What impact do massive stellar associations have on their environment through the power of their ionising radiation and mechanical energy? Gathered together in this volume are papers presented at the first IAC-RGO meeting, held in Puerto Naos La Palma, Spain, dedicated to exploring all aspects of this burning question. This volume examines the formation and evolution of new clusters of stars, and explores all the possible consequences in a wide variety of objects where massive stellar bursts have occurred. Thus it presents an alternative model to that of supermassive black holes as the power houses of active galactic nuclei; it analyses the impact of Wolf-Rayet stars, stellar winds and supernovae on their host galaxy; and it provides evidence of massive superassociations and of supersonic velocity dispersions which result from photo-ionisation by violent star formation. For graduate students and researchers, this volume provides a valuable overview and a timely update on all aspects of violent star formation in a host of objects - from 30 Doradus, the largest HII region in the Large Magellanic Cloud, to nuclear starbusts and QSOs.
Violent Star Formation From 30 Doradus to QSOs
Violent Star Formation From 30 Doradus to QSOs Proceedings of the first IAC-RGO meeting, held in Puerto Naos La Palma, Spain November 8-12, 1993 Edited by G. TENORIO-TAGLE Instituto de Astrofisica de Canarias, Tenerife, Spain
CAMBRIDGE UNIVERSITY PRESS
Published by the Press Syndicate of the University of Cambridge The Pitt Building, Trumpington Street, Cambridge CB2 1RP 40 West 20th Street, New York, NY 10011-4211, USA 10 Stamford Road, Oakleigh, Melbourne 3166, Australia © Cambridge University Press 1994 First published 1994 Printed in Great Britain at the University Press, Cambridge A catalogue record of this book is available from the British Library Library of Congress cataloguing in publication data available ISBN 0 521 47277 6 hardback
Contents Participants Preface
xi xiii
Violent Star Formation in 30 Doradus R. C. Kennicutt & Y.-H. Chu
1
The Initial Mass Function of the Center of 30 Doradus E. M. Malumuth & S. R. Heap
13
The Nature and Kinematics of the Emission Nebulae in the Cyg OBI Supershell T. G. Sitnik & V. V. Pravdikova
19
Asymmetry in the Vertical Distribution of Giant Molecular Clouds in the Carina Arm E. J. Alfaro ei al
23
Supersonic Motions in Giant HII Regions C. Munoz-Tunon
25
A Kinematical Study of NGC 604 N. S. P. Sabalisck et al
39
UV Spectroscopy of Giant Extragalactic HII Regions: the Case of NGC 604 L. Sanz Fernandez de Cordoba
41
Evolution of GEHRs: The Effects Caused by Champagne Flows J. A. Rodriguez-Gaspar & G. Tenorio-Tagle
43
Bursts of Star Formation in Central Disks of Galaxies V. Surdin
45
Violent Star Formation G. Tenorio-Tagle
50
Superassociations as Star Complexes with Violent Star Formation Yu. N. Efremov Violent Star Formation Driven by Shock-Shock Collisions A. Chernin & Yu. N. Efremov
61 65
The Search for Hierarchical Structure inside M31 Superassociations P. Battinelli et al
68
A Stochastic PSF Model: Smooth Spirals in Differentially Rotating Disks J. Palous & B. Jungwiert
70
Spatiotemporal Pattern Driven by a Self-Regulating Mechanism of Star Formation A. Parravano
75
Detection of an Age Gradient along the z-Axis in a Star-Forming Region E. J. Alfaro et al
77
Abundances of HII Regions and the Chemical Evolution of Galaxies M. Peimberi ei al
79
vn
viii
Contents
Galaxy Properties in Different Environments: Star Formation in Bulges of LateType Spirals M. G. Pasioriza et al
94
Star Formation in Galaxies in the Bootes Void D. Weistrop
100
Physical Properties of Giant Extragalactic HII Regions A. I. Diaz
105
The Giant HII Region NGC 2363 R. Gonzalez-Delgado et al
117
Photometric Diagrams of NGC 2366 A. Aparicio et al
123
Spectroscopical Imaging of Star-Forming Regions J. M. Mas-Hesse et al
125
A Study on the HII Regions of NGC 4449 0. Fuentes-Masip et al Metallicity Effects on the Properties of Very Young Star Clusters M. L. Garcia- Vargas et al
131 133
The Prototype Starburst Galaxy NGC 7714: Physical Conditions of the Gas and the Stellar Populations R. Gonzdlez-Delgado et al 139 Tracing Violent Star Formation: HI Observations of Nearby Galaxies E. Brinks
145
Massive Star Formation and Supergiant Shells in the Irregular Galaxy NGC 55 D. J. Bomans & E. K. Grebel
156
Galactic Supershells 5. A. Silich el al
162
Violent Star Formation in Dwarf Irregular Galaxies E. Skillman
168
Very Metal-Poor Galaxies and the Primordial Helium Abundance E. Terlevich et al
182
Implications from HI Composition and Lya Emission of HII Galaxies D. Kunth et al
192
The IR- and X-Radiation of the Starburst Dwarf Galaxy UGCA 86 G. M. Richler el al
200
High-Resolution CCD Photometry of HII Galaxies E. Telles & R. Terlevich
202
Formation of Narrow Hell A4686 Emission in HII Galaxies: Link with X-Ray Emission C. Motch et al 208 Environmental Effects in Star-Forming Dwarf Galaxies J. M. Vilchez
214
Contents
ix
Theory of Starburst and Ultraluminous Galaxies B. Elemegreen
220
Colliding Galaxies, Shocked Gas, and Violent Star Formation S. A. Lamb et al
243
Violent Star Formation in Merger Remnants U. Frilze-von Alfvensleben
249
UV Variability of IRAS 13224-3809 J. M. Mas-Hesse et al
256
Star Formation in Polar-Ring Galaxies V. Reshetnikov & F. Combes
258
Infrared Spectroscopy of IR-Luminous Galaxies A. Siernberg et al
263
Application of the Multiphase Model to the Galactic Bulge M. Molld et al
268
Stellar Populations and Population Gradients in Spiral Bulges M. Balcells & R. F. Peletier
270
Implications of Galaxy Alignment for the Galaxy Formation Problem W. Godlowski
275
Annular Structure Analysis of the Starburst Spiral Galaxy NGC 7217 A. M. Varela et al
277
How a Dust Concentration Mimics Dynamical Signatures around the Nucleus of NGC 7331 F. Prada et al 279 UGC 5101: An Ultraluminous IRAS Galaxy with Circumnuclear Star Formation V. Reshetnikov & F. Combes
281
The Stellar Content of Nearby and Distant Starbursts C. Leitherer
283
WR Stars in the Giant HII Region NGC 4236III R. Gonzdlez-Delgado & E. Perez Consequences of High Mass Loss Rates on Wolf-Rayet Populations in Starbursts G. Meynet
289 291
Optical and Ultraviolet Morphology of the Starburst Regions in Wolf-Rayet Galaxies W. D. Vacca 297 The Starburst Nucleus of M83 5. R. Heap
303
Spectrophotometry of Haro Starburst Galaxies 5. Steel et al
309
Starbursts in the Irregular Galaxy VV 523 J. Hecquet et al
317
x
Contents
Long-Slit Spectroscopy of the Central Regions of Starburst Galaxies Henize 2-10 and Markarian 52 H. Sugai & Y. Taniguchi 319 Star Formation in Active Galactic Nuclei: the Cases of NGC 5135, NGC 6221 and NGC 7130 H. R. Schmili et al 325 Metallicity Effects on Starburst M. Cerviiio & J. M. Mas-Hesse
327
From 30 Doradus to QSOs R. Terlevtch
329
Distance Indicators to Low-Luminosity AGN /. Aretxaga & R. J. Terlevich
343
Broad- and Narrow-Band Imaging of the CfA Seyfert Sample A. M. Perez Garcia & J. M. Rodriguez Espinosa
345
Type Transitions in Starburst-Powered AGN /. Aretxaga & R. J. Terlevich
347
Stellar Ionization of Low-Luminosity Active Galactic Nuclei J. C. Shields
353
Line Profiles in Compact Supernova Remnants and Active Galactic Nuclei R. Cid Fernandes & R. Terlevich
365
Composite Galactic Nuclei B. Boer
377
The Nature and Origin of X-Ray Emission in Active Galaxies H. Netzer
