be changed back to the M2 unit within a day then it would not compromise the use of all four UTs for VLTI at all and only remove one instrument station (assuming the spectrographs fit on one Nasmyth platform and are never dismounted).

New Developments in Integral Field Spectroscopy Anthony Horton, Joss Bland-Hawthorn and Simon Ellis

1 Introduction The development of wide field adaptive optics (AO) capabilities greatly increases the potential information content (spatial resolution elements × spectral resolution elements) per telescope pointing. This is illustrated in Table 1, which lists the wide field AO systems for the VLT. The multi-conjugate adaptive optics (MCAO) demonstrator MAD has already proven that a 1 corrected field of view (FoV) with point spread function (PSF) full width half-maximum (FWHM) of 0.1 is possible [1]. This corresponds to a number of spatial resolution elements, Nsp = 4 × 104 . The two ground-layer adaptive optics (GLAO) systems currently under development promise comparable numbers. Comparing these figures with the seeing limited case for the maximum VLT FoV and good seeing we see that in terms of the number of spatial resolution elements the wide field AO foci of the VLT will be just as ‘wide field’ as any seeing limited instrument. This simple comparison highlights the need for large spatial format AO instruments, including integral field spectrographs (IFS’s), to fully exploit the scientific potential of wide field AO. However, large format IFS’s are complex and technically challenging instruments. A good example of the current state of the art is MUSE, a 300 × 300 format, R = 2000–4000 IFS for wavelengths of 0.465–0.93 µm, giving a total of 7 × 107 resolution elements. One of the most obvious difficulties in building such instruments is simply physical size, MUSE will essentially fill the Nasmyth platform at the VLT. The technical difficulties are compounded for near infrared (NIR) IFS. The need for more extensive cooling leads to even greater size and complexity. Furthermore NIR IFS is less capable than optical IFS because the NIR sky background is much higher, which limits NIR IFS to high surface brightness objects. However, AAO technology developments, in particular integrated photonic spectrographs and OH suppression fibres, offer possible solutions. Anthony Horton () Anglo-Australian Observatory, PO Box 296, Epping NSW 1710, Australia e-mail: [email protected] A. Moorwood (ed.), Science with the VLT in the ELT Era, Astrophysics and Space Science Proceedings, doi: 10.1007/978-1-4020-9190-2_76, © Springer Science + Business Media B.V. 2009

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Table 1 Wide field AO systems for the VLT AO system/inst.

AO type

FoV

FWHM

Nsp

GALACSI/MUSE

GLAO

1 /7.5

0.35 /0.04

4 × 104

MAD

MCAO

1

0.1

4 × 105

GRAAL/Hawk-I

GLAO

7.5

0.2–0.5

1–5 × 106

–

7.5

0.5

8 × 105

Seeing limited

2 Integrated Photonic Spectrograph An integrated photonic spectrograph (IPS) is an integrated optics device which combines all the functions of a spectrograph, i.e. ‘aspectrograph on a chip’ [2].

2.1 Basis The IPS idea is based on devices developed for dense wavelength division multiplexing (DWDM) applications for telecommunications. The IPS device consists of a wafer of glass (or other material) containing a planar waveguide circuit. Light input is via a few or single mode optical fibre, and the output spectrum is edge coupled to a detector array. Two main technologies are available, photonic echelle gratings (PEGs) and array waveguide gratings (AWGs). AAO IPS development is focused on AWGs due to easier fabrication.

2.2 Instrument Concepts Instruments will consist of an ensemble of IPS’s, one per spatial sample. This can be viewed as the natural conclusion of the trend toward modular construction for large spectroscopic instruments. Individual IPS’s can be tightly packed, with multiple IPS’s printed on a wafer and individual wafers arranged in stacks. Detectors arrays are edge coupled to output regions of the wafers, either one linear array per IPS or 2D arrays fed by stacks of wafers.

2.3 Benefits Several benefits stem from diminutive size. An individual IPS with R ∼ 1000 is typically only centimetres in size, and can be even smaller [3]. As IPS’s can be densely packed the overall instrument will also be small in size. This greatly reduces the difficulties associated with cooling, temperature control, and flexure issues. Size

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may be reduced sufficiently that deployable integral field spectrographs (as opposed to deployable integral field units) become possible. As a compact, solid state device an IPS can be expected to be highly stable. An IPS based instrument is intrinsically modular and scalable. Finally, the manufacture of IPS’s lends itself to true mass production, with resulting reductions in costs.

2.4 Development Programme The first prototype device has already been produced by the AAO and a commercial partner. Shown in Fig. 1, it is a modified DWDM AWG with single mode fibre input and an output spectrum formed in free space. This device has enabled the first observation of a continuous spectrum produced by an AWG. More details can be found in Bland-Hawthorn and Horton [4]. An ongoing development programme is planned, focusing on lower resolution devices (R = 2000, 1000, 500, 250) with flattened blaze envelopes and broader spectral response. Future prototypes will have flattened focal planes to facilitate direct coupling to detector arrays. We will also pursue the stacking of wafers and bonding to 2D detector arrays. Few-mode input will be investigated too.

3 OH Suppression Another AAO technology of potentially great benefit is the suppression of OH airglow lines using asynchronous fibre Bragg gratings (AFBGs) [5]. Uniquely high fidelity suppression is possible with this technique which offers strong suppression

Fig. 1 IPS prototype device. A colour dx.doi.org/10.1007/978-1-4020-9190-2_76

version

of

this

figure

is

available

at

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(30 dB) of a large number of lines (> 150) at high resolution (R > 10 000) with high interline throughput (∼ 90%). A key benefit over existing OH suppression techniques is that suppression takes place prior to any dispersion, which greatly reduces the contribution from scattering wings of OH lines to the interline background. Truly effective high resolution OH suppression upstream makes sensitive NIR spectroscopy possible at low resolutions (R < 2000). For more detail see Ellis et al. [6] and references therein.

4 The Importance of AO Both IPS and OH suppression fibres are subject to modal dispersion, i.e. different waveguide modes have different effective refractive indices. For IPS this results in spectral shift between modes reducing the effective resolution, while for OH suppression the suppression notches are shifted in wavelength. Modal dispersion can be dealt with, but at a cost. OH suppression fibres can use wider notches, multiple gratings in a single fibre, or MMF-SMF converters and multiple SMF AFBGs [7]. For IPS we can use low resolution, design to minimise modal dispersion or use MMF-SMF converters and multiple IPS’s. To minimise the complexity and maximise the performance of IPS or OH suppression fibres we need to minimise the number of modes. However, incoming light is filtered by the mode ensemble of the input fibre [8] so the coupling efficiency is greater for a larger number of modes or for less aberrated wavefronts. Good AO therefore allows higher throughput with fewer modes, improving the cost/benefit ratio of both IPS and OH suppression.

5 Science with OH Suppressed IPS Science enabled by OH suppression in general, including MOS science, is discussed in Ellis et al. [6]. NIR IFS specific programmes which could benefit from (OH suppressed) IPS include the following: • • • • •

Resolved solar system bodies. Exoplanet detection, using the same principle as SPHERE IFS. YSO/proplyd outflow dynamics from H2 lines. In depth investigation of nearby (U)LIRGS/starbursts. Spatially resolved studies of high redshift galaxies: – Kinematics i.e. discs and mergers, dynamical mass measurements. – Metallicity. – Star formation rate. – Separation of AGN/starburst contributions. • Epoch of reionisation surveys for z > 7 Lyman-α emitters, via lensing caustic observations. • MUSE style 3D surveys in the NIR.

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6 Highly Configurable Coverage While IPS’s can enable a large increase in the number of simultaneous spatial samples for NIR spectroscopy there are still significant gains in observational efficiency possible from flexible instrument designs. Small, tileable, deployable IFUs allow almost arbitrary spatial coverage of the focal plane, so that the utilisation of spatial samples can be optimised even for fields with diverse target sizes, shapes and separations. With modular spectrographic instruments there is also scope for optimisation in the use of spectral resolution. An IPS based instrument could consist of several types of device offering different resolutions, enabling resolution to be chosen separately for each target to best suit object flux and science objectives. Similarly you could utilise different spatial resolutions, for example a hybrid MOS/IFS instrument or instruments incorporating dIFUs with different spatial sampling.

7 Conclusion There is a strong need for large format, AO-fed NIR IFS, but building such instruments is not easy. The IPS is a scaleable, modular technology for building NIR IFS’s which promises large formats at reduced size and cost. The first prototype IPS has been produced, and a development programme is under way. The complementary technology of OH suppression fibres will deliver greatly increased NIR spectroscopic sensitivity, further strengthening the case for large format NIR IFS and hence IPS. Flexible, diversified, modular instruments with highly configurable spatial and spectral coverage and several spectral and spatial resolutions can give large gains in observing efficiency though optimised deployment of limited detector real estate.

References 1. 2. 3. 4. 5. 6.

E. Marchetti, R. Brast et al., ESO Messenger 129, 8 (2007) J. Bland-Hawthorn, A.J. Horton, Proc. SPIE 6269, 62690N (2006) P. Cheben et al., Opt. Express 15, 2299–2306 (2007) J. Bland-Hawthorn, A.J. Horton, AAO Newsletter 112, 27 (2007) J. Bland-Hawthorn, M. Englund, G. Edvell, Opt. Express 12, 5902–5909 (2004) S.E. Ellis, J. Bland-Hawthorn, A.J. Horton, R. Haynes, in Science with the VLT in the ELT Era, ed. by Alan Moorwood. Astrophysics and Space Science Proceedings (Springer, Netherlands, 2009), p. 437 7. S.G. Leon-Saval, T.A. Birks, J. Bland-Hawthorn, M. Englund, Opt. Lett. 30, 2545–2547 (2005) 8. A.J. Horton, J. Bland-Hawthorn, Opt. Express 15, 1443–1453 (2007)

ULTRAPHOT Françoise Roques, Isabelle Guinouard, Jean-Tristan Buey, Alain Doressoundiram, David Horville and Michel Marteaud

An additional configuration of FLAMES is proposed for multi-objects very fast photometric observations. UltraPhot will conduct photometry of faint objects in a large field of view, on timescales smaller than one second. Variations of faint sources on timescales longer than a few seconds are explored by conventional CCD instruments. Few instruments allow high-speed observations, but they are limited to very small field of view (UltraCam). UltraPhot will allow high speed and high precision photometry of large number of faint objects on timescales of seconds to milliseconds in a large field (25 diameter). Potential interested scientific fields are numerous: Exploration of Outer Solar System Objects (Kuiper Belt, Oort Cloud) by stellar occultations: This method, commonly used to study solar system objects, is able to detect objects invisible by direct observation. Blue stars must be chosen to minimise their angular size, and two colours allow to get the diffraction signature of the occultation. The aim is to explore the small end of the size distribution and the outer part of the Solar System. A challenge is the possibility to explore the Oort Cloud (a comet of 5 km is detectable at 10 000 AU). Extrasolar Planets Transits: Ultraphot will allow follow up of the Corot/Kepler targets. High precision photometry associated with a large field of view (1/20 of the Corot Exoplanets FOV) makes Ultraphot a powerful tool to explore the transit exoplanetary systems and to detect small planets. Moreover, precise timing of transits give access to other planets in the system by their perturbations on the orbits parameters. Study of variable objects in Globular Clusters would also benefit from this instrument. Searching for rapid periodic variations of blue objects might help to identify their nature. In particular Compact Binaries (CB) including Cataclysmic Variables are predicted to be numerous in GC not only in the cores but also in their outskirts. Cross-correlation with X-ray sources can provide further constraint on CB. Periodicities vary on a wide range depending on the nature of the objects. Blue colour is essential to isolate the CV and avoid crowding from red stars. Blue Stragglers are GC blue objects which also would be worth to look for variability. F. Roques () LESIA, Observatoire de Paris, 92195 Meudon cedex, France e-mail: [email protected] url: http://www.lesia.obspm.fr/ A. Moorwood (ed.), Science with the VLT in the ELT Era, Astrophysics and Space Science Proceedings, doi: 10.1007/978-1-4020-9190-2_77, © Springer Science + Business Media B.V. 2009

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The basic idea of this instrument is a large number (100–150) of single fibres added on the FLAMES plates. Each fibre flux is separated in two channels and positioned on few lines of a CCD. A possible choice for the CCD is the back illuminated E2V 42-10 full frame CCD sensor, with 2048 × 512 pixels and a spectral range: 200–1060 nm. This CCD allows an acquisition frequency as large as 100 Hz, with a high quantum efficiency: (85% at 500 nm) for 130 fibres. The maximum number of fibres is not limited by Ozpoz which can host up to 560 fibres per plate but by the read out time of the instrument. UltraPhot would benefit of the large FOV (25 ) of FLAMES. The instrument could be small enough (600 × 700 × 300–400 mm) to be placed inside the Giraffe box, above the spectrograph itself. The fibres aperture is important. A large aperture increases the sky flux but a too small aperture increases flux lost du to scintillation. Analysis of lightcurves computed from Ultracam/VLT images shows that a 5 circular aperture is a good compromise, and the sky background (15 mag for bright time and 19 for dark time) could be precisely measured with fibres dotted around the field. A preliminary study has shown that there are two options for the optic way, with one or two CCD. The fibres manufacturing (including manpower) is estimated to 40 k. The CCD (including electronic and PC) cost can be 70 k (100 k in the 2 CCD option). Pre-study and study could need one FTE year and the manufacturing another FTE year. The following teams are interested to participate to the consortium, with a scientific implication in the above subjects. The financial and/or manpower implications could cover partial or complete cost of the instrument: Napoli observatory, Institut d’Astrophysique de Paris, Tsing Hua University in Taiwan, GEPI and LESIA of the Paris Observatory.

References 1. F. Roques, G. Georgevits, A. Doressoundiram, The Kuiper belt explored by serendipitous stellar occultations, in The Solar System Beyond Neptune, ed. by M.A. Barucci, H. Boehnhardt, D.P. Cruikshank, A. Morbidelli (University of Arizona Press, Tucson, 2008), pp. 545–556

Super-GIRAFFE: The Next Generation High Multiplex Optical Spectrograph with d-IFUs M.D. Lehnert, I. Guinouard, D. Horville, P. Jagourel, F. Chemla, J.-P. Amans, P. Bonifacio, C. Babusiaux, F. Hammer, V. Hill, F. Royer and M. Puech

Abstract We propose to build super-GIRAFFE, a high multiplex (several hundred or more fibers), very efficient (total throughput > 20%) spectrograph with deployable IFUs (20–30) that simultaneously covers the full optical band (0.35–1.0 µm) at two spectral resolutions (R ∼ 5000 and 15 000–20 000). The total field of view of the instrument would be like that of the current FLAMES/GIRAFFE (25 ). Our focus with such an instrument is to make it highly efficient, with a throughput similar to that of a simple multi-slit spectrometer, and high multiplex. Such a spectrograph could be build cheaply and quickly by upgrading the fiber positioner of FLAMES/GIRAFFE. By building it quickly, we believe the VLT will have a unique capability for tackling such important issues as the growth of large scale structure, the tomography of the inter-galactic medium, the spatially-resolved dynamics and emission/absorption line distributions of high redshift galaxies, stellar populations in the halo of the MW and in other nearby galaxies, etc., before anyone else. Also, “keeping it simple” has the advantage that several copies could be built, allowing for truly significant surveys to start within a few years, rather than within a decade. We also discuss possible enhancements to this baseline design.

1 The Need for a New High Multiplex Spectrograph at VLT When developing a concept for an instrument, the first consideration is what types of astrophysical problems need to be addressed. In fact, what problems need to be addressed in the years of operation of the instrument. The complexity of modern astronomical instrumentation is such that one often has to gaze far into the future and make only educated guesses. Therefore, it is often useful to look at the broad issues in astrophysics that will no doubt be of keen interest for many years or decades. These pertinent questions are: How do photons escape galaxies? Where are the missing baryons and metals? Why is galaxy formation so inefficient? How did the MW form? What did the first major generation of star-formation look like? How did the metals in the Universe grow? How are the “cycled”? How does environmental density influence galaxy growth and evolution? How did large scale structures in the Universe grow? M.D. Lehnert () GEPI, Observatoire de Paris, 5 Jules Jannsen, 92190 Meudon, France e-mail: [email protected] A. Moorwood (ed.), Science with the VLT in the ELT Era, Astrophysics and Space Science Proceedings, doi: 10.1007/978-1-4020-9190-2_78, © Springer Science + Business Media B.V. 2009

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To address these large astrophysical questions robustly, any next generation ground based instrument should be flexible, with a high effective multiplex, and have the ability to look at a smaller number of objects in an integral field mode. From these considerations and focusing in on the physical phenomenon exhibited by galaxies led to the development of instruments like VIMOS [4] and KMOS [6] on the VLT. Along the same lines, we now propose that the ESO member states should build a complement to KMOS in the optical, while retaining the ability to do extremely high multiplex single object spectroscopy at moderate to high spectral resolutions. Such an instrument is likely to become the “workhorse” or “workhorses” of the VLT if several instruments like this are built.

2 Baseline Design Proceeding from the large questions posed above, to more specific projects such as measuring the absorption lines in the intergalactic medium in hundreds of QSOs and galaxies (so-called IGM tomography), to the metal abundances and ratios of stars in our halo and in nearby galaxies, to spatially resolved studies of distance galaxies and AGN hosts, etc., led us to conclude that we should propose to essentially upgrade the current combination of FLAMES and GIRAFFE [1, 2, 5] to make the most capable optical spectrograph on any 8 m class telescope. We outline the baseline concept in Table 1. Our baseline design was driven mostly by a desire to increase the overall efficiency, making it similar to that for slit-let spectrographs, to increase the instantaneous wavelength coverage to include the entire optical band (and also avoiding the strong OH lines in the near-infrared) and to increase the overall multiplex by increasing the number of fibers that can be targeted on individual sources. We believe this increase in efficiency, multiplex, and wavelength coverage can be accomplished by: (1) splitting the beam into two channels—blue and red optical. This will allow us to use detectors and gratings optimized for each channel and thus Table 1 Baseline Design for Super-GIRAFFE Function

Baseline

Goal

Number of IFUs

20

40

FOV of Each IFU

3 × 3 arcsec

5 × 5 arcsec

IFU sampling

0.3 arcsec

0.2–0.4 arcsec

Patrol Field

25 arcmin

Larger? 15 000–20 000

Resolution

5000

λtotal

0.35–1.0 µm

0.8–1.8 µm

Dichroic

0.57 µm

1.0 µm

Many Fibers Modea

∼ 250–300

> 500

Total Efficiency

20%

a This

is equivalent to the MEDUSA mode in FLAMES/GIRAFFE

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433

Fig. 1 An example of the throughput of a fiber. With such high throughput, a fiber based spectrograph can reach the efficiency of simple slit spectrographs. However, the key issue is coupling both the light in the focal plane to the fiber and coupling the fiber output to the spectrograph. We propose to use microlens arrays to minimize the focal length degradation inherent in this type of design. A colour version of this figure is available at dx.doi.org/10.1007/978-1-4020-9190-2_78

gaining a significant increase in the overall efficiency, especially λ < 0.45 µm and λ > 0.8 µm. Depending on the wavelength, using volume phase holographic grating will lead to an increase in overall efficiency of 10 s of percent, while using improved red and blue optimized detectors will lead to an increase of 10% to more than double. (2) Minimize the focal ratio degradation by coupling microlenses to both the fiber inputs and outputs. This will allow us to better match the output focal ratio of the fiber to that of the spectrograph getting an increase in overall efficiency of about 25%. Our goal is to make this instrument with a throughput that is similar to simple multi-slit spectrographs such as FORS1 on the VLT. The added advantage of superGIRAFFE is the much larger field of view (25 vs. ∼ 7 ), larger number of targets (several hundred versus about several 10 s in the single fiber mode or MEDUSA mode as it is called in the current FLAMES/GIRAFFE), less slit contention in targeting objects when using the single fibers (MEDUSA) mode than a slit-let design, and a deployable IFU mode with a significant multiplex. Compared to VIMOS [4], similar considerations apply with regards to the spectral resolution, multiplex advantage having larger numbers of targets in the single fiber mode, deployable IFUs capability, and to the large total field of view of the instrument.