379
Starbursts and Compact Supernova Remnants J. Franco et al
387
Broad-Band and Line Emission from Fast Radiative Shocks in Dense Media T. Plewa
396
Study of the Stellar Populations in AGN M. Serole-Roos et al
403
Bidimensional Spectroscopy of Seyfert Galaxies: Offset BLR in NGC 3227 S. Arribas & E. Mediavilla
405
ROSAT Detection of the Most Rapidly Varying Seyfert Galaxy Th. Boiler & J. Trumper
410
QSO Evolution: a Link with Starbursts? B. J. Boyle
413
Evolution of Elliptical Galaxies - a Chemo-Dynamical Model A. C. S. Friaca & R. J. Terlevich
424
Birth of Galaxies at z = 2 or Violent Star Formation at z = 0.4? A. D. Chernin
430
Participants ALFARO, E Instituto de Astroffsica de Andalucia, Spain APARICIO, A Instituto de Astrofi'sica de Canarias, Spain ARETXAGA, I Royal Greenwich Observatory, UK ARMAND, C Laboratoire D'Astronomie Spatiale, France ARR1BAS, S Instituto de Astroffsica de Canarias, Spain BALCELLS, M Groningen University, The Netherlands BALUTEAU, J P Observatoire de Marseille, France BENN, C Royal Greenwich Observatory, La Palma, Spain BOER, B Laboratory for Space Research, The Netherland BOLLER, T MPI fur Extraterrestrische Physic, Germany BOMANS, D J Sternwarte University Bonn, Germany BOYLE, B Institute of Astronomy, Cambridge, UK BRINKS, E National Radio Astronomy Observatory, USA CERVIO, M LAEFF, Spain CHERNIN, A Sternberg Astronomical Institute, Russia CID FERNANDES, R Institute of Astronomy, Cambridge, UK COLINA, L Universidad Autonoma de Madrid, Spain COUPINOT, G Observatoire Midi-Pyrinees, France DIAZ BELTRAN, A Universidad Autonoma de Madrid, Spain DOUGLAS, N Kapteyn Institute Groningen, The Netherlands EFREMOV, Y N Sternberg Astronomical Institute, Russia ELMEGREEN, B IBM Watson Research Center, USA FERLAND, G University of Kentuky, USA FRANCO, J Instituto de Astronomfa UNAM, Mexico FRITZE VAN ALVENSLEBEN, U Universitats Sternwarte, Gottingen, Germany FUENTES MASIP, O Instituto de Astrofi'sica de Canarias, Spain GARCIA VARGAS, M Universidad Autonoma de Madrid, Spain GERRITSEN, J Kapteyn Institute Groningen, The Netherlands GODLOWSKI, W Jagiellonian University Astronomical Observatory, Poland GONZALEZ DELGADO, R Instituto de Astrofi'sica de Canarias, Spain GRIFFITHS, R E John Hopkins University, USA HEAP, S R NASA/Goddard Space Flight Center, USA HECQUET, J Observatoire de Midi-Pyrinnes, France JOUBERT, M Laboratoire D'Astronomie Spatiale, France KENNICUTT, R Steward Observatory, USA KUNTH, D Institute D'Astrophysique Paris LAMB, S University of Illinois, USA LEITHERER, C Space Telescope Science Institute, USA MAIOLINO, R Osservatorio Astrofisico di Arcetri, Italy MALUMUTH, E Goodard Space Flight Center, USA MAS-HESSE, J M LAEFF, Spain MEDIAVILLA, E Instituto de Astrofi'sica de Canarias, Spain MENON, T K Max Planck Institute fur Radioastronomie, Germany MEYNET, G Observatoire de Geneva, Switzerlan MIRABEL, F CEN, Saclay, France MOITINHO, A Universidad de Lisboa, Portugal MOLLA, M Universidad Autonoma de Madrid, Spain MOTCH, C Observatoire de Strasbourg, France xi
Participants
Xll
MUNCH, G MUNOZ-TUNON, C NETZER, H PAKULL, M PALOUS, J PARRAVANO, A PASTORIZA, M PEIMBERT, M PELETIER, R F PEREZ, A PEREZ, E PEREZ OLEA, D PLEWA, T PRADA MARTINEZ, F PRIETO, M REIMERS, D RESHETNIKOV, V RICHTER, G M ROBERTS, W J RODRIGUEZ ESPINOSA, J M RODRIGUEZ GASPAR, J A SABALISCK, N SANZ FDEZ de CORDOBA, L SCHMITT, H R SEROTE, M SHIELDS, J SILICH, S SILLANPAA, A SITNIK, T SKILLMAN, E STEEL, S STERNBERG, A SUGAI, H SURDIN, V TELLES, E TENORIO-TAGLE, G TERLEVICH, E TERLEVICH, R VACCA, W VILCHEZ, J M WAGNER, S WEISTROP, D WHITE, S ZINNECKER, H
Instituto de Astrofi'sica de Canarias, Spain Instituto de Astrofi'sica de Canarias, Spain Tel Aviv University, Israe Observatoire de Strasbourg, France Astronomical Institute, Czechoslovakia Universidad de Los Andes, Venezuela Universidade Federal Rio Grande Do Sul, Brazil Instituto de Astrononna UN AM, Mexico Royal Greenwich Observatory, La Palma, Spain Instituto de Astrofi'sica de Canarias, Spain Instituto de Astrofi'sica de Canarias, Spain Universidad Autonoma de Madrid, Spain Warsaw University Observatory, Poland Instituto de Astrofi'sica de Canarias, Spain Instituto de Astrofi'sica de Canarias, Spain Hamburger Sternwarte, Germany Observatoire de Paris, France Astrophysical Institute Potsdam, Germany Space Telescope Science Institute, USA Instituto de Astrofi'sica de Canarias, Spain Instituto de Astrofi'sica de Canarias, Spain Instituto de Astrofi'sica de Canarias, Spain LAEFF, Spain Universidade Federal Rio Grande Do Sul, Brazil Observatoire de Paris-Meudon, France Ohio State University, USA Main Astronomical Observatory, Ukrania Turku University, Finland Sternberg Astronomical Institute, Russia University of Minnesota, USA University College Dublin, Ireland Tel Aviv University, Israe National Astronomical Observatory, Japa Sternberg Astronomical Institute, Russia Institute of Astronomy, Cambridge, UK Instituto de Astrofi'sica de Canarias, Spain Royal Greenwich Observatory, UK Royal Greenwich Observatory, UK University of California, Berkeley, USA Instituto de Astrofi'sica de Canarias, Spain Landssternwarte, Heidelberg, German University of Nevada, USA Institute of Astronomy, Cambridge, UK University Wurzburg, Germany
EDITOR'S PREFACE
Violent Star Formation from 30 Doradus to QSOs was the most recent international conference organized jointly by the Institute de Astrofisica de Canarias (IAC) and the Royal Greenwich Observatory (RGO). The meeting took place in Puerto Naos, La Palma (Canary Islands, Spain), in November 1993. This volume contains most of the invited talks and papers presented at the conference. It deals with observations, analysis and theory of violent star formation in its full dimension: from giant HII regions, dwarf and HII galaxies, to starburst galaxies, IRAS ultraluminous galaxies and interacting pairs, all the way up to violent star formation in active galactic nuclei and QSOs. Several critical reviews look at the implications of violent star formation from a variety of angles: ionization structure, interaction with the ISM, hydrodynamics, triggering mechanisms, chemical evolution, luminosity functions and the starburst model for AGNs. On behalf of the organizing committee I would like to thank all the staff and colleagues at the IAC and the RGO who contributed to the organization of the meeting. We are particularly grateful to Monica Murphy (IAC), Judith de Araoz (IAC), Silvia Figueroa (RGO-La Palma) and Rachel Miles (RGO-La Palma) for their efficiency and their keen collaboration to warrant a very pleasant and successful conference. I also want to express my gratitude to Terence Mahoney for his help with the preparation of each of the papers in this volume. The Local Organization was eased by the definite commitment from the Local Government through their Cabildo Insular de La Palma, and the Cabildo Insular de Santa Cruz de Tenerife. It is a pleasure to acknowledge sponsorship from "SOL Hoteles" and the "Grupo SOL", and the local savings bank "Caja General de Ahorros de Canarias", as well as the generous financial support from the Spanish General Directorate for Scientific and Technical Research (DGICYT) and the International Science Foundation (ISF) without which the celebration of the meeting and the attendance of many young colleagues would have been impossible.
Guillermo Tenorio-Tagle Instilulo de Astrofisica de Canarias, Tenerife Spain, April 1994
xin
Violent Star Formation in 30 Doradus By R. C. KENNICUTT 1 AND Y.-H. CHU2 Steward Observatory, University of Arizona, Tucson, AZ 85721, USA 2
Department of Astronomy, University of Illinois, Urbana, IL 61801, USA
The 30 Doradus nebula is the nearest example of a giant HII region, and as such it offers a unique laboratory for studying in detail the structure, stellar content, and dynamics of a starburst region. We begin with an overview of the 30 Doradus region on scales of 0.1-1000 pc, and then discuss two current problems of particular relevance to this conference, the stellar content and IMF in 30 Dor, and the violent dynamics of its interstellar medium.
1. Introduction It is a pleasure to open a conference where 30 Doradus defines the bottom end of the star formation scale! The 30 Doradus region offers a most appropriate starting point for a conference on star formation in galaxies. It is the nearest example of a bona fide giant extragalactic HII region (GEHR), and it is the largest star forming region in the Local Group. It is large enough to exhibit many of the properties of the most luminous starbursts, yet close enough so that its physical structure and stellar content can be studied in detail. As such 30 Dor and other nearby GEHRs provide several crucial pieces of information about starbursts in general. They are the only regions where the embedded stellar population can be resolved and studied directly; this provides a unique stellar census of a starburst, which can be used to test the synthesis models which must be applied to more distant, unresolved GEHRs and starbursts. Imaging and spectroscopy of the gas provides a detailed picture of the physical conditions in the interstellar medium (ISM), hence these regions are the ideal laboratories for studying the complex dynamical interactions between massive stars and the surrounding ISM. Such feedback processes shape the large-scale structure of the ISM, and may well regulate the global evolution of galaxies. 30 Doradus offers all of the ingredients—a highly resolved concentration of thousands of massive stars in a dense, bright, dynamic, ionized medium—to serve as a "starburst Rosetta Stone" (Walborn 1991). Excellent reviews of 30 Doradus have been published recently by Melnick (1987), Walborn (1991) and Meylan (1993), and we refer the reader to those references for a comprehensive review of recent work. In this paper we begin with a brief overview of the 30 Dor region (Section 2), then discuss two specific areas which are especially relevant to this conference, the stellar content and IMF in 30 Dor (Section 3) and the dynamics of its violent ISM (Section 4). 2. 30 Doradus as a prototype giant HII region A useful definition of a giant HII region is any region which contains a sufficient number of massive stars so that its integrated spectra, physical properties, and evolution are dictated by the composite properties of its stellar population (e.g. total stellar mass, IMF, metallicity), rather than the individual properties of its most massive stars, or the local initial conditions in its parent star forming region. By this criterion 30 Doradus easily qualifies as a GEHR, because it contains so many stars that its IMF is fully populated to >100 M@. In this section we describe the 30 Dor region on different spatial scales, and compare it to other nearby star forming regions. 1
Kennicutt k Chu: Violent Star Formation in 30 Doradus
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FIGURE 1. 30 Doradus and the star cluster NGC 2070. The top panels show a 23' (345 pc) region in Ha and red continuum. The bottom panels are expanded photographs of the central star cluster. Boxes indicate the sizes of subsequent panels. From Kennicutt & Chu (1988).
2.1. Overview of the 30 Doradus region
Figures 1 and 2 show Ha and continuum images of the 30 Doradus region on various scales. The central HII region (top panels in Figure 1) is roughly 15-20 arcmin or 200300 pc across (assumed distance 50 kpc). Projected within the boundary of the nebula are several star clusters and associations, including the central cluster NGC 2070. Figure
Kennicutt & Chu: Violent Star Formation in 30 Doradus
3
2 is a deeper Ha image of a 1600 pc region around 30 Dor, and it reveals that the nebula is located in the midst of a much larger superassociation. Two large HII/OB complexes, 30 Dor B (visible to the SE in Figure 1) and the supershell 30 Dor C, are located in the periphery of 30 Dor and are often considered to be part of the main HII region. Over 20 other large HII regions, OB associations, and star clusters surround this region, mostly to the south of 30 Dor. The ages of these stars range from 0 Myr (including obscured protostellar regions in N159 and 30 Dor itself) to 10-12 Myr (Melnick 1992). Two large expanding supergiant shells, LMC-2 and LMC-3, surround the eastern and western parts of the region, respectively (Meaburn 1980). A faint halo of Ha emission envelops the entire 1 kpc region (Kennicutt k. Hodge 1986); if observed from a large distance the whole complex might well be treated as a single supergiant HII region. The picture of the 30 Dor region which emerges from these data is quite different from the simple single-age stellar population model which is usually applied to more distant GEHRs. Stellar photometry (e.g. Melnick 1987; Lortet & Testor 1991; Schild k Testor 1992; Parker 1993) shows that the region is actually a complex composite of several distinct stellar populations. Even the central nebula contains subclusters with ages ranging over at least 10 Myr. This means that the integrated spectrum of the region (or any similar GEHR) will be dominated by different stellar components at different wavelengths. While the emission-line spectrum of the photoionized gas will trace the youngest stellar population (<3 Myr), the underlying near-UV and visible continua may well be dominated by evolved supergiants from the older population (Searle 1971). The bottom two panels of Figure 1 show enlarged views of the stars in the nebular core. The lower left panel (5.8 arcmin across) is comparable to the region which has been surveyed by Parker (1993), and contains over 2000 stars down to V = 18, not including the compact core cluster around R136. The bottom right panel shows the central 3 arcmin; the unresolved subcluster in the center is R136. It is interesting to note that at a distance of 15 Mpc the bottom panels in Figure 1 would subtend 1.2 and 0.6 arcsec, respectively! Even with the Hubble Space Telescope these regions are unresolvable beyond the Local Group. This emphasizes the importance of objects like 30 Dor for understanding the nature of starbursts on all scales. 2.2. Integrated properties and comparison with other objects
The integrated properties of 30 Dor were tabulated by Kennicutt (1984), and these provide a more quantitative picture of a prototypical GEHR. The extinction-corrected Ha luminosity within a diameter of 370 pc is 1.5 x 1040 ergs s" 1 , which corresponds to an ionizing luminosity of 1.1 x 1052 photons s" 1 . The Ha flux within R < 4' (60 pc) is roughly 1/3 of this total, and is consistent with the ionizing luminosity expected from the ~2400 OB stars within that area (Parker 1992). The total mass of ionized gas in the nebula is roughly 8 x 105 M©. These values are 2-3 orders of magnitude larger than well-known HII regions such as Orion, and several-fold larger than W49 and NGC 3603, the largest HII regions in the Galaxy (Kennicutt 1984). The difference between the LMC and Galactic populations is consistent with a general trend along the Hubble sequence. Although 30 Dor is unrivaled in the Galaxy, it is similar to the largest HII regions in other Magellanic irregulars of comparable mass, and several times less luminous than the brightest GEHRs in giant Sc galaxies such as M51 and M101 (Kennicutt 1988). In addition to the differences in absolute scale, the physical conditions in the GEHRs are quite distinct from those in smaller, compact Galactic regions. The electron density ranges from ~100-300 cm" 3 in the densest filaments to of order 1-10 cm" 3 in the other parts of the nebula (Kennicutt 1984; Mathis, Chu, & Peterson 1985; Rosa & Mathis 1987). This compares to typical densities of order 102 - 104 cm" 3 in compact HII