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Fig. 2 Here we show a possible configuration for matching the speed of the telescope input to the fiber entrance and then matching the speed of the fiber output to the spectrograph using microlens arrays. This is only illustrative of several possible solutions from minimizing focal ratio degradation. One advantage of sub-sampling the fiber input with microlenses is that it allows us to keep fixed spectral resolutions in both the “many fibers mode” and in the IFU mode and allows us to possible reach high spectral resolutions (R = 15 000–20 000) in the many fibers mode of observing. A colour version of this figure is available at dx.doi.org/10.1007/978-1-4020-9190-2_78

3 Ideas Beyond the Baseline Our current baseline is provided by maintaining a fairly standard spectrograph with an upgrade of FLAMES from about 130 fibers to more than doubling this number and increasing the number of deployable IFUs which sampling better matched to the seeing on Paranal. We are emphasizing high efficiency and low cost and attempting to avoid any major new technological developments. We are simply taking advantage of better technology, two channels and larger detectors compared to the

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current FLAMES/GIRAFFE. Obviously, there are changes in the baseline design that could be made, such building a totally new fiber positioner with a significantly larger number of fibers (say 600). We think having a high spectral resolution mode, R = 15 000–20 000 is critically important for stellar chemical abundances and absorption lines in the intergalactic medium. The question yet to be determined is how to combine the two resolutions and what should be the wavelength coverage in this mode. And finally, it might be of interest to consider making another type of superGIRAFFE that has near-infrared wavelength coverage, say out to about 1.7 µm. This of course comes with the disadvantage of brighter night sky lines and more costly detectors. However, if multiple copies of the super-GIRAFFE concept could be built, one could consider building the first version cheaply following the baseline concept while working on a more capable section version with higher multiplex or red optical to near-infrared wavelength coverage.

4 Conclusions We believe that if one or several spectrographs like the one outlined here are built, ESO will have a truly exceptional facility for addressing some of the most outstanding problems in astrophysics. The simplest of the concepts presented here could be built relatively cheaply (less than a few MEuros) and quickly (few years) based on already proven techniques and technologies. Going beyond this, for example a higher multiplex and IFUs with larger field of views or a near-infrared channel will of course require more time and money. However, building something more modest may also allow the ESO community to address important astrophysical questions within a few years instead of within a decade and solve many problems associated with selecting the best targets for detailed studies with the ELT, ALMA, and JWST. This “prime mover” advantage should not be underestimated and should be one of the main considerations in driving the design of any next generation spectrograph on the VLT.

References 1. 2. 3. 4. 5. 6.

G. Avila, I. Guinouard, L. Jocou, F. Guillon, F. Balsamo, SPIE 4841, 997 (2003) L. Jocou, I. Guinouard, F. Hammer, H. Lenoir, SPIE 4008, 475 (2000) O. Le Fèvre et al., ESO Messenger 109, 21 (2002) O. Le Fèvre et al., ESO Messenger 111, 18 (2003) L. Pasquini et al., ESO Messenger 110, 2 (2002) R. Sharples et al., ESO Messenger 122, 2 (2005)

FLEX (The First Light Explorer)—The Science Case for a Fully OH Suppressed IFU Spectrograph Simon Ellis, Joss Bland-Hawthorn, Anthony Horton and Roger Haynes

1 The Need for Deep Near-infrared Spectroscopy The ability to achieve deep near infrared spectroscopy is of great importance to the future of astronomy. Our understanding of the early Universe depends on our ability to observe highly redshifted spectroscopic diagnostic features. Observations of H-α at wavelengths 0.9 ≤ λ ≤ 1.8 µm would allow accurate star-formation rates to be measured over the period 0.4 ≤ z ≤ 1.7—a crucial period in the formation of galaxies. Observations of Lyman-α at the same wavelengths would probe 6 ≤ z ≤ 13—the epoch at which the Universe underwent a major (and poorly understood) phase change from neutral to ionised. Similarly our understanding of low mass stars and unbound planets relies on NIR spectroscopy, since these objects emit most of their light at NIR wavelengths. Traditionally deep NIR observations have been hampered by noisy detectors and very bright sky backgrounds. Detector technology has now matured to the point where NIR arrays can deliver high quantum efficiency, low dark currents (≈ 0.002 e− s−1 ) and low read noise (≈ 8 e− for Fowler sampling read out) [9]. The sky however remains a problem. The H band night sky is 8 mags brighter than the V band. Furthermore the sky brightness fluctuates rapidly, both temporally and spatially, severely hindering accurate sky subtraction. The source of ≈ 99% of the J and H band background is line emission from rotationally excited hydroxyl radicals at ∼ 90 km altitude [3]. Whilst these OH lines are exceedingly bright, they are also intrinsically very narrow [3]. Thus, in principle there should exist very dark regions of sky between the lines, which if made accessible whilst simultaneously rejecting the OH lines could yield very deep NIR observations. The challenge is to exploit these regions whilst avoiding scattering from the OH lines and maintaining high throughput over the full J or H bands. S. Ellis () AAO, PO Box 296, Epping, NSW 1710, Australia e-mail: [email protected] A. Moorwood (ed.), Science with the VLT in the ELT Era, Astrophysics and Space Science Proceedings, doi: 10.1007/978-1-4020-9190-2_79, © Springer Science + Business Media B.V. 2009

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Fig. 1 The top plot shows the forest of OH lines responsible for the extremely bright NIR sky background (taken from [8]). The bottom plot shows that there exist dark regions between the lines which can in principle be exploited to achieve very deep NIR observations

2 Scattering and the Need for Photonic OH Suppression At first sight it would seem that the OH lines could be removed via high resolution spectroscopy and software subtraction. However two obstacles combine to confound such an approach: (i) the extreme brightness of the lines, (ii) instrumental scattering. The OH lines are so bright compared to the interline continuum that any significant scattering will flood the dark window with background light. This can be seen in Fig. 2 which shows the scattering-only PSF for a diffraction grating [10] convolved with 4 OH lines in the H band compared to the zodiacal scattered light, which is believed to be the dominant component of the interline continuum (see [3] for details). Even well between the lines the scattered OH light dominates the background, despite the lines being chosen to contain a large ‘dark window’, indeed this is the region in which the interline continuum of [7] was measured. To avoid the problem of scattering the OH lines must be suppressed prior to dispersion. Methods which suppress the light after dispersion will help to reduce the total intensity of the background light, however the will not help in reducing the background between the lines—arguably the most important region, since this is where scientific measurements will be made. We will now describe such a method, being pioneered at the Anglo-Australian Observatory and the University of Sydney, and show that it promises enormous benefit to near-infrared astronomy.

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Fig. 2 The scattered light from 4 OH lines in the H band (continuous line) compared to the zodiacal scattered light (dotted line). The scattered light dominates the background

3 Fibre Bragg Gratings We have pioneered a radical new approach to OH suppression employing fibre Bragg gratings [2]. Fibre Bragg gratings (FBGs) are optical fibres with a slowly varying refractive index in the core. As light propagates down the fibre it suffers Fresnel reflections due to the varying refractive index. By properly controlling the index variation it is possible to produce strong reflections at specific wavelengths. It has been shown that it is possible to produce a low loss filter defined by a series of irregularly spaced, narrow notches with suppression depths as high as 30 dB, whist maintaining high throughput outside of the notches (losses ≈ 0.03 per 100 nm wavelength coverage) [2]. Two steps are required in order for this technology to be useful for astronomical observations: the ability to suppress many lines, and the ability to work with multimode fibres. These two requirements are not easily compatible. FBGs achieve peak performance within a single-mode fibre (10 micron core), whereas astronomical fibres, with their larger core diameters, are highly multi-moded. Therefore a successful implementation of photonic OH suppression requires that the many modes of the input fibre be transformed into the same number of single modes acting in parallel. To this end, we developed the photonic lantern [1, 6] which decomposes the spatial modes of the input fibre into a parallel degenerate array of single-mode fibres.

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Identical gratings are printed into each of the fibres. Once the filtered light exits the fibre array, these are recombined by an identical taper back to a multi-mode output fibre. It has been demonstrated that fibres with as few as 5 modes can have high coupling efficiency, paving the way for AO fed FBGs in the near future [5]. The combination of these two innovative technologies means that the development of FBGs for astronomical purposes is now at a stage that is compatible with astronomical requirements. FBGs are now being developed to remove ∼ 150 lines in each of the J and H bands. In light of the discussion in Sect. 2, an important advantage of FBGs is that they can be incorporated into a spectrograph before the dispersive elements, thus the OH light is strongly rejected with very little scattering, offering for the first time the capability to achieve extremely deep NIR observations without the requirement to put a telescope and instrument into space.

4 Simulated Performance of Fibre Bragg Gratings We have thoroughly modelled an entire observing system from object to detector of an instrument incorporating fibre Bragg gratings. Our model includes all known losses and sources of background, temporal variation in the sky background and systematic wavelength calibration errors. For a complete description see [3]. This system results in a reduction of the total background by a factor of 42 and 17 in the H and J bands respectively, corresponding to 4 and 3 magnitudes. It is important to realise that this reduction applies across the entire spectrum including the interline regions. We emphasise that it is the magnitude of the interline reduction that is scientifically important, and that the advantage of FBGs over other methods of OH suppression is their ability to suppress the OH light before the light enters the spectrograph, and therefore before any significant scattering surfaces are encountered. Figure 3 shows an illustrative simulation of an H = 26 AB mag quasar at z = 11, observed at a resolution of R = 1000 for 70 hrs. For the system including fibre Bragg gratings, the Gunn-Peterson absorption trough [4] is clearly visible; for the identical system with no FBGs the signal is lost in the noise. Fibre Bragg gratings will allow spectroscopic identification of the sources which existed in the epoch of reionisation. Number counts, source-type and the ionisation fraction of the inter-galactic medium at this crucial period of the history of the Universe will all be feasible goals for a FBG fed spectrograph. In order to expedite the use of FBGs we envisage retrofitting an existing spectrograph with an OH suppression unit, FLEX (The First Light Explorer). Flex will consist of a 1 × 1 91 element IFU feeding FBGs optimised for the follow-up spectroscopy of very high redshift candidates. Such a system on the VLT (e.g. FLEX + NACO) would revolutionise NIR spectroscopy and maintain the competitive edge of 8 m class telescopes well into the ELT era.

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Fig. 3 A simulated spectrum of an H = 26 AB mag, z = 11 QSO. The exposure time was 70 hr, and the spectral resolution was R = 1000. The top panel shows a system with FBGs and the bottom panel shows an identical system without FBGs. The red lines indicate the true object spectrum. A colour version of this figure is available at dx.doi.org/10.1007/978-1-4020-9190-2_79

References 1. 2. 3. 4. 5. 6. 7.

J. Bland-Hawthorn, AAO Newsl. 108, 4 (2005) J. Bland-Hawthorn, M. Englund, G. Edvell, Opt. Express 12, 5902 (2004) S.C. Ellis, J. Bland-Hawthorn, Mon. Not. R. Astron. Soc. (2008), in press J.E. Gunn, B.A. Peterson, Astrophys. J. 142, 1633 (1965) A. Horton, J. Bland-Hawthorn, Opt. Express 15, 1443 (2007) S. Leon-Saval, T. Birks, J. Bland-Hawthorn, M. Englund, Opt. Lett. 30, 19 (2005) T. Maihara, F. Iwamuro, T. Yamashita, D.N.B. Hall, L.L. Cowie, A.T. Tokunaga, A. Pickles, Publ. Astron. Soc. Pac. 105, 940 (1993) 8. P. Rousselot, C. Lidman, J.-G. Cuby, G. Moreels, G. Monnet, Astron. Astrophys. 354, 1134 (2000) 9. R. Smith et al., in High Energy, Optical, and Infrared Detectors for Astronomy II, ed. by D.A. Dorn, A.D. Holland. Proc. SPIE, vol. 6276 (2006), p. 62760R 10. T.N. Woods, R.T. Wrigley III, G.J. Rottman, R.E. Haring, Appl. Opt. 33, 4273 (1994)

An N -Band Integral Field Spectrometer Survey Instrument for the VLT A.C.H. Glasse, D.M. Henry and D. Lee

We present the design of an instrument which is capable of providing a legacy survey of the dust mineralogy of the galactic interstellar medium. It will measure the N band spectra and images of a wide range of galactic targets, selected from the MSX and Akari missions, with an instantaneous field of view of 5 × 10 arcseconds2 . An integral field unit will sample this field at a spatial resolution of 0.2 arcseconds, with a grating spectrometer then dispersing the full 7 to 13 µm waveband onto a 1024 × 1024 pixel2 detector at a spectral resolving power of around 700. The absence of cryo-mechanisms and the utilisation of an existing design of integral field unit all contribute to the instrument having low development, build and operational costs.

1 Introduction The spectral region of good atmospheric transmission from 7 to 13 µm (the N -band), includes emission and absorption features which provide key diagnostic information on the chemistry and mineralogy of the circumstellar environment for all stages of stellar evolution, ranging from proto-planetary disks and YSOs through to AGB stars and planetary nebulae. For example, the crystallinity and chemical composition of the silicate dust component of the ISM [1] can be determined by spectroscopy of the broad feature centred at 9.7 µm, whilst the narrower PAH [4] bands at 7.7, 8.6, and 11.3 µm probe the carbonaceous content of the ISM contained in small grains. The need to chart the composition of gaseous and solid materials at the different stages of star formation has been identified as a key question for astronomy in the next decade in the recent ‘Astronet’ report [7]. Until now, the systematic study of the thermal infrared spectra of these objects has been restricted to specific classes, and samples numbering no more than a few dozen. The IRS instrument on the Spitzer Space Telescope has recently begun to show the benefit of increasing sample sizes to a hundred or so. For example, the study of T Tauri stars in Taurus [2] has allowed morphologically based sequences to be derived which give new insights into the generalised properties of circumstellar disks. A.C.H. Glasse () The UK Astronomy Technology Centre, Blackford Hill, Edinburgh EH9 3HJ, UK e-mail: [email protected] A. Moorwood (ed.), Science with the VLT in the ELT Era, Astrophysics and Space Science Proceedings, doi: 10.1007/978-1-4020-9190-2_80, © Springer Science + Business Media B.V. 2009

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We propose extending this work to build up a complete picture of the composition of the ISM in a wide range of environments in our galaxy, by carrying out a survey of several thousand objects. A dedicated integral field spectrometer will measure the full N -band spectrum in a single exposure at a fully sampled spectral resolving power of 700, sufficient both to study the detailed shape of the dust features and to measure the unresolved fine structure nebular emission lines with high sensitivity. The instantaneous field of view will be 5 by 10 arcseconds in extent, with spectra obtained over a fully spatially sampled grid, at 0.2 arcsecond intervals. By comparison with the delivered performance of existing mid-infrared instruments on 8 metre class telescopes, we estimate that the spectra will achieve a signal to noise ratio of 10 after 1 minute of integration on a point source with a 10 µm continuum flux of 0.1 Jansky. In the background limited noise regime that will be seen at the VLT, rebinning along the wavelength axis will allow broad band images of the target region to be reconstructed at sensitivities comparable to those expected for a dedicated imager (i.e. 5 milliJansky for 10 sigma in 1 minute). The targets for the proposed survey will be selected from existing space-borne photometric surveys. For example, the MSX (Midcourse Space Experiment) survey of the Galactic Plane [6], has been followed up by a campaign of multi-wavelength observations aimed at selecting massive YSOs [5]. The resulting set of some 3000 objects, categorised as MYSOs, HII regions, OH-IR stars and planetary nebulae, have positions known to much better accuracy than the 5 arcseconds of the original survey, and all have fluxes brighter than 0.1 Jansky. The recently completed Akari IRC survey [3] is expected to extend the MSX survey by having all-sky coverage down to a flux limit of 0.08 Jansky flux at 9 µm. Once again, the source positions will be measured with sufficient accuracy that they can be acquired within the field of view of our proposed VLT instrument with minimal need for raster searching or peaking up. The VLT is well suited to this project by virtue of its demonstrated capability for thermal infrared astronomy. Its large diameter aperture provides both the spatial resolution which is needed to separate the multiple sources which can confusion limit space-based observations, and the sensitivity required to complete a scientifically valuable survey on a reasonable timescale. We expect to be able to measure the spectra of several dozen sources selected from the above surveys in a single night.

1.1 Instrument Design The optical layout of the proposed instrument (excluding the spectrometer optics), is shown in Fig. 1. The all-reflective optics and their support structure will be made from aluminium and cooled to below 10 K using a closed cycle refrigerator. The mirrors will be of gold coated, diamond finished aluminium, in order to achieve high transmission (typically we expect only 1.5% transmission loss per surface). The choice of zinc selenide for the cryostat entrance window is made possible by the decision to restrict the wavelength coverage to include only the N -band. ZnSe

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Fig. 1 Optical layout from the cryostat window to the spectrometer input focal plane. Only five out of the full complement of twenty four slices are shown for clarity. A colour version of this figure is available at dx.doi.org/10.1007/978-1-4020-9190-2_80

is a relatively hard material, which is easy to polish and anti-reflection coat, and is lower in cost than other materials which transmit out to longer wavelengths. The N -band filter and pupil stop contribute to the instrument’s excellent control of stray light, along with several other locations along the optical path where field and pupil baffles can be positioned. We note that the filter is fixed. There are no moving parts in the proposed design. Next in the optical train come a series of mirrors which provide anamorphic magnification of the 5 × 10 arcsecond field and present it to the image slicer mirror. This slices the field into 24, 10 arcsecond long by 0.2 arcsecond wide strips, only five of which are shown in Fig. 1 for clarity. These strips are re-imaged onto the entrance focal plane of a grating spectrometer, which disperses the 7 to 13 µm waveband to form 24 stacked slit spectra on a 1024 by 1024 pixel detector. The optics shown in Fig. 1 are very closely modelled on the IFU which has been developed at the UKATC for the mid-infrared instrument “MIRI” [8]. MIRI is being built by a consortium of European and US Institutes to fly on the James Webb Space Telescope (JWST). In keeping with the exceptionally demanding performance requirements which are placed on such a key mission, a great deal of effort has been expended by the MIRI team in qualifying the IFU. We will be able to exploit this effort for the proposed VLT instrument, saving on design, development and manufacturing costs, and giving a high degree of confidence in achieving good image quality and throughput performance. A preliminary design of the spectrometer has been produced which uses a single collimating mirror and a three mirror anastigmat for the camera optics. It has successfully demonstrated that there are no technical obstacles in the way of achieving diffraction limited image quality across the full field at the detector. In practice, we

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Fig. 2 The JWST-MIRI Integral Field Unit. A colour version of this figure is available at dx.doi.org/10.1007/978-1-4020-9190-2_80

intend adding a fourth powered mirror to the camera in order to reduce the image distortion to a negligible level. Currently the detector is baselined to be a 1k × 1k pixel2 “Aquarius” Si:As IBC device, but any array with comparable size and format could be used. The exposure time per frame is estimated to be around 1 second, which places no stressing requirements on the array controller, or due to the data volume generated. We note that the entire optical train including the detector can be fitted comfortably within a volume of 0.5 × 1.0 × 1.5 metres3 .