Kennicutt & Chu: Violent Star Formation in 30 Doradus
FIGURE
2. Ha image of the region surrounding 30 Dor, from Kennicutt & Hodge (1986).
regions. The visual extinction over most of 30 Dor is low, only 0-2 mag (Parker 1993; Dickel et al. 1994), though a few dusty condensations with protostars have been identified (Hyland et al. 1992). As will be discussed extensively at this conference, the ionized gas in GEHRs is much more turbulent, with supersonic velocity dispersions that are 2-3 times higher than in smaller HII regions. These differences are probably related to the concentration of massive stars, which facilitates the transfer of mechanical energy to the ISM from multiple stellar winds and supernovae. We return to this point in Section 4.
Kennicutt & Chu: Violent Star Formation in 30 Doradus
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FIGURE 3. H-R diagram for the central 7' region of 30 Doradus, from Parker & Garmany (1993). Sample includes 1230 stars with low photometric errors.
3. Stellar content and IMF As mentioned earlier the proximity of 30 Doradus provides a unique opportunity to measure the stellar population and the form of the initial mass function (IMF) in a starburst region. Numerous authors have resorted to various forms of biased IMFs— top-heavy IMFs, top and bottom-deficient IMFs, or bimodal IMFs— when attempting to account for the extraordinary luminosities of GEHRs and starbursts. The 30 Doradus region is an ideal testing ground for these biased IMF models. Ground-based UBV photometric surveys have been made by several investigators (Lee 1990; La Pierre & Moffatt 1991; Melnick 1992; Meylan 1993; Parker 1993), and HST observations of the core around R136 have been published by Campbell et al. (1992) and Malumuth & Heap (1994). The most complete analysis to date is the thesis study of Parker. He surveyed a 50 arcmin2 region in the center of 30 Dor, as part of a larger survey of the LMC by the Boulder group. CCD photometry in UBV was obtained for a complete sample of 2395 stars brighter than B = V — 18, and used to produce a reddening corrected H-R diagram. Classification spectra were obtained or compiled for ~140 stars brighter than V = 15.5, to provide accurate effective temperatures and bolometric corrections for the hottest and most massive stars. This provided a complete census of the upper end of the IMF for the
6
Kennicutt k Chu: Violent Star Formation in 30 Doradus
center of 30 Dor, except for a few Wolf-Rayet stars and the crowded region immediately around R136 (see below). The final results of the analysis are summarized in Figure 3, which displays the H-R diagram for 1230 stars with the highest quality observations (Parker k Garmany 1993). Evolutionary tracks are superimposed on the data, and the present-day mass function can be derived simply by binning the stars in the respective mass intervals. Parker k Garmany find that the mass function above 12 MQ is well represented by a power law with slope F = d(\ogN)/d{\ogm) = -1.5 ± 0.2. If the cluster is younger than ~3 Myr, as appears to be the case, this PDMF represents the IMF in the region. The slope is considerably shallower than found for the solar neighborhood (F = —2.2) or the LMC outside 30 Dor (F = -1.9) by Blaha k Humphreys (1989), and it is shallower than the mean slope (F = —2.0) in 14 smaller LMC associations measured by Hill, Madore, & Freedman (1994). RecentffST observations allow this analysis to be extended to the dense stellar core of 30 Dor, around the central object R136 (Campbell et al. 1992; Malumuth k Heap 1994). As reported by Malumuth at this conference, the IMF outside of R136 itself (R > 3'-'3 or 0.8 pc) is consistent with the Parker and Garmany results (F = —1.8), but the slope becomes much shallower in the R136 core itself, with F = —0.9 for R < 0.8 pc (Malumuth k Heap 1994). For details we refer the reader to the paper by Malumuth k Heap in this volume. These results taken together lend some support for a top-heavy IMF in GEHRs relative to smaller star forming regions. However it is interesting to note, as pointed out by Parker k Garmany, that previous measurements of the IMF in 30 Dor from UBV photometry (see references above) range over F = —0.8 to —2.0 in roughly the same region! Many of the existing data sets clearly are influenced by systematic errors, and complete homogeneous surveys are needed to test for systematic IMF variations, even in relatively accessible regions such as the Galaxy and the Magellanic Clouds. Highresolution infrared observations, which could extend the mass coverage from ~10 M© (the completeness level of most current surveys) to <1 MQ would be especially valuable (Zinneker 1994, this volume). As long as there is uncertainty in the IMFs determined for highly resolved GEHRs such as 30 Dor, any claims of anomalous IMFs in more distant, unresolved GEHRs and starbursts should be approached with a healthy degree of skepticism. 4. The Violent interstellar medium in 30 Doradus The chaotic filamentary structure of 30 Doradus testifies to its violent internal dynamics. Early Fabry-Perot measurements revealed a supersonic velocity dispersion of 41 km s" 1 FWHM, much higher than those for other HII regions in the Galaxy or Magellanic Clouds (Smith k Weedman 1972), while consistent with a general correlation between linewidth and nebular luminosity among GEHRs (Melnick 1977). High-resolution slit spectra resolved these motions into expanding shells on scales of 1-100 pc (Meaburn 1984, 1988). The detection of several extended X-ray sources in 30 Dor by Einstein (Long, Helfand, k Grabelsky 1981; Chu k Mac Low 1990; Wang k Helfand 1991a) revealed hot gas which is presumably shocked by stellar winds and supernovae embedded in the nebula. These observations point to a dynamical environment which is profoundly different from that in smaller HII regions. In a typical Galactic disk region massive stars form in relatively small associations, and at any particular location the time interval between successive Wolf-Rayet stars or supernovae is usually longer than the lifetime of an in-
Kennicutt & Chu: Violent Star Formation in 30 Doradus
FIGURE 4. Mosaic of 37 echelle line profiles in Ha, covering the central 9' x 9' in 30 Doradus, from Chu fc Kennicutt (1994). Orientation is N top, E left. Each 9' spectrum is separated by 45" N-S. Velocity range in each spectrum is ±300 km s The continuum source near the center is R136.
dividual interstellar bubble or supernova remnant (SNR). In 30 Doradus the situation is dramatically different, with several hundred or thousand OB stars located within a radius of ~l-20 pc. In such environments the surrounding ISM will be influenced by the collective stellar winds from the association, and subsequent supernovae will erupt within large wind-blown bubbles or previously formed SNRs. Such regions are the likely seed environments for the development of the large interstellar bubbles, supershells, and superwinds in galaxies. The ultimate effect of such a GEHR on its ambient ISM is likely to be quite different from an environment where the same number of massive stars are distributed in small, spatially isolated regions.