1.2 Operations The use of an IFU makes the proposed instrument exceptionally efficient to operate, and ideally suited to its primary role as a survey tool. The automatic provision of a deep broad band image of the 5 × 10 arcsecond2 field around the target position, means that no peakup is required, with absolute source positions readily determined by accurate telescope offsetting from nearby astrometric standard stars. With only one optical configuration available, observations are fully defined by the choice of telescope chop and nod pattern, which would in turn be similar to those used by existing VLT thermal infrared instrumentation. The production of calibrated spectra from the exposures taken at these dither positions will be achieved using a spectral image processing pipeline which will take advantage of the algorithms currently under development for the analysis of data

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produced by the functionally very similar JWST-MIRI instrument. We anticipate that both the operational support and the off-line production of the reduced survey will follow the pattern of low cost which is anticipated for the instrument itself.

References 1. J. Bouwman, G. Meeus, A. de Koter, S. Hony, C. Dominik, L.B.F.M. Waters, Astron. Astrophys. 375, 336, 950 (2002) 2. E. Furlan, L. Hartmann, N. Calvet, P. D’Alessio, R. Franco-Hernández, W.J. Forrest, D.M. Watson, K.I. Uchida, B. Sargent, J.D. Green, L.D. Keller, T.L. Herter, Astrophys. J. Suppl. 165, 568 (2006) 3. D. Ishihara, T. Onaka, H. Kataza, T. Wada, H. Matsuhara, M. Ueno, N. Fujishiro, W. Kim, K. Uemizu, Y. Ohyama, S. Oyabu, Y. Ita, I. Sakon, H. Fujiwara, S. Hasegawa, T. Matsumoto, I. Yamamura, T. Tanabe, H. Murakami, H. Watarai, M. Cohen, Mem. Soc. Astron. Ital. 77, 1089 (2006) 4. A. Leger, J.L. Puget, Astron. Astrophys. 137, L5 (1984) 5. S.L. Lumsden, M.G. Hoare, R.D. Oudmaijer, D. Richards, Mon. Not. R. Astron. Soc. 336, 621 (2002) 6. S.D. Price, M.P. Egan, S.J. Carey, D.R. Mizuno, T.A. Kuchar, Astron. J. 121, 2819 (2001) 7. The Science Vision Working Group, in A Science Vision for European Astronomy, ed. by P.T. de Zeeuw, F.J. Molster (2007) 8. M. Wells, D. Lee, A. Oudenhuysen, P. Hastings, J.-W. Pel, A. Glasse, in Space Telescopes and Instrumentation. Proc. SPIE, vol. 6265 (2006), p. 14

High Resolution Visible Imaging on the VLT Craig Mackay

Abstract A new approach to achieving diffraction limited imaging in the visible from ground-based telescopes is described. It uses lucky image selection to choose the sharpest images from a sequence taken at high speed. When used behind a low order adaptive optics system it is possible to work with much larger diameter telescopes than would normally be expected. A further development involving aperture segmentation and image re-synthesis promises to allow these techniques to be extended to even larger telescopes.

1 Introduction The Very Large Telescope (VLT) of the European Southern Observatory is arguably the finest telescope ever built. Its theoretical optical performance is quite remarkable and the mirror quality makes it entirely capable of diffraction limited imaging in the visible. Even on an excellent observing site such as Paranal the atmosphere restricts resolution in the visible to typically somewhat less than one arcsecond in diameter. The diffraction limited resolution of the telescope in V-band (550 nm) is about 16 milliarcsec while in I-band (850 nm) it is about 25 milliarcsec, or a factor of about 40 higher than permitted by the atmosphere. At longer wavelengths the turbulent phase errors are smaller so better imaging performance may be obtained although the intrinsic resolution of the telescope is also worsened. Most efforts to overcome this have used adaptive optics techniques which are described in detail elsewhere in this volume. The success of AO systems has been mixed. Good results have been obtained in the near infrared and in particular Kband (2.2 microns) on several telescopes. For example, on the Keck telescope [1] delivers Strehl ratios as high as 60% with a bright natural guide stars, somewhat lower using laser guide stars. Similar results have been obtained on the VLT with the NAOS-CONICA system which also has a laser guided facility [2]. In K-band the resolution is comparable with that delivered by the Hubble Space Telescope in the visible. However no one has yet delivered Hubble resolution in the visible from ground-based Hubble-sized telescope using adaptive optics. The principal problem is that in the visible phase variations occur on such short length scales across the diameter of the telescope that many sensors and actuators are required in principle to deliver acceptable correction. The turbulent structure can C. Mackay () Institute of Astronomy, University of Cambridge, Cambridge, UK e-mail: [email protected] A. Moorwood (ed.), Science with the VLT in the ELT Era, Astrophysics and Space Science Proceedings, doi: 10.1007/978-1-4020-9190-2_81, © Springer Science + Business Media B.V. 2009

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change radically on very short timescales. The full width of a bright star image can change by a factor of two within a few tens of milliseconds. Atmospheric turbulence is significantly non-Kolmogorov [3]. Fried [4] suggested that provided there are not too many turbulent cells covering the aperture of the telescope then, using very short exposures, eventually individual frames will be obtained that are close to diffraction limited. We have developed a complete imaging system for lucky imaging that uses dedicated high-speed noiseless CCD cameras developed in Cambridge. This camera (LuckyCam) has been used on telescopes around the world to deliver Hubble resolution images in the visible routinely, something not achieved hitherto by any other imaging method. An additional benefit of the lucky imaging technique is that the size of the isoplanatic patch is greatest at those moments when the variance in phase across the telescope aperture is at its lowest. A typical lucky imaging isoplanatic patch diameter of about one minute of arc at 850 nm compares favourably with typical isoplanatic patch sizes reported in the visible by Sarazin and Tokovinin [5] equivalent to about 4–5 arcsec at the same wavelength. An additional advantage of the lucky imaging technique is that it is possible to use the smallest percentage to give the very highest resolution and a larger percentage to give better sensitivity at the expense of resolution. This trade-off may be done later simply by being exposures into bins of different quality. Many scientific results have already been published using these techniques. Lucky Imaging is very efficient in its use of telescope time. It requires a reference star brighter than about 16.5 magnitude in I band. With simple modifications of our existing camera (the use of deep depletion CCDs, and the use of additional out of band light recorded simultaneously with a separate detector) it is expected that reference stars as faint as 18 magnitude can be used, a magnitude that almost guarantees that will be a reference star within the one arc minute isoplanatic patch even at the highest galactic latitudes. Indeed it is the magnitude of the reference star for any of these of image correction techniques that is going to make or break it. Any approach that requires reference stars substantially brighter than this will be of more limited application.

2 Lucky Imaging on Larger Telescopes Resolution higher than that produced by the Hubble Space Telescope is only possible with a larger diameter telescope. Under typical conditions the probability of recording a near diffraction limited image on a Hubble sized telescope is acceptable, and in the range of 5–30%, but with larger telescopes the probability of getting good images rapidly becomes vanishingly small as the number of turbulent cells covering the aperture increases. Most of the turbulent power in the atmosphere is concentrated on the longer scales. In principle by removing these low order errors the net turbulent power will be substantially reduced allowing lucky imaging methods to work more successfully. We used the Palomar 5 m telescope which has a low order adaptive optics

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Fig. 1 Core of globular cluster M13 with natural seeing (0.65 arcsec) on left and with LuckyCam behind the PALMAO system (40 milliarcsec resolution). A colour version of this figure is available at dx.doi.org/10.1007/978-1-4020-9190-2_81

system (PALMAO), behind which we mounted the lucky camera (LuckyCam). Under conditions of 0.6 arcsec seeing we recorded images with about 40 milliarcsec resolution and Strehl ratio of about 17% in I band (Fig. 1). This image is the highest resolution image ever taken in the visible or in the near infrared on a telescope on the ground or indeed in space. We have also images taken in the V-band which also show excellent resolution with nebular images. The disadvantage of using an imaging system behind an adaptive optics unit is that most AO units require relatively bright reference stars to work. This is particularly acute for high order systems where the aperture of the telescope is divided into a large number of relatively small cells that have to be read out much faster than the typical de-correlation time of a few milliseconds. The requirement to get a satisfactory signal-to-noise out of each element of Shack Hartmann sensor is that there needs to be of the order of a hundred detected photons per read and these photons may only be gathered from a patch of the image as small as 25 cm in diameter. It is not uncommon for adaptive optics systems to be run at many hundreds of hertz. This immediately demonstrates why lucky imaging is able to use such faint reference stars. With a 2.5 m diameter telescope we have 100 times the photon gathering capacity and by using a frame rate of a few tens of frames rather than a few hundreds of frames per second we have an overall sensitivity for a reference object that can be as much as 1000 times better than the typical AO system. By using a lucky imaging camera behind a typical AO system we lose this enormous advantage.

3 High Resolution on the Largest Telescopes Having an AO system in front of a lucky camera seriously limits guide star sensitivity. The great majority of adaptive optics systems use the Shack Hartmann sensor

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approach. An alternate approach is to use a curvature sensor. A comparison of realised adaptive optics systems using both Shack Hartmann sensors and curvature sensors has been made by Racine [6] which explains that although curvature sensors can be difficult to build they prove to be particularly effective in delivering high efficiency correction systems. The efficiency with which curvature sensors may be constructed has been greatly improved by the availability of photon counting back illuminated CCD detectors. They are ideal as curvature sensors because the selection of sub aperture geometry may be handled in software rather than hardware. Compact high resolution MEMS-type flexible mirrors such as those manufactured by Boston Micro-Machines allow the design of very simple, compact systems free from many of the complexities traditionally associated with curvature sensors. The combination of a low order adaptive optics system and a lucky imaging camera is unlikely to work with the next generation of very large telescopes, with diameters > 10 m. A different approach must be used if we are to maintain our luck in our image selection strategies. This can be done if the aperture of the telescope is broken into sub apertures as in Fig. 2. The light from sets of four sub apertures is directed to one detector and other sets of four sub apertures directed to other detectors. Each then sees the image of the sky as it would appear with a single sub aperture with compact objects crossed by Michaelson fringes arising from the interference between sub apertures (Fig. 3). The low order AO system will ensure that the light from a compact object from each sub aperture would be overlaid. Residual turbulence will cause the fringe visibilities to be variable. Images may be selected for the highest visibility. When three or more sub apertures are used together then the phases of the fringe patterns may be measured and closure phase methods used to combine data to allow a near diffraction limited image to be synthesised. This technique was demonstrated on the Palomar 200 inch telescope nearly 20 years ago [7] giving 65 milliarcsec resolution images. In those days the limitations of available photon counting detectors and computer processing power made the technique impractical for general use. It is clear from the

Fig. 2

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work reported here that this is no longer the case and that we can now expect to obtain near diffraction limited images in the visible from current large ground-based telescopes as well as even high resolution from the next generation of extremely large telescopes.

References 1. 2. 3. 4. 5. 6. 7.

Van Dam et al., Keck AO Note 489 (September 2007) ESO User Manual, version 81.1 (December 2007) Tubbs, Proc. SPIE 6272, 93 (2006) Fried, J. Opt. Soc. Am., A, Opt. Image Sci. 65, 1372 (1978) Sarazin, Tokovinin, in Conference Beyond Conventional AO, Venice, 2001 R. Racine, Publ. Astron. Soc. Pac. 118, 1066 (2006) Nakajima et al., Astron. J. 97, 1510 (1989)

Life on the Fast Lane: The Burst Mode at the VLT at Present and in the Future Andrea Richichi, Octavi Fors, Elena Mason, Marco Delbó, Jörg Stegmaier and Gert Finger

Abstract The recent implementation of the high-speed burst mode at the ISAAC instrument on UT1, and its propagation to other ESO instruments, has opened the door to observational capabilities which hold the potential for a wealth of novel results. In the ELT era, when the accent will likely be on lengthy programs aimed at the best sensitivity and angular resolution, the VLT telescopes could continue to play a significant and largely unique role by performing routinely observations of transient events at high temporal resolution. In our contribution, we provide details on two such kinds of observations, namely lunar occultations of stars and of asteroids. For the first ones, we report on two passages of the Moon in regions with high stellar density as the Galactic Center. The VLT-UT1 telescope was used for the first time to record successfully 53 and 71 occultations on March 22 and August 6, 2006, with an angular resolution of 0.5–1 milliarcsecond and K ∼ 12.5 limiting magnitude. We note that the angular resolution is superior to that achieved at present by Adaptive Optics on any telescope, and also superior to that foreseen for the ELT at the same wavelength. LO are also very efficient in terms of telescope time. We present some of the results, including the discovery of close binaries, and the detection and study of compact circumstellar components of cool giants, AGB stars and embedded IR sources. Concerning asteroidal occultations, we aim at observations starting in P80 which would permit high-accuracy, direct determinations of asteroid sizes for bodies larger than ≈ 50 km. This is a critical information to improve our understanding of the physical properties of these bodies. It will allow us an independent, crucial calibration of the indirect techniques commonly used to derive estimates of asteroid sizes and albedo, namely radiometry (Harris and Lagerros [Asteroids III, ed. by W.F. Bottke, A. Cellino, P. Paolicchi, R.P. Binzel (Univ. of Arizona Press, Tucson, 2002), p. 205]) and polarimetry (Cellino et al. [Icarus 179, 304 (2005)], and references therein). Lunar occultations can be used also to detect asteroid binary systems, which have been found recently to be not very rare. Binary systems are invaluable to estimate asteroid masses and densities, parameters that are at present very poorly known. A. Richichi () European Southern Observatory, Garching, Germany e-mail: [email protected] A. Moorwood (ed.), Science with the VLT in the ELT Era, Astrophysics and Space Science Proceedings, doi: 10.1007/978-1-4020-9190-2_82, © Springer Science + Business Media B.V. 2009

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1 Scientific Motivations Our main scientific driver are lunar occultations of stellar sources, and we provide here some details on the goal and the technical requirements. Later we will mention other astrophysical topics which also require fast, or very fast, photometry. Lunar occultations (LO) represent a very powerful and yet extremely simple method to obtain high angular resolution information on sources covered by the Moon during its apparent motion as seen from a specific site. The obvious disadvantages of LO are that we cannot choose the targets at will, and that they are fixed time events. Figure 1 illustrates the basic geometry of an occultation event. The position angle (PA) defines the scan direction of the lunar limb across the source. The local limb slope ψ can influence this by a few degrees. The contact angle CA defines the effective speed VP of the limb motion, which can be significantly slower than the apparent lunar motion VM , especially for LO approaching grazing conditions. The parameter VP , in conjunction with the observing wavelength, defines the rate of the diffraction fringes. The model light curve of Fig. 1 has been generated for a typical speed of 0.5 m/s, in a broad-band K filter. The detailed analysis of the diffraction light curve of a LO is outside the scope of this contribution, but we mention that it is rather simple and relatively quick. It can recover information on the occulted source with angular resolutions of ≈ 1 mas, and can even be used to reconstruct brightness profiles of complex sources without any model-dependent parameters. These characteristics make LO competitive with other high-angular resolution methods such as AO or long-baseline interferometry. However, the light curve must be sampled at millisecond rates. This was easily achieved with simple photometers: single pixel InSb detectors have been routinely used for this work in the near-IR. On the other hand, these photometers presented the disadvantage of integrating the light over the whole diaphragm, thus resulting in added noise from the rather high background normally encountered during LO observations. The progressive introduction of array detectors has significantly changed the landscape of near-IR instrumentation. Nowadays, almost all telescopes of medium

Fig. 1 Left: Scheme of a lunar occultation event. Symbols are explained in the text. Right: A typical (near-IR) model occultation light curve. Note the temporal scale of the diffraction fringes, which requires high speed photometry for proper sampling

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to large size are equipped with such detectors, which emphasize area and read-out noise (RON) but, with the possible exception of specialized detectors for wavefront sensing, are usually slow to read-out. However, suitable compromises can be realized by reading out fast only parts of the arrays. A comparison of the two approaches was made by Richichi [7], and we refer to this paper for an analytic and quantitative comparison of the signal-to-noise ratio (SNR) for the photometer and the array cases under various combinations of telescope size, source brightness and background intensity. The comparison included 4 m and 8 m telescopes: while these latter might have seemed a remote possibility ten years ago, LO have now been observed on one of the 8.2 m VLT telescopes equipped with the ISAAC instrument in the so-called burst mode (Richichi et al. [8]). Presently, the burst mode is limited to sampling times of 3–5 ms. This is suitable for LO work, as well as for other applications of moderately fast photometry such as occultation and transit phenomena in the solar system, transient and oscillatory phenomena in stars and binary stars, and some aspects of the highly time variable properties of white dwarfs, neutron stars and pulsars—where of course sensitivity is the main barrier. A list of high-time resolution astrophysical topics and available instruments has recently been presented by Redfern and Ryan [5].

1.1 The Burst Mode of ISAAC A fast readout mode on a user-defined subarray has recently been implemented, tested and commissioned on the Aladdin detector of the ISAAC instrument [8]. This burst mode is optimized for speed, at the expense of data organization which needs to be performed offline. The raw data consist in a sequential write of successive reads of the subarray, i.e. they are double the size of the standard FITS cubes and do not have a FITS structure. A similar mode, so-called Fastjitter, can achieve standard data organization but is several times slower. The main characteristics of the burst mode are listed in Table 1. We have used the burst mode on the occasion of two series of lunar occultations in crowded regions near the galactic center, on March 22 and August 6, 2006. In particular for the second event, we were able to observe 71 occultations over a period of few hours. For experienced observers, it is possible to offset the telescope, start the observation, record the data in a total of about two minutes. Preliminary Table 1 Main parameters of the ISAAC burst mode Window size (px)

Field of view ( )

Min DIT (ms)

Max NDIT

32 × 32

4.7 × 4.7

3.2

16 000

64 × 64

9.5 × 9.5

6.4

16 000

128 × 128

19 × 19

14

4000

256 × 256

38 × 38

37

1000

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results from these runs have been reported by Fors et al. [2]. Among them, we mention 7 small separation binaries, 5 angular diameters, and 4 sources with extended circumstellar emission. An example is shown in Fig. 2. The study by Richichi [7] predicted that an 8 m telescope equipped with a nearIR array detector would reach between K = 12 and 14 mag, depending on the lunar phase and background, with an integration time of 12 ms at SNR = 10. The preliminary results with ISAAC reported by Richichi et al. [8] show a limiting magnitude K ≈ 12.5 at SNR = 1 and 3 ms integration time. Within the uncertainty of the lunar background, the results are in perfect agreement with the decade-old prediction. Thanks also to the introduction of massive predictions based on IR survey catalogues such as 2MASS and of automated data pipelines based on new methods of light curve characterization based on the wavelet theory (Fors et al. [3]) it is now possible to imagine programs of routine LO observations also at the VLT. We have prepared one such program as a filler for brief unused telescope times, which has been approved for P80 and has been resubmitted for P81.

Fig. 2 An anonymous star observed on August 6, 2006 (2MASS 17524903-2822586), shown to be a small separation, low contrast binary. The main panel shows the data, the best fit with a binary star, and two sets of fit residuals: one for a single source and one for the binary model. The residuals are offset by arbitrary amounts and enlarged for clarity. The inset is a profile reconstruction using a model independent method (Richichi [6]). The companion has a projected separation of 12.8 ± 2 mas, and a brightness ratio of 1 : 35 (magnitudes of the two components are K = 5.23 and K = 9.12, respectively). A colour version of this figure is available at dx.doi.org/10.1007/978-1-4020-9190-2_82

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1.2 Asteroid Occultations The interest in LO of asteroids stems from the possibility to derive the sizes of these bodies from the duration of the occultation event. We also note that for objects with angular sizes of ≈ 20 mas or larger, the LO light curves can be analyzed by simple geometrical (rather than diffraction) optics. In these cases, details of the brightness profile, connected with the distribution of albedo and irregularities on the surface, can be measured. Moreover, the technique is ideally suited to discover new binary systems. Asteroids are moving objects and the prediction of the epoch and geometry of their lunar occultations, along with the visibility of the events from Paranal, must be calculated for each of object. The K ≈ 12.5 limiting magnitude of the high-speed burst mode of ISAAC on the UT1, roughly scales to a limiting asteroid magnitude of V ≈ 14.5. Using a computer code to predict the lunar occultations involving the first 5000 numbered asteroids (asteroids with number > 5000 are in general likely to have V > 14.5), we find that there are about 40 lunar occultations of asteroids observable from Paranal every year. The unique combination of ISAAC and UT1 makes it possible to observe occultations events of asteroids with diameters in the range between 100 and 50 km in the Main Belt and obtain direct determination of their sizes. In this size range, asteroids display angular extensions between 50 and 100 mas. At the average lunar proper motion their disappearance last 100 ms (minimum). By tuning the integration time to yield SNR > 10 the sizes of these bodies can be determined with the unprecedented accuracy of some percent.

References 1. A. Cellino, R.G. Hutton, M. di Martino, P. Bendjoya, I. Belskaya, E.F. Tedesco, Icarus 179, 304 (2005) 2. O. Fors, A. Richichi, E. Mason, J. Stegmaier, T. Chandrasekhar, in Highlights of Spanish Astrophysics IV, ed. by Figueras et al. (2007) 3. O. Fors, A. Richichi, X. Otazu, J. Nuñez, Astron. Astrophys., submitted, arXiv:0711.0537 (2007) 4. A.W. Harris, J.S.V. Lagerros, Asteroids III, ed. by W.F. Bottke, A. Cellino, P. Paolicchi, R.P. Binzel (Univ. of Arizona Press, Tucson, 2002) p. 205 5. M. Redfern, O. Ryan, in Towards the European ELT, Marseille (2006) 6. A. Richichi, Astron. Astrophys. 226, 366 (1989) 7. A. Richichi, in Very High Angular Resolution Imaging. IAU Symp. 158 (1997) p. 71 8. A. Richichi, O. Fors, E. Mason, J. Stegmaier, Messenger 126, 24 (2006)

High Resolution Near Infrared Spectroscopy: Prospects for 10 and 40 m Class Telescopes E. Oliva and L. Origlia

Abstract The recent development of high sensitivity near infrared (NIR) spectrometers with high resolution (HR) capabilities is opening new windows in our understanding of several hot topics of modern astrophysics. This field is one of the youngest branches of astronomical research. Only few dedicated instruments are available worldwide, but with remarkably limited capabilities in terms of spectral coverage. In this paper we discuss some of the most important scientific applications of HR-NIR spectroscopy and present the design of a simple and very powerful instrument which can be easily adapted to existing 8–10 m class telescopes (e.g. VLT) and to future, larger facilities (e.g. ELT). The expected performances and limiting magnitudes are analyzed in quite some details.

1 Introduction High resolution near infrared spectroscopy is a relatively new branch of observational astronomy, with a huge scientific potential. Indeed, the 0.9–2.5 µm spectral range has a wealth of diagnostic absorption and emission lines, which are crucial for a thorough understanding of several hot topics of modern astrophysics. The much reduced extinction at these wavelengths also allows NIR spectrometers to pierce the dust embedding several Galactic and extragalactic objects, which are heavily obscured in the optical. Moreover, at high redshift several emission and absorption spectral features, commonly exploited when studying local galaxies, are shifted into the NIR. Existing instruments on large telescopes suffer by a major and unresolvable limitation, namely the narrow spectral coverage in one shot (see Table 1). This implies that extremely long overall integrations are needed to obtain reasonable complete spectra. These instruments are thus suitable for “few-lines” science, but inefficient for a proper quantitative spectroscopy aimed for example at obtaining a complete screening of chemical abundances in cool stars or extremely accurate radial velocity measurements for extra-solar planet search. A wide spectral coverage is also E. Oliva () INAF – Osservatorio di Arcetri, largo E. Fermi 5, 50125 Firenze, Italy e-mail: [email protected] E. Oliva INAF – Telescopio Nazionale Galileo, Roque de Los Muchachos Observatory, 38700 La Palma, Spain A. Moorwood (ed.), Science with the VLT in the ELT Era, Astrophysics and Space Science Proceedings, doi: 10.1007/978-1-4020-9190-2_83, © Springer Science + Business Media B.V. 2009

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Table 1 List of all the HR-IR spectrometers in the world Instrument

Telescope

Rmax

Slit width for Rmax

Spec. coverage in single frame

CRIRES

VLT (8 m) + AOa

1 · 105

0.2

λ/70

PHOENIX

Gemini (8 m)

8 · 104

0.2

λ/200

GIANOb

TNG (3.5 m)

5 · 104

0.5

λ/1.3

NIRSPEC

Keck (10 m)

3 · 104

0.3

λ/10

CSHELL

IRTF (3 m)

3 · 104

0.5

λ/240

a Coupled

to a dedicated adaptive optics system

b commissioning

at the telescope foreseen in 2008

crucial to unveil the nature and physical properties of poorly explored objects, like e.g. very low mass dwarfs or transient objects.

2 Science Cases We identify three major astrophysical subjects where the unique combination of high spectral resolution and wide spectral coverage in the near NIR is crucial. (i) The search for extrasolar planets is one of the new, rapid developing field of research [1]. Most planet candidates have been detected by measuring the variation of the stellar radial velocity due to the orbiting body [7]. For a given planet mass, the amplitude of the Doppler shift is larger in the lowest mass stars, which however, are intrinsically very red and too faint in the visual range. High-resolution spectroscopy in the NIR domain is thus crucial to detect rocky planets down to a few Earth masses around 0.1 M and possibly with habitable conditions. Wide spectral coverage and high stability are crucial to measure radial velocity with ≈ 1 m/s accuracy. Existing 4–8 m class telescopes can allow the detection of such systems within ≈ 200 pc. An ELT will allow to push the search as far as ≈ 1 kpc. (ii) The most suitable laboratory to study stellar physics, evolution and chemical enrichment in the high metallicity domain is the inner Galaxy, where dust obscuration makes optical observations virtually impossible. HR-NIR spectroscopy over a wide spectral range is thus crucial for a systematic study of the chemistry and kinematics of metal rich stellar populations [3, 4]. Existing 4–8 m class telescopes allow to study the bright, evolved stellar populations of red giants and super-giants. However, pristine abundances, which are best recorded in non evolved (and much fainter) main sequence stars, can only be measured with an ELT. (iii) Understanding the metallicity evolution of galaxies is one of the most important and challenging issues of modern astronomy. While in the local Universe the metallicity has been studied in detail for several systems, at high-z, where

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the metallicity evolution should be strongest and cosmologically more important, there are only a few studies, often plagued by large uncertainties [5]. Damped Lyα and Lyα Forest systems are especially interesting since associated to proto-galaxies lying along the QSO or GRB line of sight [6, 8]. At z > 4, these systems can be probed only by means of HR-NIR spectroscopy. Existing 4–8 m class telescopes allow only to reach the brightest ones. An ELT is thus crucial to explore metals in low density Lyα systems, an making a breakthrough in our understanding of the metallicity and dust evolution in the early Universe. Finally, we mention that HR-NIR spectroscopy is also a fundamental tool to study several other relevant topics in astrophysics, namely stellar magnetic fields and winds, star-forming regions, brown dwarfs, obscured super star clusters in starburst galaxies, and black holes in obscured AGNs.

3 Expected Performances The expected performances can be quantified in terms of the limiting magnitudes achievable in the various photometric bands for a given combination of integration time, telescope and instrumental parameters. To properly evaluate their effects it is convenient to express the limiting magnitudes as analytical functions of the various parameters. Our modeling procedure operates as follows. The basic parameters of the instruments, which are maintained constant in the computations, are (see Sect. 4) as follows: – the projected slit width on the array: 2 pixels, – the product between pixel area and aperture of the beam on the array: A · = 0.55 m2 arcsec2 . This translates into a sky-project pixel size of 0.19 , 0.09 and 0.019 for a 4 m, 8 m and 40 m telescope, respectively, – the resolving power with a 2-pixels slit: λ/λ = 1 · 105 . We assume that the spectrograph will be fed by means of an optical system (e.g. adaptive optics, image slicer) which concentrates a fraction SLE of the light within the spectrometer slit. This parameter is treated as a free parameter independent from the telescope diameter. We adopt a reference value of SLE = 0.5, which is appropriate for a 4 m-class telescope operating with 0.9 natural seeing or a larger telescope equipped with moderate adaptive-optics correction combined with a 1 × 2 image slicer. The spectrum is extracted summing NY pixels along the slit. This parameter takes also into account the slicing-factor. Typical values are NY = 3 (no slicing) and NY = 10 (1 × 2 slicing). We adopt a reference value of NY = 6. The telescope and instrument have an overall efficiency η. This parameter includes the array quantum efficiency and all the light losses due to atmospheric absorption, optical reflections (mirrors), transmissions (lenses/prisms) and diffraction (grating). It does not include the slit-efficiency which is given, instead, by the parameter SLE above. We adopt a reference value of η = 0.15.

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The detector is read-out and reset every DIT sec. Each frame is affected by a noise of DN electrons. This noise includes the read-out noise, the dark-current and whatever other source of noise which is generated by the detector independently of the amount of light impinging on the array. Typical values of these parameters for present HgCdTe detectors are DIT = 15–30 min and DN = 10 el. We assume a value of DIT = 30 min and a reference noise of DN = 10 el. The integration lasts a total of Thr hours. This parameter must be an integer multiple of DIT. We adopt a reference value of Thr = 1 hr. The extracted spectrum is re-binned (smoothed) along the dispersion by NX pixels. We adopt a reference value of NX = 1 (i.e. no re-binning). The sky background is assumed to be μref (Y) = 17, μref (J) = 17, μref (H) = 16 and μref (K) = 13 mag/arcsec2 . The dependence of limiting magnitude upon variations of the sky brightness is expressed by the parameter μsky = μref − μsky . We adopt a reference value of μsky =0. The above parameters influence the limiting magnitudes as follows. The limiting flux for a given S/N ratio is always proportional to the inverse of the slit-losses (SLE) and of the telescope area, i.e. mlim = const. + 2.5 log(SLE) + 5.0 log(DTel ). Note that this relationship has general validity because of the way in which the parameter SLE is defined. The other parameters have different effects depending on the relative importance of the array-noise, the sky-background (bck) noise, and the Poisson noise of the photons from the object (obj-noise). In particular: mlim = const. + 2.50 log(Thr · NX)

if obj noise array + bck noise,

mlim = const. + 1.25 log(Thr · NX)

if obj noise array + bck noise,

mlim = const. + 0.00 log(NY)

if array-noise bck + obj noise,

mlim = const. + 1.25 log(NY)

if array-noise bck + obj noise,

mlim = const. + 1.25 log(η) − 0.5 · μsky

if bck-noise array + obj noise,

mlim = const. + 2.50 log(η) − 0.0 · μsky

if bck-noise array + obj noise.

In practice, the situation is always in between the limiting cases just depicted. However, handy relationships can be derived by numerically determining limiting magnitudes and applying a simple interpolation on the data. The resulting fitting formula, which is accurate to within a few tenths of magnitude for most circumstances, is:

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Table 2 Limiting magnitudes and coefficients describing their dependencies on the instrumental parameters (see Sect. 3) Band

S/N

mref

α(η)

α(DN)

α(Thr )

α(NY)

α(sky )

Y

10

19.6

2.2

−1.7

1.5

−1.0

0.10

J

10

19.3

2.2

−1.7

1.5

−1.0

0.10

H

10

18.8

2.2

−1.7

1.5

−1.0

0.10

K

10

18.1

1.9

−1.1

1.5

−1.0

0.20

Y

30

18.0

2.2

−1.5

1.7

−0.8

0.10 0.10

J

30

17.8

2.2

−1.5

1.7

−0.8

H

30

17.3

2.2

−1.5

1.7

−0.8

0.10

K

30

16.6

1.9

−0.9

1.6

−0.9

0.20

Y

100

16.1

2.3

−1.0

1.9

−0.6

0.05

J

100

15.9

2.3

−1.0

1.9

−0.6

0.05

H

100

15.4

2.3

−1.0

1.9

−0.6

0.05

K

100

14.8

2.1

−0.6

1.8

−0.7

0.15

DTel SLE η + 2.5 · log + α(η) · log 42 m 0.5 0.15 DN NY + α(DN) · log + α(Thr ) · log(NX · Thr ) + α(NY) · log 10 el 6

mlim = mref + 5.0 · log

+ α(sky ) · μsky . Table 2 lists the values of mref and of the α coefficients appropriate for different signal to noise ratios and photometric bands.

4 The Spectrometer The design of the instrument was developed based on the experience that we recently gained with the GIANO-TNG instrument [2]. The spectrometer delivers a complete 0.95–2.5 µm spectrum on an arrays-mosaic with a total of 6K × 2K pixels. The spectral format, typical of echelle cross-dispersed spectra, is displayed in Fig. 1. The resolution of R = 105 is achieved with a 2 pixels wide slit whose sky-projected size is 0.19 or 0.04 for a 8 m or 40 m telescope, respectively. In practice, it has the same slit-resolution product (Rθ ) of CRIRES-VLT, which is also the maximum Rθ which can be achieved with a single (non-mosaic) grating. To increase the fraction of light passing through the slit we envisage a simple 1 × 2 slicer (e.g. a BowenWallraven prism). The spectrometer layout and ray-tracing is displayed in Fig. 2. In the design we gave particular attention to the practical feasibility, mechanical mounting and alignment of all the elements included in the instrument. All the elements are at

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Fig. 1 Simulated HR-IR spectrum as seen by the 6K × 2K arrays mosaic. The echellogram covers the complete spectral range in a single exposure. The slit parameters and limiting magnitudes are summarized in the upper inserts. A colour version of this figure is available at dx.doi.org/10.1007/978-1-4020-9190-2_83

cryogenic temperature and fixed, i.e. there are no moving parts. The collimated beam is 180 mm and the grating is a commercial R2 echelle with 23.2 lines/mm. The grating works in a quasi-Littrow configuration, with a small (1.4◦ ) off-axis in the direction perpendicular to dispersion. The core of the instrument consists of a three-mirrors anastigmatic (TMA) which acts both as collimator and camera. The entrance and exit beams are F/4.9, i.e. the focal length of the TMA is 882 mm. The off-axis aluminum mirrors are simple conical surfaces with maximum diameters well below 600 mm. Therefore, they can be easily manufactured using standard diamond-turning machines. Cross dispersion is performed by means of 3 prisms in double pass. The first prism (CRD1) is made of infrasil, while the others are made of ZnSe. This combination of materials is the same used for GIANO-TNG and XSHOOTER. It is the only one which practically guarantees the feasibility of large enough prisms. The detector is a mosaic of three 20242 Hawaii-II RG arrays with 18 µm pixels. The arrays are butted in the direction perpendicular to the plane of Fig. 2. The spectrometer can also be adapted to a mosaic of two 40482 arrays with 15 µm pixels, a format which may soon become the standard for HgCdTe arrays. Given the trapezium-like distribution of the light on the focal plane (see Fig. 1)

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Fig. 2 Layout of the cryogenic HR-NIR spectrograph with ray-tracing. The rays towards the grating are marked in green, while the blue and red lines show the path of the dispersed light at 0.95 and 2.5 µm. The bottom-left insert shows the spots diagrams taken at the center and edges of several orders. The numbers on top of the squares (18 µm, the pixel size) are the pixel ensquared energy (in %). A colour version of this figure is available at dx.doi.org/10.1007/978-1-4020-9190-2_83

some array quadrants are not illuminated. Therefore, one could conveniently select arrays with poor performances in these quadrants. The entrance slit is conveniently positioned on a very slow (F/60, say) beam, which simplifies the construction and alignment of the image slicer. After the slit the light passes through FL, a very weak (and with non-critical alignment) fieldlens which re-images the pupil image close to the grating. A simple elliptical mirror (FR), fed through a flat mirror (PM1), acts as focal reduced and creates a F/4.9 slit image about 6 mm on the side of the array. The light enters into the TMA collimator through a flat mirror (PM2) and reaches the grating after having been crossdispersed once. The light then returns to the prisms for a second cross-dispersion and re-enters the TMA system which focuses the echellogram onto the array with superb image quality (see the spot-diagrams in Fig. 2).

References 1. 2. 3. 4. 5.