8
Kennicutt & Chu: Violent Star Formation in 30 Doradus 4.1. Kinematic structure of 30 Doradus
The 30 Doradus region offers a unique opportunity to study these collective interactions between massive stars and the ISM, and the early development of a supergiant shell. As part of a larger survey of the dynamics of the HII regions and SNRs in the Magellanic Clouds, we have mapped the kinematics of 30 Dor using the echelle spectrograph on the CTIO 4-m telescope in a single-order, long-slit mode (Chu et al. 1992; Chu &, Kennicutt 1994). To illustrate the kinematics of the ionized gas in the central 9' x 9' core, Figure 4 shows a mosaic of 37 E-W observations of the Ha line, separated by 3' in the E-W direction and 45" N-S. The velocity range in each individual spectrum is ±300 km s" 1 . R136 is visible as the bright continuum feature near the center of the mosaic. Similar observations were obtained at 53 other positions outside the core. For several regions of interest, echellograms of the [SII] doublet were observed for shock diagnostics and density measurements. The echellograms in Figure 4 reveal the extraordinary complexity of the violent motions in the center of 30 Dor. The outer parts of the HII region (mostly outside the region shown in Figure 4) are characterized by a relatively smooth velocity field, but with a much higher velocity dispersion than in other HII regions in the LMC. As one approaches the center of the nebula the velocity field becomes more complex, with multiple-velocity components at all positions. Several distinct types of kinematic features have been identified (Meaburn 1984; Chu & Kennicutt 1994). These include slow expanding shells with Wexp < 100 km s" 1 and sizes from a few to ~100 pc, and fast expanding shells with Vea;p = 100-300 km s" 1 and sizes from a few to ~30 pc. There are also discrete high-velocity (AV > 100 km s"1) features; some of them are isolated, while others are concentrated within large slow expanding shells, forming complex expanding networks. Several fast expanding shells and complex expanding networks are evident in Figure 4. The slow expanding shells are commonly seen in other regions in the Galaxy or the Magellanic Clouds. The smaller objects (<15 pc) are probably stellar wind-blown bubbles, since the energetic requirements can be met with winds from one W-R star or a small number of OB stars. The larger slow shells (e.g. 30 Dor C outside the main nebula) can be accelerated by a combination of stellar winds and supernovae. The most massive of these, a 20 pc shell expanding at 40 km s~l from the central star cluster, requires an impressive input energy (~ 1052 ergs), but this is not unreasonable given the concentration of massive stars in this region. The fast expanding shells are the most spectacular kinematic features in 30 Doradus, and are only rarely seen elsewhere. No known wind-blown bubbles can match such high expansion velocities. The shell masses range from 100-5000 MQ, and the kinetic energy in the visible ionized shells ranges over 1049 — 1051 ergs. The energy required to produce these shells, whether from stellar winds or supernovae, is roughly an order of magnitude higher, or 1O50 — 1052 ergs. Stellar winds are a less plausible energy source for the fast shells, as the large energy requirement and short expansion ages (104 —105 yr) are usually inconsistent with the number of embedded stars. These fast shells most likely are SNRs inside 30 Dor. The complex expanding networks are unique to 30 Doradus (within the LMC). Each network consists of a large slow expanding shell with additional high-velocity, shocked material within the shell boundary. This structure can be naturally explained by a winddriven superbubble, with SNRs interacting with its shell walls. Chu & Mac Low (1990) have proposed the same mechanism to explain X-ray bright superbubbles in the LMC. Indeed, bright X-ray emission has been detected within the 30 Dor networks, as well as in the aforementioned fast shells.
Kennicutt k Chu: Violent Star Formation in 30 Doradus
9
4.2. Soft X-ray properties of 30 Doradus and surroundings Soft X-ray observations are useful in tracing regions of fast shocks and delineating bubbles of hot gas in a large HII complex. We have obtained a ROSATPSFC pointed observation of the 30 Dor region (Figure 5). The top panel shows the entire 30 Dor superassociation. The brightest sources are the X-ray binary LMC X-l, the known SNRs 30 Dor B and 0540-69.3, and two Galactic foreground stars. If the LMC were observed from a large distance these sources would be unresolved, and a significant fraction of the total flux would be contributed by LMC X-l and the SNRs. Variable interstellar absorption at these low energies (0.1-2.4 keV) further complicates the analysis. This illustrates the difficulty inherent in interpreting X-ray observations of distant GEHRs and starbursts. Figure 5b (bottom) shows a close-up, smoothed X-ray image centered on 30 Dor itself. Within 30 Dor, two types of X-ray sources are visible: point sources coincident with the stellar groupings R136 and R140, and extended diffuse emission associated with expanding shell structures. The fast expanding shells and networks identified in the echelle data have high X-ray surface brightness, while the shells without high-velocity features are much fainter in X-rays. The diffuse X-ray sources in 30 Dor can be explained by SNRs interacting with preexisting superbubbles (Chu & Mac Low 1990; Wang & Helfand 1991b). One such example is 30 Dor C, a superbubble around the OB association LH90; it has a limb-brightened shell structure in the X-ray image (Figure 5a). The SNR activity in 30 Dor C has been unambiguously confirmed by its nonthermal radio emission (Mathewson et al. 1985). A superbubble with SNRs interacting with its shell walls can emit 1035 to several times 1036 ergs s" 1 X-rays (Chu & Mac Low 1990; Chu et al. 1993). The diffuse X-ray emission from 30 Dor, at a level of ~ 2 x 1037 ergs s"1 (Wang & Helfand 1991a), can be decomposed into 5-6 X-ray bright superbubbles (Chu et al. 1994, in preparation). The X-ray observations reveal the prevalent SNR activities in 30 Dor, but also illustrate that the 3-dimensional structure of 30 Dor is dominated by several large shells. 4.3. Integrated kinematic properties We have used the uniform coverage of our echelle data in the central 9' core to estimate the total amounts of energy in the various types of velocity structures, and constrain the physical origin of the supersonic motions. When the echellograms in Figure 4 are summed to produce a single pseudo-integrated profile, the resulting profile has a remarkably smooth structure, as shown in Figure 6. This suggests that the roughly Gaussian velocity structure which is seen in the integrated profiles of more distant GEHRs (e.g. Melnick 1977) is not necessarily the result of virial motions, but rather is produced by the superposition of hundreds of smaller structures which are accelerated by a combination of stellar winds, supernovae, champagne flows, and truly turbulent motions. A point-by-point analysis of the velocityfieldin the 30 Doradus core appears to support this conclusion (Chu & Kennicutt 1994). The background smooth turbulent velocity field is the largest single contributor to the total kinetic energy of the nebula, with ~ 6 x 1051 ergs. The high-velocity shells (Vexp > 100 km s"1) and slower shells each contain approximately 3 — 4 x 1051 ergs, with the shell around R136 alone contributing about 30% of the latter. It is interesting to note that although the high-velocity gas in 30 Dor contains a considerable fraction of the nebular kinetic energy, it represents an insignificant fraction of the integrated line emission. Most of the shells have emission measures of order 100-5000 cm"6 pc, in contrast to 104 — 106 for the foreground nebulosity. The total kinetic energy in the gas, ~ 1052 2 ergs in the central 9', exceeds by at least an order of magnitude the gravitational binding energy of the complex, and suggests that most of 30 Dor, with the possible exception of the core star cluster, is literally blowing
10
Kennicutt & Chu: Violent Star Formation in SO Doradus
star
30 Dor C
FIGURE 5. The 30 Doradus region as observed by the ROSAT PSPC. The top panel shows a 80' x 60' field. The bottom panel is a smoothed expanded plot of 30 Dor itself, covering 42' x 32'. The point sources near the center are coincident with R136 and R140.
Kennicutt k Chu: Violent Star Formation in 30 Doradus 1000
-1—i—1—r
i—i
n—i—r
11
.800
.2 CO o 600 3
I" 400
200
0 -300
, , I -200
-100
0 Velocity (km/s)
100
200
300
FIGURE 6. Ha velocity profile of the central 9' core of 30 Doradus, derived by summing the 37 echellograms in Figure 4. The dashed line shows the profile expanded by a factor of 20, to show the line wings. The peak at -270 km s"1 is geocoronal Ha.
itself apart. On the other hand these energy requirements are in rough accord with the power available from stellar winds and supernovae, as expected if the latter are ultimately responsible for most of the gas motions. For that reason we remain skeptical whether gravitational motion can account for the observed kinematics of GEHRs, though it may contribute to some extent to the lower velocity "turbulent" motions. The mechanical energy in 30 Dor itself is still factors of 5-10 lower than that required to produce the largest supergiant shells in the LMC, such as the large supershell LMC-4 in Shapley Constellation III (Dopita, Mathewson k Ford 1985). However the supergiant shells have characteristic dimensions of 1 kpc (Meaburn 1980), and if one includes the wind and supernova power from the entire 30 Dor superassociation (Figure 2) over a time scale of ~ 107 yr, the energies are comparable (~ 1053 ergs), suggesting that the 30 Dor region will likely form a similar supershell over the next 10-30 Myr. We are very grateful to Eliot Malumuth and Sally Heap for providing preliminary results prior to publication, and to Joel Parker for permission to reproduce Figure 3. RCK is supported in part by the U.S. National Science Foundation through grant AST9019150. YHC acknowledges the support of NASA Grants NAG 5-1900, NAG 5-2112 and GO-4497.02-92A.
BLAHA,
C. &
HUMPHREYS,
REFERENCES R. M. 1989 Astron. J. 98, 1598.
12
Kennicutt & Chu: Violent Star Formation in 30 Doradus
B. ET AL. 1992 Astion. J. 104, 1721. CHU, Y.-H. & KENNICUTT, R. C. 1994 Astrophys. J. In press. CHU, Y.-H., KENNICUTT, R. C , SCHOMMER, R. A. & LAFF, J. 1992 Astion, J. 103, 1545. CHU, Y.-H. & MAC LOW, M.-M. 1990 Astrophys. J. 365, 510. CHU, Y.-H., MAC LOW, M.-M., GARCIA-SEGURA, G., WAKKER, B. & KENNICUTT, R. C. 1993 Astrophys. J. 414, 213. CAMPBELL,
DICKEL, J. R., MILNE, D. K., KENNICUTT, R. C , CHU, Y.-H. & SCHOMMER, R. A. 1994
Astron. J. In press. DOPITA, M. A., MATHEWSON, D. S. & FORD, V. L. 1985 Astrophys. J. 297, 599. HILL, R. J., MADORE, B. F. & FREEDMAN, W. L. 1994 Astrophys. J. In press. HYLAND, A. R., STRAW. S., JONES, T. J. & GATLEY, I. 1992 Mon. Not. Roy. Astron. Soc. 257, 391. KENNICUTT, R. C. 1984 Astrophys. J. 287, 116. KENNICUTT, R. C. 1988 Astrophys. J. 334, 144. KENNICUTT, R. C. & CHU, Y.-H. 1988 Astron. J. 95, 720.
R. C. & HODGE, P. W. 1986 Astrophys. J. 306, 130. LA PIERRE, N. & MOFFAT, A. F. J. 1991 In The Formation and Evolution of Star Clusters (ed. K. Janes). ASP Conference Series, vol. 13, p. 155. LEE, M. G. 1990 PhD thesis, University of Washington. LONG, K. S., HELFAND, D. J. & GRABELSKY, D. A. 1981 Astrophys. J. 248, 925. LORTET, M.-C. & TESTOR, G. 1991 Astro/i. Astrophys. Suppl. 89, 185. MALUMUTH, E. M. &C HEAP, S. R. 1994 Astron. J. In press. KENNICUTT,
MATHEWSON, D. S., FORD, V. L., TUOHY, I. R., MILLS, B. Y., TURTLE, A. J. & HELFAND,
D. J. 1985 Astrophys. J. Suppl. 58, 197. MATHIS, J. S., CHU, Y.-H. & PETERSON, D. E. 1985 Astrophys. J. 292, 155. MEABURN, J. 1980 Mon. Not. Roy. Astron. Soc. 192, 365. MEABURN, J. 1984 Mon. Not. Roy. Astron. Soc. 211, 521. MEABURN, J. 1988 Mon. Not. Roy. Astron. Soc. 235, 375. MELNICK, J. 1977 Astrophys. J. 213, 15. MELNICK, J. 1987 In Observational Evidence of Activity in Galaxies (ed. E. Ye. Khachikian, K. J. Fricke & J. Melnick), p. 545. Springer. MELNICK, J. 1992 In Star Formation in Stellar Systems (ed. G. Tenorio-Tagle, M. Prieto k F. Sanchez), p. 253. Cambridge Univ. Press. MEYLAN, G. 1993 In The Globular Cluster-Galaxy Connection (ed. G. H. Smith & J. P. Brodie). ASP Conference Series, vol. 48, p. 588. PARKER, J. W. 1992 PhD thesis, University of Colorado. PARKER, J. W. 1993 Astron. J. 106, 560. PARKER, J. W. & GARMANY, C. D. 1993 Astron. J. 106, 1471. ROSA, M. & MATHIS, J. S. 1987 Astrophys. J. 317, 163. SCHILD, H. & TESTOR, G. 1992 Astron. Astrophys. Suppl. 92, 729. SEARLE, L. 1971 Astrophys. J. 168, 327. SMITH, M. G. & WEEDMAN, D. W. 1972 Astrophys. J. 172, 307. WALBORN, N. R. 1991 In Massive Stars in Starbursts (ed. C. Leitherer, N. R. Walborn, T. M. Heckman & C. A. Norman), p. 145. Cambridge Univ. Press. WANG, Q. & HELFAND, D. J. 1991a Astrophys. J. 370, 541. WANG, Q. & HELFAND, D. J. 1991b Astrophys. J. 373, 497.
The Initial Mass Function of the Center of 30 Doradus By ELIOT M. MALUMUTH 1 !, AND S A R A R. HEAP 2 'Astronomy Programs, Computer Sciences Corporation 2
The Laboratory for Astronomy and Solar Physics, Goddard Space Flight Center
We report on new Planetary Camera observations of the central region of 30 Doradus. These PC images are the first "deep" HST exposures of 30 Doradus that have appropriate photometric calibration. With R136a at the center of the PC6 CCD chip, the image reveals over 800 stars in a 35"x35"area, and over 200 stars within 3'.'3 of the center of R136a. We used the PSF-fitting method of Malumuth et al. (1991) to measure the magnitudes of all detected stars on the PC6 chip. We used these new B magnitudes, along with U and V magnitudes derived from archived PC images, to derive the luminosity function, mass density profile, and Initial Mass Function of the 30 Doradus ionizing cluster. We find that the mass distribution is like that of a King model, with a core radius of O'/96 (0.24 pc), and a total mass of 17,000 solar masses. Both the luminosity function and the IMF show evidence for mass segregation in the sense that the central region has a higher fraction of massive stars than the outer region of 30 Doradus.