G.W. Marcy, R.P. Butler, Annu. Rev. Astron. Astrophys. 36, 57 (1998) E. Oliva, L. Origlia, C. Baffa et al., SPIE 6269, 41 (2006) L. Origlia, R.M. Rich, S. Castro, Astron. J. 123, 1559 (2002) R.M. Rich, L. Origlia, E. Valenti, Astrophys. J. 665, 119 (2007) M. Pettini, in Cosmochemistry. The Melting Pot of the Elements, ed. by Esteban et al. (Cambridge University Press, 2004), p. 257

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6. J.X. Prochaska, H.-W. Chen, M. Dessauges-Zavadsky, J.S. Bloom, Astrophys. J. 666, 267 (2007) 7. G.A.H. Walker, Astrophys. Space Sci. 223, 103 (1995) 8. A.M. Wolfe, E. Gawiser, J.X. Prochaska, Annu. Rev. Astron. Astrophys. 43, 861 (2005)

Prospects and Needs of Micro-arcsecond Astrometry Andreas Seifahrt, Tristan Röll and Ralph Neuhäuser

Dedicated instruments for the VLTI, such as PRIMA [2, 5] or the proposed second generation VLTI instrument GRAVITY [4], are designed for phase referenced imaging and aim in providing an astrometric accuracy of about 10 µas for two targets within the isoplanatic angle (PRIMA) or within the 2 arcsec FOV of the VLTI (GRAVITY). While interferometric measurements provide highly accurate astrometry, they are also time consuming and technically challenging. In the niche where precise but not necessarily accurate astrometry is needed, such as for measuring the reflex motion of stars orbited by extrasolar planets, single aperture measurements using AO imagers could provide an alternative. Providing astrometric precisions in the order of 100–200 µas, such measurements are feasible with current AO imagers and are a much easier and cheaper way to confirm and detect extrasolar planets. This article focuses on the possibilities and limitations of this technique, outlines the difference to interferometric measurements and shows its possible impact on current and future VLT instrumentation.

1 Accuracy vs. Precision Interferometric measurements yield accurate astrometric information since distances in the sky (i.e. angles) are measured as differential optical path differences on the ground. Hence, the measurement is directly related to a physical distance that can (and has to) be measured accurately. Achieving µas accuracy on the sky translates into nanometre accuracy in the determination of the optical path differences of the interferometer. Single aperture measurements, i.e. measurements of the separation between two or more targets in conventional images, depend on the accuracy and precision of A. Seifahrt () Institut für Astrophysik, Universität Göttingen, Friedrich-Hund-Platz 1, 37077 Göttingen, Germany e-mail: [email protected] A. Seifahrt Astrophysikalisches Institut und Universitäts-Sternwarte, Universität Jena, Schillergässchen 2, 07745 Jena, Germany A. Moorwood (ed.), Science with the VLT in the ELT Era, Astrophysics and Space Science Proceedings, doi: 10.1007/978-1-4020-9190-2_84, © Springer Science + Business Media B.V. 2009

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the plate scale of the telescope and the instrument. In conjunction with the active optics system of the VLT and the adaptive optics systems of the imager (needed to suppress atmospheric turbulence) the plate scale can effectively only be determined on astronomical objects. The choice of such objects is, however, severely limited when aiming for sub-milli-arcsec accuracy and precision. Accurate astrometric coordinates of double or multiple stellar systems that fit into the FOV of an AO imager (i.e. smaller than the isoplanatic angle) are scarce and basically limited to interferometric measurements, such as masers observed with the VLBI. Such measurements offer a final accuracy of about 1 mas when taking the typical uncertainties in proper motion into account. Hence, the accuracy in determining the plate scale of an imager is usually not sufficient to warrant accurate astrometric measurements far below the milli-arcsec level. However, if the accurate position (or separation) of the target(s) is not important but only a precise measurement is needed to monitor small changes in the separation of two (or more) objects, the problem is reduced to keeping the plate scale constant at the needed precision, instead of determining its accurate value. As outlined in Seifahrt et al. [9], the dense cores of globular clusters are offering a rather high number of bright sources with an intrinsically high astrometric stability. Object densities are sufficient to fill the typical FOV of an AO imager with more than a dozend bright sources. Such systems are ideal for monitoring the stability of the plate scale by measuring the separation of various pairs of stars in the chosen field. The final precision is than given as the intrinsic variation of the separation measurements divided by the square root of independent measurements. The later is the product of the number of images times the number of stars in the FOV.

2 Breaking the Milli-arcsec Barrier PSF centroids can be determined with high precision in the optical from individual frames [6]. In the near infrared (the wavelength regime of AO imagers), the typical precision in obtaining a stellar position in a well sampled and high-S/N frame taken with an AO imager is of the order of 1 mas. This is, however, far below the expected precision from photon noise. A further improvement is thus only possible on a statistical basis. Obtaining a high number of observations, thus performing statistically independent measurements, improves the final precision when the individual values (e.g. the measured separation of a binary in many frames) are normally distributed. This technique resembles the basic principle of determining high precision radial velocities (RV). While the typical pixel scale of a high resolution spectrograph is matched to its spectral resolution and is thus in the order of 1–2 km/s, the final precision needed to detect e.g. extrasolar planets is of the order of a few m/s. Hence, the equivalent centroiding precision of the spectral lines would be about 1/1000 pixel. Only the measurement of many stellar (and reference) lines relaxes this requirement with the square root of the number of lines, hence with the number of measurements. Applying this technique to the determination of stellar positions in images allows in principle to determine the separation between two objects with a precision far better

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than obtained from a single measurement, hence better than 1 mas. The accuracy of this measurement is however still limited to accuracy of the pixel scale. The measurement is thus only useful if changes in the separation of objects are to be measured, but not the separation itself.

3 Early Results from a Feasibility Study Neuhäuser et al. [7] started a feasibility study to measure the reflex motion of the primary component in the stellar binary HD19994 induced by the presence of an extrasolar planet candidate, which was detected by radial velocity measurements. Since the inclination angle of the orbit is unknown and can not be determined from radial velocity measurements alone, the true mass of the planet candidate HD19994 Ab is unknown and only the lower mass limit is well determined to be M sin i 1.69 MJ (Butler et al. [1] and references therein). The one-dimensional projection of the reflex motion of HD19994 A as a periodic change in the binary separation of HD19994 A and B is expected to be about 100 µas when the minimum mass is the true mass (hence the system is close to edge-on), and even larger if the true mass of the planet candidate is higher. The separation of HD19994 A and B is about 2.4 arcsec and thus well suited for AO measurements. Since the first data were obtained in December 2004 with NACO, two more epochs in November 2006 and July 2007 allowed for a first data assessment. By making use of NACOs cube mode (i.e. writing frames taken with the minimum exposure time directly into data cubes instead of individual frames with headers) the observing efficiency could be substantially enhanced in the last two epochs and several thousand frames were taken within one hour. The binary separation of HD19994 AB has a standard deviation of about 1–2 mas but the values are normally distributed (verified via Kolmogorov-Smirnov test). Thus, the standard deviation of the mean, hence the most likely value for the true separation, can be measured with a precision of about 30 µas. Again, this measurement does not allow to give an accurate value for the binary separation in arcsec but allows to measure short and long term changes in the binary separation. Strictly speaking, this result is achieved in pixel space and not in astronomical coordinates. Only if the changes of the binary separation are small (well below 10 mas) the value can be given with an accuracy in astronomical units (arcsec) which is still competitive to the achievable precision, assuming that NACOs S13 pixel scale can be accurately calibrated to about 0.05 mas/pixel. Hence, such measurements are worthless when strong stellar motion is directly detected. For example the stars in the Galactic Center show orbital motion in excess of 150 mas. Even if the individual images would yield a high precision in determining the pixel coordinates of the stars better than 1 mas, the conversion from the pixel scale to astronomical coordinates is not accurate enough and would induce an uncertainty of approx. 600 µas when measuring a change of 150 mas in the position of an object. On the other hand, the typical accuracy of the pixel scale is high enough for measurements of extremely small perturbations, as the expected range of one to few hundred µas for HD19994 A and will not limit the performance of the measurement.

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For long-term measurements, i.e. to measure the reflex motion of HD19994 A over at least one complete period of the exoplanet candidate (approx. 540 days), one needs to assure the stability of the plate scale of NACO at a level of 4/100 000. As outlined before, only globular clusters offer the needed intrinsic stability. We chose 47 Tuc due to its well determined velocity dispersion of approx. 630 µas/year [8]. Given this stability limit, the number of star that fit into the FOV of NACOs S13 camera, and the precision in determining the stellar positions in 47 Tuc, the achievable precision in stabilising the plate scale of NACO is about 100–150 µas (for the mean separation of HD19994). Thus, the achievable precision in measuring the reflex motion of HD19994 A is not limited by the centroiding precision of HD19994 A and B or the accuracy of the plate scale but by the necessity of keeping the plate scale of NACO stable to better than 4/100 000. Between November 2006 and July 2007 a change of the plate scale was observed in excess of this level and will be taken into account in the final data reduction of HD19994.

4 Impact on Current and Future VLT Instrumentation The method outlined here is most suitable for the confirmation of known extrasolar planet candidates in binary systems. Single stars orbited by extrasolar planet candidates need to have a close-by reference star and the measurement is further complicated by the differential proper motion and parallax of the two stars. Other than the influence of the stellar binary system, which can be solved for with only two observations matched in phase to the orbital period of the planet, unrelated reference stars exhibit a higher number of freedom in their differential motion. This problem is of course also existing for future interferometric measurements. The number of known extrasolar planets in binary systems is however growing [3]. Moreover, with the advent of high resolution spectrographs with adaptive optics, such as CRIRES at the VLT, the risk of light contamination and degraded RV precision in close double stars is strongly reduced. CRIRES will be equipped with a gascell to provide accurate RV measurements, especially of late type stars and brown dwarfs, a spectral regime currently inaccessible by optical RV search programs. It seems thus promising to combine the techniques of radial velocity measurements and astrometry when searching for new extrasolar planets and when following up on existing candidates. In both cases a lower and an upper mass limit (in the best case a rather accurate true mass) can be determined from the same measurements. CRIRES is already equipped with an AO system and an imaging camera, making this instrument ideally suited to be upgraded to a high spectral and spatial resolution spectro-imager where the astrometric information are available nearly for free. CRIRES would have to be equipped with an additional optical component to capture a slightly offset version of the image in respect to the entrance slit, instead of the current straight viewing geometry. This modification would also improve the tracking precision (and thus indirectly the spectral stability), since the spectroscopic target can be observed unaffected from the entrance slit, instead of centroiding only the remains of the stellar PSF on the slitjaws.

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If this concept proves to be useful for the discovery and characterisation of extrasolar planets, future instrument concepts at the VLT for stable high resolution spectrographs could benefit from this extension as well. Even in the GAIA era, when millions of stars will have positions and proper motions accurate to a few dozend µas, relative astrometry from the ground has still its niche in the discovery and characterisation of extrasolar planets, given the inhomogeneous timesampling of GAIA. Then being able to accurately calibrate groundbased imagers to a level of accuracy that is no longer limiting the final astrometric performance, astrometry obtained from direct imaging will add important information to radial velocity measurements, also for other projects, like dynamical studies of stellar clusters.

References 1. R.P. Butler et al., Astrophys. J. 646, 505 (2006) 2. F. Delplancke, S.A. Leveque, P. Kervella, A. Glindemann, L. D’Arcio, in Proceedings of the SPIE Conference: Astronomical Telescope and Instrumentation. Munich, Germany, 25–31 March 2000. Proc. SPIE, vol. 4006, p. 365 3. S. Desidera, M. Barbieri, Astron. Astrophys. 462, 345 (2007) 4. F. Eisenhauer, G. Perrin, S. Rabien, A. Eckart, P. Lena, R. Genzel, R. Abuter, T. Paumard, Astron. Nachr. 326, 561 (2005) 5. R. Launhardt et al., in Astrometry in the Age of the Next Generation of Large Telescopes. ASP Conf. Ser., vol. 338 (2005), p. 167 6. P.F. Lazorenko, M. Mayor, M. Dominik, F. Pepe, D. Segransan, S. Udry, Astron. Astrophys. 471, 1057 (2007) 7. R. Neuhäuser, A. Seifahrt, T. Roell, A. Bedalov, M. Mugrauer, IAU Symp. 240 (2006) 8. D.E. McLaughlin, J. Anderson, G. Meylan, K. Gebhardt, C. Pryor, D. Minniti, S. Phinney, Astrophys. J. Suppl. Ser. 166, 249 (2006) 9. A. Seifahrt, T. Roell, R. Neuhaeuser, in Proceedings of the 2007 ESO Instrument Calibration Workshop, arXiv:0706.2613 (2007)

CASIS: Cassegrain Adaptive-Optics Simultaneous Imaging System for the VLT M. Kissler-Patig, M. Casali, B. Delabre, N. Hubin, H.U. Käufl, P. Jolley, M. Le Louarn, S. Oberti and J. Pirard

1 Optimally Using the Cassegrain Focus behind the Adaptive Secondary Mirror 1.1 The Need for a Cassegrain Instrument Concept The VLT Unit telescope 4 will be equipped in 2013 with an Adaptive Optics (AO) Facility, replacing the current secondary mirror of the telescope with an adaptive one. The AO Facility will allow high order AO correction at all three telescope foci. The two Nasmyth foci will make use of that facility with a ground-layer AO module for the HAWK-I wide-field near-infrared imager, and with a single-conjugated AO module for the MUSE, a large optical integral-filed spectrograph. No Cassegrain instrument making use of the AO facility is yet planed.

1.2 Optimal Choice for a Cassegrain Instrument Over the last year, we have explored how to optimally use this Cassegrain focus at the AO Facility in 2013 and beyond. The capabilities that will have to be replaced are the diffraction limited imaging ones of NACO and diffraction limited spectroscopy of single objects of SINFONI. Both instruments have their own AO modules and could be transferred to another telescope. The true advantage of the AO Facility is that it provides (a) an adaptive mirror outside the instrument which makes a wide FoV correction available without the need to reimage the telescope pupil on a post-focal DM with large and heavy optics; (b) a mirror with enough actuators to potentially provide a diffraction limited image down to near IR. The optimal use of the AO Facility is to provide diffraction limited image quality over a large field of view. At Cassegrain, that large corrected field of view is hard to exploit with spectroscopy (the corresponding instrument would be M. Kissler-Patig () ESO, Karl-Schwarzschild-Str. 2, 85748 Garching, Germany e-mail: [email protected] A. Moorwood (ed.), Science with the VLT in the ELT Era, Astrophysics and Space Science Proceedings, doi: 10.1007/978-1-4020-9190-2_85, © Springer Science + Business Media B.V. 2009

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very large and complex). However, the large field of view could clearly be optimally exploited by an imaging capability. This is our proposal: the Cassegrain focus of the AO Facility should host an imager exploiting the maximum diffraction limited field possible. Below, we present the concept of CASIS, a Cassegrain Adaptive-optics Simultaneous Imaging System.

1.3 CASIS Basic Concept The basic properties of CASIS are: five simultaneous channels (0.6–1 µm, J, H, K, L + M), a 1 × 1 field-of-view corrected by multi-conjugate AO. The channels below 2.5 µm are each sampled at 15 mas/pix with a 4k × 4k detector; the L + M channel is sampled at 30 mas/pix with a 2k × 2k detector.

2 The CASIS Science Case The science case for CASIS overlaps with the one of NACO imaging (by design) and the one of the NIRCam on the JWST (see [1]). The main science drivers are: morphology and spectral energy distribution (SED) of high-redshift galaxies, multiband faint -ray burst (GRB) follow-up, circumstellar disk and star-formation region studies, deep colour-magnitude diagrams in the Local Group, and studies of the Galactic centre. CASIS will operate in the era of the JWST and the E-ELT and therefore needs to provide a compelling science case as well as unique capabilities. Unique capabilities of CASIS are five simultaneous channels and the ability to do fast (> 10 Hz) photometry. In particular the multi-band capability provides two key advantages: (1) for rapidly varying phenomena (GRBs, flares around black holes) is presents the only possibility to accurately track SED variations; and (2) the simultaneous photometry guarantees that PSF variations (expected behind an multi-conjugate AO system) do not bias SED studies of barely resolved sources (e.g. high-redshift galaxies, circumstellar disks). As a comparison, we list in Table 1 some key properties (as collected on the web) of future diffraction limited instruments which will either compete with or be complementary to CASIS. Clearly, a multi-conjugate AO imager on the ELT (here labelled as HARCE) will surpass CASIS—although the former might not reach beyond 2.5 µm in wavelength and the later has unique science cases with the simultaneous bands and fast photometry. The NIRCam onboard JWST will reach a fainter limiting magnitude and have a large field-of-view, but does not quite reach the diffraction limit of the VLT and samples the images more coarsely. The GSAOI camera at Gemini was designed to operate in good seeing (< 0.7 ?), and has a significantly narrower wavelength range (no 5 µm capability). LINC-Nirvana on the LBT uses pyramid sensors and no Laser

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Table 1 Comparison of some key properties of future diffraction limited instruments. Note that GSAOI was build to work in good seeing (< 0.7 ) only, and that LINC-Nirvana uses Pyramid wave-front sensors and no LGS CASIS

HARCE/ELT NIRCam/JWST GSAOI/Gemini LINC-Nirv./LBT

60 × 60

30 × 30

130 × 260

80 × 80

10 × 10

λ/D in K

50 mas

10 mas

64 mas

50 mas

18 mas

Sampling

15 mas/pix

4 mas/pix

32 mas/pix

20 mas/pix

5 mas/pix

Limit KAB

27.5

30

30

?

28?

0.9–2.5 µm

0.6–5.0 µm

1.0–2.5 µm

0.6–2.4 µm

2013–2018(23)

2009?–?

2010?–?

FoV

Wavelength 0.6–5.0 µm Available

2013–2023+ 2017?–. . .

Guide Star. This restricts the sky coverage and limits the maximum Strehl ratio that can be achieved. Thus, we believe that CASIS has advantages with respect to competing projects, and could become a work-horse on the VLT, complementing the cameras on the JWST and ELT.

3 CASIS Optical Concept A first optical concept for CASIS is shown in Fig. 1. From left to right: a cold (but not cryogenic, as T ∼ 235 K is sufficient to reduce the background up to 5 µm) AO module picks off the Laser Guide Stars—the science field and Natural Guide Star are passed on to the cryogenic part of the instrument—the full 3 diameter field around the scientific field of view is passed to Natural Guide Star sensors, while the 1 square science field is relayed with a flat mirror into a F/17.3 beam from which the various bands are being picked off (not shown here for clarity, but see below). In the AO module, the light is directed onto two deformable mirrors (AO mirrors), being the second and third adaptive mirrors in the system, after the adaptive secondary mirror of the telescope. The second AO mirror is conjugated to 12 km altitude, the third to 5 km altitude (the adaptive mirror of the telescope being conjugate with the ground). Five Laser Guide Stars are foreseen, each 1.5 off-axis. Natural Guide Stars can be place in a field of 3 diameter.

4 CASIS Mechanical Concept We have verified whether the five channels can fit mechanically in the volume allocated to Cassegrain instruments. One solution is shown in Fig. 2. In Fig. 2, the light would come from the telescope located above the instrument (i.e. the optical concept of Fig. 1 is now rotated by 90 degrees clockwise). The five channels are now shown as picked off along the last optical path. The volume can

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Fig. 1 CASIS—concept for optical layout. The optical for the Laser Guide Stars is shown in green, for the Natural Guide Stars in black, and for the science light in red. A colour version of this figure is available at dx.doi.org/10.1007/978-1-4020-9190-2_85

Fig. 2 CASIS—concept for mechanical envelop (left) and shown with respect to the available volume (right): the compact design comfortably fits in the allocated volume for Cassegrain instruments. A colour version of this figure is available at dx.doi.org/10.1007/978-1-4020-9190-2_85

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easily accommodate the five channels mechanically. In Fig. 2 on the right, it is clear that CASIS uses only a fraction of the reserved volume for Cassegrain instruments.

5 CASIS Adaptive Optics Concept The most challenging part of CASIS is the adaptive optics. Their goal is to provide a good and uniform performance over 1 field-of-view from the R to the M band. This goal translates into a mean Strehl ratio in K of 50% (with goal 70%), and a Strehl uniformity of 5% rms across the field. This, in turn, sets a number of requirements on the multi-conjugated AO (MCAO) system. A first performance evaluation showed that the goal requires, in addition to the telescope AO mirror with 1170 actuators, the addition of two altitude AO mirrors in the instrument with around 41 × 41 actuators each, and five Laser Guide Stars close to the science field. A wide sky coverage allowing to work near the Galactic Pole called for the use of Laser Guide Stars for the high-order corrections. Further, three Tip/Tilt sensors based on Natural Guide Stars (of magnitude < 19, located in a 3 disk but outside the scientific FoV of 1 ) provide the good homogeneity.

6 Conclusion Around 2013, the unit telescope 4 will be equipped with an adaptive secondary mirror that will allow to correct image down to the diffraction limit over a wide field of view. This capability is best exploited at the Cassegrain focus with an imager. In order to be competitive in the era of the JWST and ELTs, the imager ought to have mode not covered by these facilities. Our concept, CASIS, a five-band simultaneous imager, meets these requirements and would optimally exploit the AO facility. CASIS would also play a key role as a pathfinder for multi-conjugate AO on the ELT and serve as a test bench for multi-conjugate AO with multiple lasers. Acknowledgements We thank Thierry Fusco (Onera) who kindly provided some preliminary MCAO simulations in support of our study.

References 1. J.P. Gardner, J.C. Mather, M. Clampin et al., Space Sci. Rev. 123 (4), 485 (2006)

The Need for a General Purpose Diffraction Limited Imager at the VLT Thomas Ott, Richard Davies, Frank Eisenhauer, Reinhard Genzel, Reiner Hofmann and Stefan Gillessen

Imaging at the diffraction limit of the VLT is a crucial aspect of many astrophysical programs. While this capability is currently provided by NACO with its unique infrared wavefront sensor, the roadmap for future instrumentation does not include a diffraction limited imager. Some of NACO’s capabilities will be taken over by HAWK-I (delivering improved seeing via GLAO), SPHERE/IRDIS (optimized for high strehl on bright optical sources), and also by JWST. However, a multitude of applications—specifically, long term monitoring projects to track variability and proper motions—will not be feasible with these instruments. Science cases include obscured star forming regions, high velocity stars, the Galactic Center, and active galactic nuclei. In order to maintain the impact of the European community on such science programs, we propose a diffraction limited 1–5 micron imager with infrared wavefront sensing optimized for faint guide star targets.

1 Diffraction Limited Instruments at the VLT Currently, there are a total of five instruments installed at the VLT which operate at or near the diffraction limit: • • • • •

HAWK-I, a wide field imager. CRIRES, an adaptive optics assisted high-resolution long slit spectrograph. VISIR, an imaging camera which operates in the Mid-Infrared. SINFONI, a diffraction limited integral field spectrograph. NACO, a diffraction limited imager and long slit spectrograph.

The Laser Guide Star Facility should also be mentioned. Installed at UT4, it feeds the adaptive optics systems of NACO and SINFONI with an artificial guide star, thus overcoming the limited sky coverage of natural guide star adaptive optics systems. Additionally, NACO is unique world-wide since it features an infrared wavefront T. Ott () Max-Planck-Institute for Extraterrestrial Physics, Garching, Germany e-mail: [email protected] A. Moorwood (ed.), Science with the VLT in the ELT Era, Astrophysics and Space Science Proceedings, doi: 10.1007/978-1-4020-9190-2_86, © Springer Science + Business Media B.V. 2009

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sensor for its adaptive optics, which makes it possible to use guide stars which are bright in the infrared, but very faint in the optical, e.g. obscured star forming regions, IR-bright galactic nuclei, or the galactic center, just to mention a few. In the next few coming years until 2012, quite a few changes are expected with respect to the diffraction limited instrumentation. What concerns the infrastructure, UT4 will be equipped with an adaptive secondary and the currently installed LGSF will be replaced by an MCAO multiple laser guide star system. In order to make the best use of these facilities, the following instruments will be installed at the VLT UT4: • MUSE, an optical integral field spectrograph. • SPHERE, a narrow field imager optimized for high strehl ratios near bright optical AO guide stars. In order to install these instruments, one instrument will have to be dismounted from UT4. It has been argued that the natural lifetime of NACO will end by 2012, so the choice most probably will be this instrument. This leaves a hole in ESO instrumentation what regards a general purpose diffraction limited imager.

2 Science Cases There are many science cases which rely on precision astrometry and/or photometry over extended time periods to measure proper motions, variability, or change in morphology. At the nearby end of the scale these include solar system objects, in particular identifying and characterizing asteroids, and on larger bodies studying meteorological and seismological events. In stellar systems, one may wish to trace the evolution of planetary nebulae or supernovae, track the orbital motion of galactic clusters, or understand the kinematics of magnetars. Variability is also important in extragalactic work and has been used to show that the hot dusty region around AGN is just larger than that of the broad line region [7]; and to search for stochastic events in galactic disks which may, for example, allow one to identify supernovae and determine supernova rates in star forming regions (e.g. [8]). Here we focus on just one topic, the Galactic Center.

2.1 The Galactic Center One of the most prominent science cases for the VLT is near infrared observations of the center of our own galaxy. Because of its proximity, the Center of the Milky Way is a unique laboratory for studying physical processes that are thought to occur generally in galactic nuclei. The central parsec of our Galaxy contains a dense star cluster, with a remarkable number of luminous and young, massive stars, as well as several components of neutral, ionized and extremely hot gas. For two decades, evidence has been mounting that the Galactic Center harbors a concentration of

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dark mass associated with the compact radio source Sgr A* (diameter about 10 light minutes), located at the center of that cluster. Measurements of stellar velocities and (partial)orbits with the ESO NTT (SHARP) and VLT (NACO) have established a compelling case that this dark mass concentration is a massive black hole of about 3.6 × 106 M [1]. The Galactic center thus presently constitutes the best proof we have for the existence of massive black holes in galactic nuclei.

Progress in IR Imaging Since the late 1980’s, when the first array cameras for the near-infrared wavelength regime became available, the spatial resolution of such observations underwent significant improvements: The very first observations were seeing-limited with a resolution of ≈ 1 arcsec. Advances in detector technology allowed for the first high-speed cameras in the early 1990’s. At this time, the MPE build the speckle camera SHARP. This camera was mounted regularly at the New Technology Telescope, at that time the most advanced (in terms of image quality) telescope available on the southern hemisphere. SHARP allowed for diffraction limited (speckle-) images in the K-Band, which corresponds to 0.15 arcsec. Only ten years later, with the commissioning of NACO, another leap in image quality was possible. Its new IR wavefront sensor makes observations of the Galactic Center straightforward, since a bright IR reference source (IRS 7) is located just 6 arcsec north of Sgr A*, whereas the optical guide star is both faint (mv = 14.6) and far away (35 arcsec).

Advances in GC Research Have Been Technology Driven The huge progress made in Galactic Center research has been correlated to the progress in instrumentation: 1. In 1990, seeing limited observations allowed to observe the light distribution of the central cluster of stars. 2. Thanks to Speckle imaging, the first measurements of proper motions of stars could be published in 1996. 3. The Hubble Space Telescope was equipped with NICMOS in 1997, but due to its limited aperture, little progress in GC research was made. 4. The first accelerations of stars were published using Speckle observations made at the Keck 10 m telescope. 5. The spectacular first measured orbit of the star S2 was measured thanks to the adaptive optics at the VLT. 6. First measurements of infrared flux from Sgr A* were measured in 2003, also thanks to the adaptive optics. 7. With SINFONI at the VLT, we were able to measure the spectral slope of such flares of Sgr A*.

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8. At the time of writing, we are using Laser Guide Star adaptive optics and have a census of 16 measured orbits of stars. The upcoming foreseeable technical advances in near-infrared instrumentation are, in 2012, multi-conjugate adaptive optics systems, which allow for a larger corrected field of view, and, at the same time, the launch of the JWST. The next big step will then be, around 2017, the inauguration of the next size telescopes, e.g. the E-ELT with a diameter of 42 m. It would of course be foolhardy to predict what exciting new discoveries will be possible.

The Need for Frequent Imaging Currently, we are imaging the GC every few weeks. This frequent sampling is governed by our aim to accurately probe the gravitational potential near the central black hole. For this, we follow individual stars as test particles along their orbits. In cases when a star is moving close to Sgr A* through pericenter, it may be necessary to observe more often up to once a week, since at such events, the speed of the stars can easily exceed 5000 km/s.

Monitor Flares—The Central Accretion Zone Around Sgr A* The current consensus model is that the Sgr A* radio to IR spectral energy distribution (SED) is due to synchrotron emission. Flares arise when electrons are locally accelerated to non-thermal Γ -factors of 103 –105 , perhaps best thought of as hot spots in the accretion flow or jet orbiting just outside the last stable orbit. The orbital motion might explain the quasi-periodic structure in the IR light curves first observed by [2]. At the same time the electrons experience synchrotron cooling which is apparent in NIR color variations of Sgr A* [3]. The same electrons also inverse-Compton upscatter the photons, mainly from the submm peak. This model naturally explains the near simultaneity of IR and X-ray flares and in particular that the latter show the same characteristic time scale as the NIR flares [4]. Stand-alone IR observations provide crucial tests of the accretion physics. Photometry at a high sampling rate will test if the 17 minute quasi-periodicity discovered by [2] truly represents a fundamental time scale, such as the orbital period at the last stable orbit. Time resolved NIR-MIR SED studies provide key insights into heating/cooling processes and magnetic field strengths. The available MIRand submm-observations suggest that the flare SED peaks close to L- or M-band. Simultaneous H- to L/M-band SEDs at good SNR would be a design goal of a third generation diffraction limited imager. VISIR N-band observations provide another important data point. Simultaneous detection in NIR and MIR will stringently constrain the properties of the emitting electron population. Finally, polarimetry of flares constrain magnetic field and accretion flow geometries. All three IR flares observed polarimetrically so far exhibit the same on-flare polarization angle [5, 6]. This suggests a stable geometry of the system and needs to be confirmed, as do substantial swings of polarization angle on a time scale of 10–20 minutes.

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3 Top Level Requirements for a VLT Diffraction Limited Imager A new diffraction limited imager for the VLT should satisfy the following requirements: • • • • •

Diffraction limited imaging from H to L with a field-of-view of 1 arcmin. Observe all Stokes-parameters simultaneously. Observe several bands simultaneously. IR WFS adaptive optics, or MCAO with multiple laser guide stars. Preferably make use of adaptive secondary.

4 Conclusions We have seen that in 2012, NACO will most likely be decommissioned, either because the instrument will have reached the end of its lifetime, or because it has to make space for another instrument (MUSE, in this case). Still, a diffraction limited imager will be highly desirable in order to make the best science possible with a diffraction limited 8-meter telescope in the near-infrared wavelength domain. While some of the diffraction limited observations will be possible using either SPHERE or HAWK-I, there are many instances where a general purpose diffraction limited imager will be necessary in order to do the best science in these highly competitive areas of research. An instrument which might fill the gap could be CASIS (presented at this conference by M. Kissler-Patig), which will satisfy the imaging needs of our science cases—polarimetry will still be an issue, but could possibly still be incorporated into the design of this camera.

References 1. 2. 3. 4. 5. 6. 7.

F. Eisenhauer, R. Genzel, T. Alexander et al., Astrophys. J. 628, 246 (2005) R. Genzel, R. Schödel, T. Ott et al., Nature 425, 934 (2003) S. Gillessen, F. Eisenhauer, E. Quataert et al., Astrophys. J. 640, 163 (2006) A. Eckart, R. Schödel, L. Meyer et al., Astron. Astrophys. 455, 1 (2006) A. Eckart, F.K. Baganoff, R. Schödel et al., Astron. Astrophys. 450, 535 (2006) S. Trippe, T. Paumard, T. Ott et al., Mon. Not. R. Astron. Soc. 375, 764 (2007) M. Suganuma et al., in The Central Engine of Active Galactic Nuclei, ed. by L.C. Ho, J.-M. Wang. ASP Conf. Ser., vol. 373 (2007), p. 462 8. G. Cresci, F. Mannucci, M. Della Valle, R. Maiolino, Astron. Astrophys. 462, 927 (2007)

Exploring the Time Axis—High Resolution Timing Observations with Present and Future Instrumentation V.D. Ivanov, C. Caceres, E. Mason, D. Naef, F. Selman, C. Melo, D. Minniti and G. Pietrzynski

1 Science Drivers for “Fast” Astronomy The technological advances and the increase of the collecting are of the modern telescopes made it feasible to obtain high signal-to-noise (S/N) observations of “fast” phenomena, which was not possible earlier, even for bright objects. Although the natural application of the high timing resolution technology is to study the “fast” events, it helps to increase the S/N for events with fixed length because—indeed, the “fast” observations are only meaningful with small “dead” times between the individual integrations. Therefore, the “fast” instrumentation guarantees high filling factor, or high cadence. In other words, for a fixed duration events, such as transits, occultations, etc. the observer can collect more data with respect to the “slow” instrumentation. Last but not least, the shorter integration times feasible with “fast” instruments allow to bridge the calibration gap between 1–2–4–8–42 m telescopes that rises from the fact that most primary calibrators used today were developed on smaller telescopes. The most common science case that requires “fast” instruments are various compact objects. Their small physical size leads to event on time scale of seconds to minutes. For example, the tomography of X-ray binaries and pulsar timing do require observations on these scales. Other applications of the “fast” instrumentation are the lunar occultations whose total duration is a few seconds, the occultations of other solar system bodies, as well as various transients and even the astroseismology.

2 Transiting Extrasolar Planets About 270 of extrasolarplanets are known up to date and about two dozen show photometric transits. Transits allow a direct determination of the orbital plane inclination and the planet radius (assuming that the primary star radius is known). When V.D. Ivanov () ESO, Alonso de Cordova 3107, Santiago 19, Chile e-mail: [email protected] A. Moorwood (ed.), Science with the VLT in the ELT Era, Astrophysics and Space Science Proceedings, doi: 10.1007/978-1-4020-9190-2_87, © Springer Science + Business Media B.V. 2009

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combined with radial velocity data, they give access to the real planet mass and to its mean density revealing its nature: gaseous, icy or rocky. Future space-based transits will allow to find Earth mass undetectable via radial velocity. Furthermore, the transits allow timing measurements, similar to the technique that brought the discovery of the first extrasolar planet around the pulsar PSR 1257 + 12 [26]. [1] make some fairly robust general predictions: (i) the transit time deviations are largest in case or resonant orbits; (ii) the transit time deviations are proportional to the mass of the perturbing planet; (iii) the transit time deviations are proportional to the period of the transiting planet; (iv) the transit time deviations vary with time (see Fig. 6 in [1] and the interval between two sequential transits can change by as much as few minutes. For example, an Earth mass planet in the system of HD 209458, in a 2 : 1 resonance with HD 209458b will lead to transit time variation of order of 3 min. The insufficient timing accuracy of individual transits limited the planet searches with this method so far: most timing data came from small telescopes with large “dead” time between the images, up to 2/3 of the observation duty-cycle, i.e. [4]: 5–6 sec integrations, 12–90 sec readout times. The uncertainties of the individual transit times vary from 19 to 150 sec. [12] demonstrates the need to improve the cadence: 23 sec integration, 32 sec readout, leading to 9–15 sec errors in the timing of individual transits. Similarly, [20] and [21] obtained ∼ 1 FORS exposure per min. Many of these observations required defocussing which leads to contamination from fainter neighboring stars, sometimes equal to the flux of the host star [4]. The first detailed timing study so far [24] used 12 transits of TrES-2b [2, 7] to search for additional planets in this system. They find no convincing evidence for a second planet and set an upper mass limit for planets in low order resonances comparable or lower than the Earth mass, making these timing data the first that are sensitive to Earth mass perturbing planet. [12] studied the system of OGLETR-113b—they detect a variation of 43 sec (2.5σ ), consistent with 1–7 Earth mass planet. [4] derived residuals for individual transits for HD 189733b: ranging from 0.7 sec to 302 sec or 0.0045–5.9σ , and the errors of individual transit times are 19– 150 sec. They are consistent with perturbations from 0.15 MJ mass planet at 2 : 1 resonance orbit that would remain undetected in radial velocity observations. Here we report the first results from our timing study of individual transits with infrared (IR) detectors that allow to obtain imaging with minimum “dead” time for readout (≤ 0.1%). The IR detectors read out fast by design because the high background forces the usage of short exposures and the IR array technology has evolved to achieve reset/readout times of order of microseconds. This gives us the following advantages: (i) we observed with unprecedented time resolution of ∼ 0.1–0.2 sec; images of the host stars on individual images have S/N ∼ 50–100; (ii) we observed bright planet hosting stars without defocussing, minimizing the contamination from nearby sources; (iii) we carried out a multi-instrument campaign that allowed us to cross-correlate the transit light curves, obtaining independent verification for small scale variations due to “sunspots”, moons, rings or planet atmosphere asymmetry [5, 6, 22, 23]; (iv) we relied on the ESO timing system that provided us with uniform time accurate to better than 0.1 sec, minimizing any systematics—an advantage over other studies that collect data from various sources.

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The target selection is difficult because the hosts of the OGLE and SWEEPS planets are too faint and most other planets are too far to the North, leaving us with: (1) XO-1b [17] is a 0.90 ± 0.04 MJ planet with radius of 1.30 ± 0.11 RJ , density 1.0 ± 0.24 g cm−1 , and inclination of 87.7 ± 1.2 deg, on a 3.941534 ± 0.000027 day orbit around the G1V star GSC 02041–01657 (R∗ = 1.0 ± 0.08 R , M∗ = 1.0 ± 0.03 M , [Fe/H] = 0.015 ± 0.04; B = 11.7, V = 11.3, J = 9.9, H = 9.6, K = 9.5 mag) at 200 ± 20 pc from the Sun. XO-1b has the third largest period among the known transiting exoplanets, maximizing the transit time variations [1]. (2) WASP-2b [8] is a 0.88 ± 0.11 MJ planet with radius of 0.95 ± 0.3 RJ , density 1.3 ± 0.3 g cm−1 , and inclination angle 87 ± 3 deg, on a 2.15226 ± 0.00004 day orbit around the K1V star WASP-2 (R∗ = 0.78 ± 0.06 R , M∗ = 0.79 ± 0.15 M ; B = 13.0, V = 12.0, J = 10.2, H = 9.8, K = 9.6 mag).