1. Introduction 30 Doradus is one of the most interesting and important objects in the nearby part of the Universe. Walborn (1991) goes as far as to call the 30 Doradus region of the LMC a Rosetta stone for the interpretation of similar, more distant regions. It is no coincidence that 30 Doradus was among the first objects observed with the Hubble Space Telescope. And from the very first HST images, our view of the 30 Dor ionizing cluster began to change. The dense inner core of the cluster, R136a, once thought to be a single supermassive star (Savage et al. 1983), was seen to be a densely packed cluster of several hundred bright, but essentially normal, stars (Malumuth et al. 1991; Weigelt et al. 1991). While much ground based work on the 30 Dor cluster has been done recently (e.g. Parker & Garmany 1993), it is only with HST observations that the nature of the central region can be probed. Indeed, recent papers by the members of the WFPC team (Campbell et al. 1992) and the FOC team (De Marchi et al. 1993) have been published or are awaiting publication. This work is based on a PC image obtained in January 1992. This image is the first deep image of 30 Dor to have a proper flat-field calibration. 2. Photometry The data were obtained using the PC6 chip of the PC and the F439W (hereafter, B-band) filter. Four 180-s and one 30-s images were obtained and combined into a single deep image. The final image is five times deeper than any of the images reported by Campbell et al. (1992). Deep (/-band (F368M) and V-band (F547M) images were obtained as part of the Science Assessment Observation (SAO) program in August 1990. These image were obtained prior to the first WFPC UV flood procedure, so there are no appropriate in-orbit flat-field images. A comparison of our photometry with that of Campbell et al. (1992) demonstrates that it is important to have a deep image in t Address: Code 681, Goddard Space Flight Center, Greenbelt, MD 20771. 13
Malumuth & Heap: The IMF of the Center of SO Doradus
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1. Comparison of our V-band photometry to that of Campbell et al. (1992) for stars inside R136 (central 7'.'3 x 8'.'7).
order to do accurate photometry in the densest regions, even for the brightest stars. For stars outside of R136 there is no systematic difference, except for a small 0.05-mag color difference, in our photometry of the SAO V-band image and the Campbell et al. V-band photometry. However, for the stars within R136, a 7'.'3 x 8'/7 area, there is a strong systematic difference in the sense that the Campbell et al. magnitudes are smaller (see Figure 1). This difference grows with decreasing brightness, and is what would be expected from an image which is not very deep. Campbell et al. overestimate the brightness of the stars they do detect because of the light contributed by the fainter stars they do not detect. Our photometry suffers from a similar effect; however experiments with test images at the same signal-to-noise ratio as our 5-band image indicate that this bias will be small for stars brighter than about 16.5 mag. For that reason we will use 16.5 mag as a faint end cut-off for determining the luminosity and initial mass functions.
3. Luminosity function In order to compare the properties of the dense inner region of 30 Dor with the less dense outer region we divided our sample in two. All of the stars which are within 3'/3 of R136al comprise the inner sample, while those further away make up the outer sample. For the 5-band image, there are 213 stars in the inner sample and 629 stars in the outer sample. Figure 2(a) shows the luminosity function for the stars in the outer sample. A power law, slope= -0.83, has been fitted to the luminosity function between B =1 2.5 mag and B = 16.5 mag. The fit is displayed in Figure 2(o) as the straight line. The luminosity function for the inner sample is shown in Figure 2(6). Again a power law has been fitted to the data between B = 12.5 mag and B = 16.5 mag, and is displayed as a straight line. The slope of the luminosity function for the inner sample is flatter (slope= -0.66) than that of the outer sample, indicating that there is a larger percentage of bright stars in the inner sample than in the outer sample. Using the
Malumuth & Heap: The IMF of the Center of 30 Doradus 30 Doradus Outer Luminosity Function
15
-0.83
CD
9 1
30 Doradus Inner Luminosity Function -0.66
cn 9 1
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FIGURE
22
2. Comparison of the B-band luminosity function (a) for stars outside of Rl36a (17'.'5 > R > 3'.'3), and (6) for stars within R136a (R <3'.'3).
Kolmogorov-Smirnov test, which has the advantage that the data are not binned, we can reject the hypothesis that these distributions are the same at the 98% confidence level.
4. Initial mass function Our goal of this work is to determine the initial mass function (IMF) for the center of the 30 Doradus cluster. Since it is not feasible to obtain spectra of more than a few stars in R136, we must rely on evolutionary models of the cluster. First, we convert the observed B magnitudes to absolute magnitude, MB, using a distance modulus m — M = 18.57 (Panagia et al. 1991) and an average reddening E(B — V) = 0.44 (Parker k Garmany 1993). Next, we assume a cluster age of 3 Myr. The presence of Wolf-Rayet stars both within and outside R136a indicate an age greater than 2.8 Myr, while the lack of late O-type supergiants indicate an age less than 5 Myr. Finally, we use the program CLUSTMOD (Landsman 1991) to calculate MB for stars with a range of zero-age masses at 3 Myr using Maeder's (1990) evolutionary tracks for Z = 0.005 and Kurucz' (1991) flux distributions for stars with low Z. Each star is assigned a zero-age mass based upon its magnitude MB . The seven Wolf-Rayet stars (3 in the inner sample and 4 in the outer
Malumuth k Heap: The IMF of the Center of 30 Doradus
16
7.5 T = -1.82
5.5
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1
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Outer IMF
1
1
T = -0.90
Inner IMF
u a. 7.0 o
6.5 1.2
1.4
1.6 Log Mass (Mo)
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FIGURE 3. Comparison of the IMF for (a) stars outside Rl36a, and (6) stars within Rl36a (R <3'.'3). In each figure, the solid line represents a power-law fit to the data between B = 12.5 mag (~70 MQ) and B = 16.5 mag (~20 MQ)- The slope of the IMF within Rl36a is much flatter than outside. sample) are put on the high-mass branch, while all of the other stars are placed on the lower-mass branch. For stars which are not Wolf-Rayet stars, the upper mass cutoff is about 70 M Q , however there are few stars more massive than 60 M Q . The resulting IMFs for the outer and inner samples are shown in Figures 3(a) and (6), respectively. The IMF, denoted by £, is denned as the number of stars per logarithmic mass interval per unit area. The slope of the IMF is denned as
where the standard (Salpeter 1955) IMF has a slope, T = -1.35. The slope of the IMF for the outer sample is T = -1.82±0.41. This is in keeping with the slope found by Parker & Garmany (1993), F = - 1 . 5 , for stars in 30 Dor more massive than 12 M 0 , using ground
Malumuth & Heap: The IMF of the Center of 30 Doradus
-1.0
•0.5
0.0
17
0.5
Log r (pc) 4. The radial profile of the surface mass density for the central 17'.'5 of the 30 Doradus cluster. For comparison, the solid line shows a King model with a core radius of 0.24 pc.
FIGURE
based data. The slope derived for the inner sample is much flatter, T = -0.90 ± 0.38. There is a larger percentage of very massive stars within R136a. Again the KolmogorovSmirnov test indicates that we can reject the hypothesis that these distributions are the same at the 95.4% confidence level. Most of the excess massive stars are located right at the very center of the cluster. The IMF for stars in the central ^.'O x 2/.'O area is essentially flat (Heap et al. 1992). Either massive stars are more likely to be formed in the high-density center of the cluster or mass segregation due to dynamical evolution has taken place within the 30 Dor cluster.
5. Mass distribution Given the estimate of the initial mass of each star, we can determine the surface mass density profile by summing the masses of the stars in rings 20 pixels (0/.'88) wide and dividing by the area of the ring. Figure 4 is a plot of the log of the surface mass density as a function of log radius. The line is the projection of a King model with a core radius of 0//96 (0.24 pc). Our core radius is a factor of 4 larger than that derived by Campbell et al. (1992) from their PC images. The large difference is most likely due to their fitting of a King model to the surface brightness profile, instead of the surface mass density profile. As we saw in §3 and §4, the luminosity function and the IMF are both different near the center of the cluster than in the rest of the cluster. Furthermore, the optical luminosity of an OB star does not increase linearly with mass, but instead goes as M 2 2 . As a consequence, the luminosity profile is more sharply peaked than the mass density profile. While 21% of the total B-band luminosity is included in the central l'.'O, only 10% of the total mass is. The question of what is the correct core radius becomes important when we consider the dynamical state of the cluster.