Fig. 1 KS band light curve of XO-1b, obtained with ISAAC. The areas with higher scatter correspond to reduced sky transparency due to clouds

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3 Observations and Data Reduction The data for this project were acquired with ISAAC [19] at the VLT and with SofI [18] at the NTT. Unfortunately, we were hampered by poor weather conditions, with seeing up to 1.3–1.7 arcsec. We used the FastPhotJitt mode. The data were stored in 3-dimensional data cubes, each slice of which was a frame, corresponding to a single detector integration time (DIT). The typical data volume per transit was 30– 50 Gb. The detector window size had to meet a number of conflicting requirements: (i) as small as possible to reduce the minimum DIT; (ii) large enough to include a reference star; this is absolutely necessary to achieve photometric accuracy of mmags; (iii) “technological” constraints: only one window can be defined, it is centered on the detector’s center, bad pixels have to be avoided; (iv) the DITs were selected to avoid the non-linearity regimen; it depended on the filter, the seeing and the sky background. First, we removed dark/bias taken with the same windowing. Next, we applied a flat field correction. Unlike the “classical” IR data reduction, we did not construct and remove the sky because we did not jitter. Instead, during the final step—the aperture photometry—we subtracted an estimate of the sky from circular anulae centered on the target/reference. Having only two bright stars in the field we could

Fig. 2 χ 2 minimization curve for the fit for two XO-1b transits. A colour version of this figure is available at dx.doi.org/10.1007/978-1-4020-9190-2_87

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not use PSF fitting and the PSF varied too much from frame to frame to allow us to create a PSF model from averaging a few frames taken close in time. This possibility will be studied in the future. A sample light curve is shown in Figure 1.

4 Preliminary Results: Transit Timing Measurements We have obtained data for six transits of two extrasolar planets. Two transits were observed in parallel by both ISAAC and SofI, and we expect to observe two more. The first step in our timing measurements was to fit theoretical light curves to the observations. For now, we used the parameters derived in [8, 17] to generate analytic light curves following the algorithm of [16]. The actual fit was done minimizing χ 2 , weighted by the flux (Fig. 2) to account for the transparency variations that change the S/N during the observations. The best accuracy we achieve timing a single transit is ∼ 1 min. This is much larger than the few tens of seconds, predicted by our Monte Carlo simulations. Currently, we are working to reduce further the systematical effects and to establish longer time baselines.

References 1. 2. 3. 4. 5. 6. 7. 8. 9. 10. 11. 12. 13. 14. 15. 16. 17. 18. 19. 20. 21. 22. 23.

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Advanced Calibration for Quantitative Astrophysics: 2nd Generation VLT Instruments and Beyond Florian Kerber, Paul Bristow and Michael R. Rosa

1 Introduction We describe the activities of the Calibration and Modeling Support Group (CMG) in ESO’s Instrumentation division related to the calibration and physical description of the instrument with the objective to support the data reduction of science data and to facilitate operations. We summarize the concept and lessons learned with a view to future applications in 2nd generation VLT instruments and beyond.

2 Building and Operating an Instrument The life cycle of an instrument can be described as follows: 1. 2. 3. 4. 5. 6. 7.

Definition based on Science Requirements. Optical Design (code V, Zemax). Scientific & Technical Optimization. Build, Test and Commission. Operation and Data Flow. Calibration of Instrument. Scientific Data and Archive.

Experience shows that it is difficult to ensure that the know how and expertise that went into designing and building the instrument (steps 1–3) also is fully used in its calibration and scientific operations (steps 6 and 7). A case in point is the wavelength calibration in which well-understood physics is employed to design a spectrograph with an optimal format while during operations the dispersion solution is then derived in a purely empirical manner e.g. fitting polynomials to a sparse calibration line spectrum. F. Kerber () ESO, Karl-Schwarzschild-Str. 2, 85748 Garching, Germany e-mail: [email protected] A. Moorwood (ed.), Science with the VLT in the ELT Era, Astrophysics and Space Science Proceedings, doi: 10.1007/978-1-4020-9190-2_88, © Springer Science + Business Media B.V. 2009

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3 Advanced Calibration: VLT 2nd Generation, E-ELT, . . . One way to overcome the limitations of classical instrument building and operations is to capture all the engineering information in a physical model-based description of the instrument. This model accompanies the instrument throughout its life cycle and is used to drive the science data reduction pipeline. In our concept the model is combined with validated physical data of the instrumental components and calibration reference data.

3.1 Instrument Physical Modeling In the case of wavelength calibration the task is to relate x, y positions of photons on the detector to the original position of the incoming light at the entrance slit and its wavelength, hence providing physical meaning to the information from the astrophysical source. The concept of using the relevant physics to describe the performance of an astronomical spectrograph was developed more than 10 years ago by M. Rosa and P. Ballester [2] at ESO and has matured into the following strategy. We first construct a streamlined model of the dispersive optics based on the (e.g. Zemax/Code V) optical design. This model has a parameter file that describes the orientation, relative positioning and optical properties (e.g. grating constants, detector array dimensions) of the relevant components. The core of the model is a function that will return the detector co-ordinates for a given wavelength and entrance slit position. Further modules facilitate the calculation of throughput and the addition of noise and so on. By making iterative calls to the core function for a range of wavelengths and slit positions we can create simulated 2D data (Fig. 1) and other products such as detector array wavelength maps. The model parameter set can be optimized to reflect the performance of the operational instrument with suitable calibration data in a similar way that a polynomial dispersion solution would be fit. The difference is that the parameters optimized here have physical meaning and represent the actual configuration of the instrument. In order to perform the optimization we iteratively call the core model function for a list of standard wavelengths, comparing the results of each iteration to the measured centroids for these wavelengths as measured in the calibration data. We employ the simulated annealing technique [12, 14] to continually adjusts the model parameters until the best match between predicted and measured centroids is found. As it has evolved, variations of this strategy have so far been applied to FOS and STIS (where the application was recognized by a NASA group achievement award) on HST and UVES, CRIRES and X-shooter on the VLT.

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Fig. 1 2D simulated observation of a Th-Ar hollow cathode lamp with X-shooter. A colour version of this figure is available at dx.doi.org/10.1007/978-1-4020-9190-2_88

3.2 Calibration Reference Data Calibration reference data serve as “ground truth” to operational instrument calibration efforts. Our goal is to use only data that are traceable to laboratory standards as calibration reference data. All such data need to be fully described by meta-data explaining how they were obtained and what accuracy can be expected. In the following we give some recent examples.

Spectro-photometric Standard Stars for the Near Infrared (IR) In preparation of the operations of X-shooter [5] a need for better IR flux calibration standards was identified. The aim of this ESO observatory program (PI: S. D’Odorico) is to establish a set of spectro-photometric standard stars over a very wide wavelength range from 320 nm to 2500 nm. The strategy is to extend the useful range of the well established optical flux standards [6, 7] into the near IR using SINFONI integral field unit (IFU) without adaptive optics (FOV 8 × 8 ). This instrument provides sufficient spectral resolution to observe through the forest of telluric absorption and OH emission lines while avoiding uncertainties related to slit losses. Flux measurements are done in a few well chosen windows where atmospheric transmission is > 99%. Interpolation between these measurements is done with state-of-the-art stellar atmosphere models for the white dwarfs (see T. Rauch’s web site http://astro.uni-tuebingen.de/~rauch/). Three HST white dwarfs serve as primary reference [3]. For these agreement between flux measured above the atmosphere and stellar models is 1–2%. The goal for the ground based standard star observations was to achieve about 10% accuracy. Early results indicate that a precision of about 5% is within reach for J, H and K band. Once fully

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validated this new set of standard stars will be useful to instruments other than Xshooter as well.

Wavelength Calibration Sources for the Near-IR Traditional approaches to wavelength calibration in the near-IR relied on atmospheric features [15], in particular the rich OH emission spectrum of the night sky [1]. This approach has significant limitation for high resolution spectroscopy and we decided that in order to realize the potential accuracy of ESO’s Cryogenic High-Resolution IR Echelle Spectrometer (CRIRES) [8] an external calibration source would be desirable. To this end ESO, in collaboration with the Space Telescope European Co-ordinating Facility (ST-ECF) and the US National Institute of Standards and Technology (NIST), embarked on a project to establish ThAr wavelength standards in the 950–5000 nm operating range of CRIRES. Th-Ar hollow cathode lamps (HCLs) are a well establish wavelength standard in the ultraviolet (UV) and Visible providing a rich spectrum [13]. Since there is only one isotope in nature and the nuclear spin is zero the use of Th avoids the complexity of isotope- and hyperfine structure in the spectrum. Dedicated laboratory measurements at NIST have resulted in about 2400 lines between 750 and 4800 nm with highly accurate (accuracy: 0.001 cm−1 for strong lines) wavelengths. The calibration accuracy relative to laser measurements of Th [4] is ≈ 1.4 × 10−8 . Details of this effort and the operations of HCLs are described in [9] while the line list is published in Kerber et al. (2008 submitted). The Th-Ar HCLs provide a high density of sharp well-characterized emission lines combined with the ease and efficiency of operation of a commercial discharge lamp. They now are the backbone of wavelength calibration of CRIRES up to 2500 nm. For longer wavelengths gas cells (N2 O and OCS) are being established as calibration sources. With these developments wavelength calibration in the near-IR will become very similar to the UV-visible region, and it is possible to support high accuracy absolute wavelength calibration without having to rely on atmospheric features.

3.3 Optimizing Calibration Systems A combination of laboratory measurements and a physical instrument model is a very powerful tool for assessing the predicted performance of an instrument or its calibrations subsystem. For the selection of the best-suited wavelength calibration sources for the near-IR arm of X-shooter we did an in-depth analysis [10, 11]. As a result we have been able to identify a combination of the noble gases Ne, Ar and Kr as the best three-lamp combination. Our analysis provides a quantitative order by order prediction about the number of lines available from a given source, their relative intensities—including the effect of the blaze function—and an estimate of the line blending between sources.

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We expect that such an analysis will become routine for the design phases of future instruments. In the case of X-shooter we had been limited to the noble cases by the availability of IR line data. Since the data on the IR spectra of many elements is sparse we have launched a project to study a significant number of elements with the aim of producing a comprehensive database that will help us to select the optimal calibration source for a given combination of spectral resolution and wavelength range for future IR spectrographs.

4 Summary and Outlook The activities of the Calibration and Physical Modeling Support Group in ESO’s Instrumentation Division try to combine several tasks and methods in order to support instrument development and science operations. The optical model of the spectrograph based on engineering information provides predictive power and a quantitative understanding of the instrument’s performance combined with a solid estimate of the associated errors of the measurements. Calibration reference data traceable to laboratory standards provide the ground truth needed for quantitative calibration. A combination of the modeling techniques and calibration reference data can be used to optimize instrument performance throughout all phases of the life cycle of an instrument: design, manufacture, testing and operations. In addition the group supports the implementation of its products in the operational science data reduction software. Furthermore we fully participate in the integration and testing, commissioning and science verification in close collaboration with the instrument team. Key to success and to achieving the best science product is an integrated approach that combines the above tasks early on from the design phase. Planning for the 2nd generation of VLT instruments and E-ELT instruments should make full use of these capabilities available at ESO. Acknowledgements collaboration.

We thank our colleagues at NIST and in the Instrument Teams for the good

References 1. 2. 3. 4. 5. 6. 7. 8. 9. 10. 11. 12.

M.C. Abrams, S.P. Davis, M.L.P. Rao et al., Astrophys. J. Suppl. Ser. 93, 351 (1994) P. Ballester, M.R. Rosa, Astron. Astrophys. Suppl. Ser. 126, 563 (1997) R. C Bohlin, in ASP Conf. Ser., vol. 364, ed. by C. Sterken (2007), pp. 315–330 W. DeGraffenreid, C.J. Sansonetti, J. Opt. Soc. Am. B19, 1719 (2002) S. D’Odorico, H. Dekker et al., Proc. SPIE 6269, 33-1 (2007) M. Hamuy, A.R. Walker, N.B. Suntzeff et al., Publ. Astron. Soc. Pac. 104, 553 (1992) M. Hamuy, N.B. Suntzeff, S.R. Heathcote et al., Publ. Astron. Soc. Pac. 106, 566 (1994) H.-U. Käufl, P. Ballester, P. Biereichel et al., in Proc. SPIE, vol. 5492, ed. by A.F.M. Moorwood, M. Iye (2004), pp. 1218–1227 F. Kerber, G. Nave, C.J. Sansonetti et al., in ASP Conf. Ser., vol. 364, ed. by C. Sterken (2007), pp. 461–478 F. Kerber, F. Saitta, P. Bristow, XSH-TRE-ESO-12000-0163 (2007a) F. Kerber, F. Saitta, P. Bristow, Messenger 129, 21 (2007b) S. Kirkpatrick, C.D. Gelatt Jr, M.P. Vecchi, Science 220, 671 (1983)

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13. B.A. Palmer, R. Engleman Jr., Los Alamos Rep., LA-9615 (1983) 14. W.H. Press, B.P. Flannery, S.A. Teukolsky, W.T. Vetterling, in Numerical Recipes, The Art of Scientific Computing (Cambridge Univ. Press, 1986) 15. P. Rousselot, C. Lidman, J.-G. Cuby et al., Astron. Astrophys. 354, 1134 (2000)

Quantitative Near-IR Spectroscopy of OB Stars M.F. Nieva, N. Przybilla, A. Seifahrt, K. Butler, H.U. Käufl and A. Kaufer

Abstract Massive OB-type stars are associated with star-formation regions. Their high luminosity allows us to derive present-day chemical abundances over large distances, in the Milky Way and the Magellanic Clouds. We discuss first quantitative results from an analysis of high-quality spectra of early B-type stars in the solar vicinity obtained with CRIRES on the VLT/UT1. This work includes the identification of metal lines never resolved before and benchmark tests of spectrum synthesis, which is more challenging than in the optical because of the amplification of nonLTE effects in the near-IR. Results from the visual and near-IR spectra agree for those elements with reliably calibrated model atoms. The ionisation equilibria of He I / II and C II / III are established simultaneously for both spectral ranges, which confirms the self-consistency of the analysis. This will allow us to derive reliable chemical abundances in (dust-enshrouded) regions throughout the Galactic disk, which are observable only at near-IR wavelengths. Moreover, an accurate modelling of early-type star spectra is important for the use of these objects as telluric standards for other spectroscopic investigations in the near-IR. This pioneering work is part of our preparation for the science to be done with the next generation of ELTs.

1 Analysis, Preliminary Results & Prospects Our dataset consists—so far—of high-S/N and high-resolution spectra for 3 apparently slow-rotating early B-type dwarfs and giants located in the solar vicinity, obtained with CRIRES [1] in the J, H, K and L bands. A hybrid non-LTE approach based on hydrostatic metal-blanketed model atmospheres in plane-parallel geometry and non-LTE line-formation with DETAIL and SURFACE, as discussed by [2], was chosen for the spectrum synthesis. State of the art model atoms for H [3], He [4] and C [5, 6] are extended to an application in the near-IR. The stellar parameters have been previously derived in [5, 6] from high-quality optical spectra obtained with FEROS. There, a careful and self-consistent analysis was performed, matching the Balmer lines and the He I / II and C II / III / IV ionisation equilibria, thus providing highly accurate atmospheric parameters and chemical abundances. The same non-LTE level populations and atmospheric parameters as derived from the optical spectrum were employed here for the calculations in M.F. Nieva () Dr. Remeis-Sternwarte Bamberg, Univ. Erlangen-Nürnberg, Sternwartstr. 7, 96049 Bamberg, Germany e-mail: [email protected] A. Moorwood (ed.), Science with the VLT in the ELT Era, Astrophysics and Space Science Proceedings, doi: 10.1007/978-1-4020-9190-2_89, © Springer Science + Business Media B.V. 2009

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Fig. 1 Line fits to hydrogen, He I and C II / III. A colour version of this figure is available at dx.doi.org/10.1007/978-1-4020-9190-2_89

the near-IR. Excellent quantitative agreement (within small 1σ -uncertainties) from the analysis of both spectral regions is obtained, which is not trivial because of enhanced non-LTE effects at near-IR wavelengths. Some emission lines, as in the case of C II / III shown in the right panel of Fig. 1, can be successfully modelled only in non-LTE, whereas they cannot be reproduced even qualitatively in LTE. The agreement between the optical and near-IR results is a basic requirement for avoiding systematic errors in further abundance analyses of stars only observable at near-IR wavelengths, such as those located in dust-enshrouded star-formation regions. This work is in its preliminary phase and will find broad applications not only in our Galaxy with the present generation of large telescopes, but also in further studies of extragalactic objects with the new generation of ELTs, which will resolve crowded fields with adaptive optics only in the near-IR. Moreover, this modelling and analysis of OB-stars in the near-IR at high resolution can be useful for telluric line corrections in the reduction process of other science spectra. These stars are ideal telluric standards, as they show few lines in the near-IR. Some improvements in the data reduction and in the model atoms for nonLTE calculations are being implemented, and the identification of metal lines is in progress. Other regions of the near-IR spectrum with stronger telluric contamination have to be further investigated. Studies of more stars will be necessary in order to constrain our analysis methodology within a larger stellar parameter range.

References 1. 2. 3. 4. 5. 6.

H.U. Käufl, P. Ballester, P. Biereichel et al., SPIE 5492, 1218 (2004) M.F. Nieva, N. Przybilla, Astron. Astrophys. 467, 295 (2007) N. Przybilla, K. Butler, Astrophys. J. 609, 1181 (2004) N. Przybilla, Astron. Astrophys. 443, 293 (2005) M.F. Nieva, N. Przybilla, Astrophys. J. 639, L39 (2006) M.F. Nieva, N. Przybilla, Astron. Astrophys., accepted (2008); arXiv:0711.3783

The Very Large Telescope Interferometer in the ELT Era M. Schöller, F. Delplancke, A. Glindemann and A. Richichi

1 VLTI—The Current Status The Very Large Telescope Interferometer (VLTI, see Fig. 1) in its basic configuration of four 8 m and four 1.8 m telescopes, spanning baselines between 8 m and 200 m, will be at the forefront of infrared interferometry for at least the coming two decades. It is currently the only interferometer which gives access to closure phases and more than one baseline with 8 m-class telescopes. Arguably, it is also the only existing interferometer which can easily adapt to different science cases with its movable smaller telescopes. Its suite of current and envisioned instruments is absolutely unique and unlike other interferometers it has a large user base, running in the VLT operations model. Today, the VLTI uses all four 8.2 m Unit Telescopes (UTs), which are each equipped with an adaptive optics system (MACAO). These telescopes span six baselines between 47 m and 130 m. They are accompanied by four 1.8 m Auxiliary Telescopes (ATs), which can be moved to 30 different stations, giving access to baselines between 8 m and 202 m. The delay line tunnel holds currently six delay lines, with room for two more. Thus, with new instruments, VLTI will be able to combine six or even eight telescopes in the future. In the laboratory, the tip/tilt tracker IRIS and the fringe tracker FINITO stabilize the two dominant atmospheric disturbance modes. The two science instruments MIDI and AMBER, together with the commissioning instrument VINCI, make up the suite of beam combiners. The PRIMA dual feed facility is being installed during 2008.

2 Optical Interferometry—Some Facts We would like to look first at the unique resolving power of optical interferometers. There seems to be quite some confusion about spatial resolution of optical systems, especially when comparing single telescopes and interferometers. M. Schöller () ESO, Casilla 19001, Santiago, Chile e-mail: [email protected] A. Moorwood (ed.), Science with the VLT in the ELT Era, Astrophysics and Space Science Proceedings, doi: 10.1007/978-1-4020-9190-2_90, © Springer Science + Business Media B.V. 2009

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Fig. 1 The summit of Paranal in early 2005. The picture is dominated by the enclosures of UT1 to UT4 from left to right. In the background between UT3 and UT4 one can see the enclosure of the VLT Survey Telescope (VST). Flanking the VST in the foreground, the first two ATs can be seen. To the left of the ATs and at the far right of the picture, the small enclosures of the two test siderostats are visible, which have been removed in the meantime. The white lids in the front to the left are covering some of the 30 station pits to which the ATs can be moved. A colour version of this figure is available at dx.doi.org/10.1007/978-1-4020-9190-2_90

Fact is that the spatial resolution of a single telescope, defined as the distance from the center of the point spread function (PSF) to the first null, is 1.22 Dλ , and λ 2B for an interferometer, with D and B the diameter and the baseline length, respectively. Thus, a full aperture telescope with the same diameter as an interferometer’s baseline length has a nearly 2.5 times worse spatial resolution (see Fig. 2). Of course, interferometers tend to have much longer baselines than typical telescopes diameters and win enormously in spatial resolution. On the other hand, the high spatial resolution comes with the price of numerous large side peaks (the fringes) and having achieved the high spatial resolution in only one direction. This can be overcome by using different baselines and making use of earth rotation, which helps to fill the uv-plane. Also, the full aperture is much more sensitive than the two smaller telescopes, where the light has to bounce off dozens of optical surfaces, and goes through beam splitters, dichroics and spatial filters before it is interfering. Interferometers are also limited to very short exposure times in at least one of their subsystems. Another aspect concerning the quality of images is the ability to calibrate the visibility and phase information in an optical system looking through the earth’s

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Fig. 2 Left: Point spread functions (PSF) for two different apertures with identical diameters. In the upper panel the aperture is completely filled. In the lower panel it resembles an interferometer, where only two holes at the extremes are used. Right: Cuts through the two PSF. One can see that the first null for the interferometer PSF is more than a factor two closer in when compared to the full aperture. A colour version of this figure is available at dx.doi.org/10.1007/978-1-4020-9190-2_90

atmosphere at every spatial frequency. Here, sparse apertures usually win over completely filled apertures which are typically corrected with adaptive optics. The best example for this is the very successful aperture masking experiment on the Keck telescopes [7] which is using data taken with the Keck AO and an aperture mask. While diffraction limited sampling of the image distribution on sky is one aspect of interferometers, astrometric information is another, very important one. Comparing astrometric accuracy achieved with single telescopes to that expected from interferometers shows that interferometers scale again with the baseline and win by a factor of 2. Spectro-astrometry gives about 1 mas astrometric accuracy on a 4 m telescope in the visible (e.g. [6])—a factor of 30 to 50 better than the diffraction limit. NACO data from the galactic center (65 mas resolution) give an overall astrometric accuracy of 450 µas, i.e. a factor of about 150 (MPE GC team, private comm.). PRIMA aims at 10 µas accuracy, which on the UTs (maximum 130 m baseline, leading to a spatial resolution of 3 mas in K) is a factor of about 300. “Superresolution” describes methods which aim at obtaining spatial information slightly beyond the diffraction limit. While this works in the laboratory with an in-

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finite amount of light, no astronomical experiments are known to the authors of this paper. Since interferometers and telescopes are sharing the same physics, a method developed for a single aperture in this context is very likely transferable to an interferometer. Another misconception comes from the fact that large apertures seem to come with a huge number of photons. While this is true for an incoherent mode, what counts for a coherent mode is the number of photons per turbulence cell, which is independent of the telescope diameter, and is determined by the brightness of the source and the atmosphere, i.e. the quality of the site. Another argument that is often used against interferometers is that they are only used by a few experts. This is not the case for the VLTI (see e.g. [5]). On 2008 February 19, the ESO telescope library lists 87 papers with data from the VLTI, by 51 different first authors. In total during the ESO observing periods P73–P80 (April 2004 to March 2008), there were 141 AMBER programs, 194 MIDI programs, or in total 309 VLTI programs scheduled, with 66 different AMBER PIs, 72 different MIDI PIs, or in total 116 different VLTI PIs from 15 countries (numbers courtesy of M. Wittkowski).

3 VLTI—The Future It is not an easy task to foresee the future of an astronomical facility like the VLTI on a timescale of ten or twenty years. We have based our projections mainly on the following two assumptions. 1. High spatial resolution gives access to unique science which cannot be tackled by any other method and there is a large enough scientific community interested in this science. 2. ESO will be pressed into cost savings on Paranal the very latest when the ELT goes into construction, maybe even already before, due to ALMA. There will not be any money for 3rd (or 4th) generation facility instruments. UTs will have one facility instrument and comparably cheap visitor instruments will be much more common. As already stated above, the VLTI is a quite unique installation and will be for some time. An overview about the other optical interferometers already in operation or under construction can be found in Table 1. A next generation optical interferometer in the style of the VLA with e.g. 27 4 m telescopes is very unlikely in the next ten to twenty years. For the future, there are several obvious upgrade paths that the VLTI will or might follow. The second generation instruments will make use of the full potential that the VLTI provides today (MATISSE), soon (Gravity), or with some minor upgrades (VSI). All of these instruments will combine more than three beams simultaneously and the current fringe tracking system has to be upgraded to allow up to six beams being tracked at the same time. Operationally, another logical step is the purchase of two more star separators for the ATs, making it possible to also combine all four ATs in dual feed mode. Finally, to use the full volume provided by the infrastructure,

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Table 1 Overview of the optical interferometers in existence or under construction Number and size of telescopes PTI NPOI

Movable

Maximum baseline

Wavelength range

3 × 0.5 m

No

110 m

NIR

6 × 0.4 m

Yes

90 m (430 m)

VIS

(4 × 1.8 m) 6×1 m

CHARA

No

330 m

VIS/NIR

MROI

≤ 10 × 1.4 m

Yes

400 m

NIR

KeckI

2 × 10 m

No

75 m

NIR/MIR

O’HANA

2 × 10 m, 2×8 m, 2×4 m

No

∼ 800 m

NIR

one could think about two additional Delay Lines and two or even four additional ATs. Other potential VLTI upgrades could come from connecting all Paranal telescopes (currently ten) with fibers similar to O’HANA on Mauna Kea [3]. Another interesting way at mid infrared wavelengths could be heterodyne detection with a frequency comb [2] which would allow to interfere the telescopes electrically. Intensity interferometry [1] is gaining sensitivity with the large UT apertures and could be a stepping stone to later combine TMT and ELT. Potential cost saving roads for the VLT include the reduction of the number of instruments, which saves both on the building of these instruments and their operations costs. Also, a reduction of the number of instrument modes offered to the astronomical community can reduce the number of specialists (i.e. astronomers) which are needed to operate them. Finally, the use of all four UTs at the same time (e.g. with VLTI or ESPRESSO) with one night astronomer can reduce costs significantly. It would make a lot of sense to team up VLTI and ESPRESSO [4] for another reason. Operationally, they seem to be a perfect match. While VLTI performance is strongly depending on seeing and coherence time, the performance of ESPRESSO is not, as long as the seeing is only slightly worse than ESPRESSO’s 1.1 fiber size. Switching from VLTI to ESPRESSO and back can be made relatively fast. Having both instruments available simultaneously and using the VLTI only under exceptionally good conditions, would make optimum use of the facility.

References 1. C. Barbieri et al., in Science with the VLT in the ELT Era, ed. by Alan Moorwood. Astrophysics and Space Science Proceedings (Springer, Netherlands, 2009), p. 249 2. A. Glindemann, H.-U. Käufl, Heterodyne interferometry with a frequency comb—the cornerstone of optical very long baseline interferometry?, in Visions for Infrared Astronomy, Paris, (2006), pp. 347–350 3. A. Seifahrt et al., in Science with the VLT in the ELT Era, ed. by Alan Moorwood. Astrophysics and Space Science Proceedings (Springer, Netherlands, 2009), p. 469

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4. L. Pasquini et al., in Science with the VLT in the ELT Era, ed. by Alan Moorwood. Astrophysics and Space Science Proceedings (Springer, Netherlands, 2009), p. 395 5. G. Perrin et al., in Science with the VLT in the ELT Era, ed. by Alan Moorwood. Astrophysics and Space Science Proceedings (Springer, Netherlands, 2009), p. 81 6. M. Takami, J. Bailey, T.M. Gledhill et al., Mon. Not. R. Astron. Soc. 323, 177 (2001) 7. P. Tuthill, J.D. Monnier, W.C. Danchi, Nature 398, 487 (1999)

VLTI and Beyond: The Next Steps in AGN Research with Interferometers Klaus Meisenheimer

Starting from the current state of infrared interferometry, this contribution outlines the scientific challenges in AGN research for the next decade. The instrumental requirements in terms of resolution and sensitivity are compared with current and 2nd generation VLTI instruments. I argue that investigations of the central engine will require telescope arrays with baselines of kilometers.

1 Where Do We Stand Today? With few exceptions, all successful interferometric observations of AGN at λ < 0.8 mm have been obtained with the MID-infrared Interferometric instrument (MIDI) at the VLTI. After early observations of NGC 1068 [2] interferometric fringes have been detected for 10 AGN, including the quasar 3C 273 at z = 0.158. In most cases the shortest VLTI baselines were used for this MIDI AGN Snapshot Survey (MASS). Thus most of the distant sources could not be resolved up to now and future observations with longer baselines are needed. Currently detailed, well resolved measurements exist only for the two nearest Seyfert 2 galaxies, Circinus and NGC 1068, and for the radio galaxy Centaurus A (see Fig. 1 and contribution by Jaffe et al., page 89). Since MIDI combines only two telescopes at a time, the phase cannot be determined and the correlated fluxes have to be interpreted by model fits. Essentially, we have reached the state-of-the-art of radio interferometry in the mid-1950s. The history of radio interferometry (Fig. 2) might be used to anticipate the future of infrared interferometry: It took more than 10 years to get images by radio interferometry and another six before in 1974, Ryle et al. presented the first map of Cygnus A, which was sufficiently resolved that the hot spots could be identified—a detection which immediately led to the still valid “twin-exhaust model” of radio galaxies [1]. With the next generation of VLTI instruments, images of the dust emission from AGN with a quality similar to radio maps in the late 1960s should be possible (lower left panel in Fig. 2). So the IR interferometry seems on track to follow the radio success, but we will need the same devotion and patience in improving resolution and sensitivity as Ryle and others showed in their time. K. Meisenheimer () Max-Planck-Institut für Astronomie, Königstuhl 17, 69117 Heidelberg, Germany e-mail: [email protected] A. Moorwood (ed.), Science with the VLT in the ELT Era, Astrophysics and Space Science Proceedings, doi: 10.1007/978-1-4020-9190-2_91, © Springer Science + Business Media B.V. 2009

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Fig. 1 MIDI results on the 10 µm emission of two Seyfert 2 galaxies, the Circinus galaxy (left) and NGC 1068 (middle), and the radio galaxy Centaurus A (right). The dust distribution of the Seyfert 2s shows a compact disk (yellow) which is embedded in a larger, geometrically thick torus [6]. The 10 µm emission of Cen A is dominated by an unresolved, non-thermal core, surrounded by a feeble dust disk [3]. A colour version of this figure is available at dx.doi.org/10.1007/978-1-4020-9190-2_91

2 The Scientific Challenges Ahead Verify the Unified Scheme for Seyfert Galaxies After MIDI has resolved the dust emission from two nearby Seyfert 2 galaxies, which we identify with the torus postulated by the unified scheme, it remains to be shown that similar dust structures are present in Seyfert 1 galaxies. The first attempt to resolve the closest Seyfert 1 galaxy, NGC 4151 at 2 µm failed [5]. As demonstrated in Fig. 3 the accretion disk in a Seyfert 1 far outshines the torus emission at λ < 3 µm. Detecting the torus emission at 10 µm seems much easier. The Seyfert 1 galaxies well accessible from the VLTI are 3 to 6 times more distant than NGC 1068 and have been observed with MIDI with the short baseline UT2–UT3 only. So, it is no surprise that they remain unresolved so far. But if similar dust tori exist in their cores they cannot escape detection with the longest VLTI baselines.

Dissect the “Dust Torus” The detailed studies of the dust distribution in the Circinus galaxy and NGC 1068 already have unveiled hints for fine-structure: besides the central concentration in a disk- (or ring-)like structure, both “wiggles” of the observed visibilities around the model fit and the low surface brightness of the extended “torus” emission point towards their makeup of clumps or filaments [6]. Likewise, hydrodynamical models of the torus predict filaments produced by cooling instabilities, which merge into a turbulent disk [4]. The expected images cannot be described by smooth models and definitely require imaging interferometric instruments. It is one of the major

VLTI and Beyond: The Next Steps in AGN Research with Interferometers

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Fig. 2 The history of radio interferometry demonstrated by the example of the classical radio galaxy Cygnus A (panels on the left). It can be expected that the development from MIDI to MATISSE proceeds at a similar pace (right panels). The reconstructed MATISSE image was provided by K.-H. Hofmann (MPIfR, Bonn). A colour version of this figure is available at dx.doi.org/10.1007/978-1-4020-9190-2_91

science goals of the 2nd generation VLTI instrument MATISSE to dissect these fine-structures in the dust torus (see reconstructed image in Fig. 2).

Investigate the Fueling Process It would be short-sighted to regard the torus merely as a light-blocking nuisance. It is rather the reservoir of gas which supplies the accretion disk with fuel. To understand the central fueling process, ideally, one would witness directly how parts of the inner torus (the turbulent disk seen in the hydrodynamical models?) are caught by the black hole’s attraction and dragged towards the center. In nearby Seyfert galaxies this process should occur within radii r < 1 pc and on time-scales of years. Observations with angular resolutions < 1 mas both of the continuum emission of the hottest dust (2 < λ < 4 µm) and of emission lines (Brα , Brγ , FeII) need to be combined for such a study.

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Fig. 3 Spectral energy distribution of Seyfert 1 (face on view i = 0◦ ) and Seyfert 2 galaxies (edge on view i = 90◦ ), calculated for a clumpy torus model. The contribution from the accretion disk (core) and the dust torus are shown separately. Note that the total torus emission needs to extend over many resolution elements to be detectable. Mid-infrared bands used by MIDI (N-band) and by MATISSE (L + N-band) are indicated. A colour version of this figure is available at dx.doi.org/10.1007/978-1-4020-9190-2_91

Resolve the Structure of the Broad Line Region Despite the huge progress made by reverberation mapping, the origin and structure of the Broad Line Region (BLR) remains one of the greatest mysteries in AGN: is it composed of spherical clouds or of elongated filaments? Does it fill a spherical volume, a flattened structure or bipolar cones governed by out-flowing winds from the accretion disk? Is it expelled from the accretion disk or accreted from material falling towards the black hole? Definite answers to these questions cannot be gained without imaging the BLR. On the other hand, a clear view of the shape and the sub-structure of the BLR would enable us to use line profiles and line ratios for the diagnosis of the central engine. Thus, resolving the BLR remains the prime challenge in AGN research. However, we know from reverberation mapping that the size of the BLR in nearby AGN does not exceed one light month (0.026 pc). In the closest Seyfert galaxies this corresponds to < 0.5 mas. Thus, a resolution of < 0.1 mas is required to dissect the BLR. In addition, very high dynamic range (> 500 : 1) and spectral resolution (R > 1000) will be needed.

3 Reality Check: What Instruments Are Required? In Table 1 the requirements that are needed to face the above scientific challenges are compared with the capabilities of existing and planned VLTI instruments. Testing the unified scheme and dissecting the dust torus are within the reach of MIDI,

VLTI and Beyond: The Next Steps in AGN Research with Interferometers

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Table 1 Instrumental requirements to meet the scientific challenges. The last column indicates whether VLTI instruments are capable of doing the job (X = no) Science goal

Requirements:

VLTI

λ [µm]

beam [mas]

Ntel

DynRange

Verify the unified scheme

8. . . 13

10

2

> 50 : 1

Dissecting the “dust torus”

3. . . 13

<3

3+

> 200 : 1

2nd

The fueling process

2. . . 8

<1

3+

> 400 : 1

X

The structure of the BLR

2. . . 13

< 0.1

3+

> 500 : 1

X

√

AMBER and MATISSE. But it is obvious that those observations which could most significantly advance our understanding of the AGN phenomenon (fueling process, structure of the BLR) clearly exceed the resolution of the VLTI by an order of magnitude. Addressing these questions will need a telescope array with a maximum baseline > 20× longer than that of the VLTI, without compromising the sensitivity and dynamic range.

4 Where to Go? To me it seems clear that there is no realistic path to such a telescope array within the next decade, unless one starts from the existing facility—the VLTI. I propose, therefore, to extend the existing VLTI into a larger array by adding 5 or 6 telescopes

Fig. 4 Schematic sketch how to extend the existing VLTI to a much larger array. A colour version of this figure is available at dx.doi.org/10.1007/978-1-4020-9190-2_91

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of the 10 m class. In Fig. 4 a schematic layout of such an “XVLTI” is sketched: its “unit baseline” would be 500 m, distributed over a large area to provide maximum baselines of > 2500 m. At λ < 5 µm fiber links and integrated optics (of highly dispersed light) would provide beam combination. At λ > 8 µm, heterodyne receivers which cover the entire N-band by many (optically dispersed) sub-bands seem the most promising way to me. The VLTI could serve as testbed to develop the required techniques.

References 1. 2. 3. 4.

R. Blandford, M. Rees, Mon. Not. R. Astron. Soc. 169, 395 (1974) W. Jaffe, K. Meisenheimer, H. Röttgering et al., Nature 429, 47 (2004) K. Meisenheimer, K. Tristram, W. Jaffe et al., Astron. Astrophys. 471, 453 (2007) M. Schartmann, Models of dust and gas tori in active galactic nuclei, PhD thesis, RupertoCarolo University, Heidelberg (2007) 5. M. Swain, G. Vasisht, R. Akeson et al., Astrophys. J. 596, L163 (2003) 6. K. Tristram, K. Meisenheimer, W. Jaffe et al., Astron. Astrophys. 474, 837 (2007)

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The Sun and Space Weather (Astrophysics and Space Science Library)

Water in the Universe (Astrophysics and Space Science Library)

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Sirius Matters (Astrophysics and Space Science Library)

The Cluster Active Archive: Studying the Earth's Space Plasma Environment (Astrophysics and Space Science Proceedings)

High Energy Density Laboratory Astrophysics (Astrophysics and Space Science)

High Time Resolution Astrophysics (Astrophysics and Space Science Library)

Solar, Stellar and Galactic Connections between Particle Physics and Astrophysics (Astrophysics and Space Science Proceedings)

Organizations and Strategies in Astronomy (Astrophysics and Space Science Library)

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Polarization in Spectral Lines (Astrophysics and Space Science Library)

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