6. Dynamical state of the 30 Doradus cluster One of the most interesting findings of the study is that there is an increase in the percentage of very massive stars near the center of the cluster. The question of why
18
Malumuth k Heap: The IMF of the Center of 30 Doradus
this is so is like the nature vs. nurture questions of human behavior. On the one hand (nature), there may be some density or position-dependent star-formation process that favors the formation of massive stars in the center of clusters or in very dense regions. On the other hand (nurture), the effects of dynamical evolution and mass segregation may cause the most massive stars to sink to the center of the cluster. Campbell et al. (1992) suggest that the relaxation time in the core of 30 Dor is given by the following equation (Spitzer & Hart 1971). = 8.7 x 1 0 6 - ^ _ (6 .2) G mp\nk mlnA Where a — ic\/4irGp0/9, rc is the core radius, m is the mass of a typical star, and p0 is the central density. For the Campbell et al. values of rc = 0.06 pc, p0 = 5 x 104 M© pc~ 3 , m = 10 M 0 , and lnA = 1.4, they get ( r = 3 x 104 yr. The cluster would be about 100 relaxation times old. However, if we use our values for the core radius, rc = 0.24 pc and central density, p0 = 9.5 x 1O3M0 pc~ 3 , we get a relaxation time of tr = 2.6 x 106 yr. The cluster is only one relaxation time old. Since core collapse happens at 12-19 median or half mass relaxation times (Binney & Tremaine 1987), Campbell et al. conclude that the cluster has undergone core collapse while we conclude that it has not. The relaxation time is small enough, however, that dynamical relaxation may be important. tr = 0.34
a
2
REFERENCES J. & TREMAINE, S. 1987 Galactic Dynamics. Princeton University Press, Princeton. CAMPBELL, B. et al. 1992 Astron. J. 104, 1721. DE MARCHI, G., NOTA, A., LEITHERER, C , RAGAZZONI, R. & BARBIERI, C. 1993 Astrophys. J. Accepted. HEAP, S. R., EBBETS, D. & MALUMUTH, E. M. 1992 In Science with the Hubble Space Telescope (ed. P. Benvenuti & E. Schreier), p. 347. European Southern Observatory, Garching. KURUCZ, R. 1991 Private communication. LANDSMAN, W. 1991 Private communication. MAEDER, A. 1990 Astron. Astrophys. Suppl. 84 139. MALUMUTH, E. M., NEILL, J. D., LINDLER D. J., & HEAP, S. R. 1991 In The First Year of HST Observations (ed. A. L. Kinney & J. C. Blades), p. 249. Space Telescope Science Institute, Baltimore. PANAGIA, N., GILMOZZI, R., MACCHETTO, F., ADORF, H.-M., & KIRSHNER, R. P. 1991 Astiophys. J. Lett. 380, L23. PARKER, J. W. & GARMANY C. D. 1993 Astron. J. 106, 1471. SALPETER, E. E. 1955 Astrophys. J. 121, 161. SAVAGE, B. D., FITZPATRICK, E. L., CASSDMELLI, J. P. & EBBETS, D. C. 1983 Astrophys. J. 273, 597. SPITZER, L. & HART, M. H. 1971 Astrophys. J. 164, 399. WALBORN, N. 1991 . In Massive Stars in Starbursts (ed. C. Leitherer, N. R. Walborn, T. M. Heckman, & C. A. Norman), p. 145. Cambridge University Press, Cambridge. WEIGELT, G. et al. 1991 In The First Year of HST Observations (ed. A. L. Kinney & J. C. Blades), p. 208. Space Telescope Science Institute, Baltimore. BINNEY,
The Nature and Kinematics of the Emission Nebulae in the Cyg OBI Supershell By T. G. SITNIK AND V. V. PRAVDIKOVA Sternberg Astronomical Institute, Universitetsky Prospect 13, Moscow, 119899, Russia Detailed Ha line investigations of the gas kinematics in the supershell around the Cyg OBI association were carried out. The supershell contains nebulae and CO-cavities around WR and Oj stars which form a hierarchical system of mutually embedded gaseous dust shells. The nebulae around WR 134, 135, 141 and 142 and the SNR G73.9+0.9 are shown to be located at the far edge of the parent molecular cloud at Vjsr ~ 5 — 10 km s" 1 . We found high negative velocities up to 70 km s" 1 and [OIII]-^Q emission stratification typical for shocks. Both could be associated with stellar wind and SNe. The collective wind and ionizing radiation of the Cyg OBI stars (especially WR) and supernova explosions must play some role in forming the supershell. There are some reasons to suppose that the gas at the sound velocity Viar ~ 12 km s"1 isflowingdowmstream of the ionization front.
Six years ago Lozinskaya & Sitnik, 1988 discovered a hierarchical system of mutually embedded gaseous-dust shells in the Cygnus arm (73 < / < 78°, - 0 * 5< 6 < 3° ). In the sky plane this system consists of several small-size shells around WR and Of stars (NGC 6888 among them) inside the supershell around Cyg OBI association (Figure 1). The supershell (diameter about 100 pc) and inner shells of different sizes are seen as optical ring nebulae, radio-shells, CO-cavities (Lozinskaya & Sitnik, 1988) and IR supershells (Lozinskaya & Repin,1991; Saken et.al, 1992). Analyzing radio and optical data, Lozinskaya & Sitnik, 1988 supposed that the hierarchical system might represent an integral formation created by the collective action of the winds of the Cyg OBI stars. The supershell has been shown to be located at the far side and above the giant molecular cloud. However, at that time we had only a fragmented knowledge about the region near Cyg OBI and we considered the results as preliminary until they could be confirmed by a detailed investigation of the complex's kinematics. Since then we have been carrying out an interferometric Ha survey using the FabryPerot interferometer and an image intensifier with a focus reducing camera mounted on the 125-cm and 60-cm telescopes. The results of these observations can be summarized as follows: (1) the whole region is dynamically very active. The FWHM of the Ha line reaches 100-120 km s" 1 in some positions inside the supershell; the mean value is about 50-100 km s" 1 ; (2) the Ha line emitted in the bright nebulae in the northern, eastern and western parts of the supershell is comparatively narrow and shows little variation; the mean radial velocity of the ionized hydrogen corresponds to V;sr ~ 9-12 km s" 1 , FWHM ~ 50 km s" 1 ; (3) in the southern part of the shell which we have studied in detail the radial velocities of the ionized hydrogen change over a wide range, from —70 km s" 1 to 45 km s" 1 . We have reported several velocity components of the ionized hydrogen in the southern nebulae that we have investigated so far: in the shell around WR 134, SNR G73.9 + 0.9 and near WR 141 and 142 (Lozinskaya, Pravdikova k Sitnik, 1993a, 1993b, 1994). To understand the origin of the different motions we have to clear up the following 19
20
Sitnik & Pravdikova: Cyg OBI supershell
FIGURE
1. The supershell around the Cyg OBI association.
questions: which velocity component corresponds to an undisturbed parent material and which one characterizes the peculiar motions. Let us answer the question using the ring nebula WR 134 as an example (Figure 2). This shell is located in the south-western part of the supershell. The ionized hydrogen here has the following velocity components: 5-7 km s"1 in the bright filaments, -42 km s" 1 to the north from the WR 134 and 12-14 km s"1 almost everywhere inside the shell. If we assume circular Galactic rotation, the interstellar clouds in the Cyg arm towards WR 134 would have radial velocities in the range 0-10 km s" 1 . Indeed, HI and CO observations (Dame et.al, 1988) confirm that the undisturbed interstellar clouds in this direction have the typical velocities from 0 to 10 km s" 1 . Thus the component Vj,r ~ 5-7 km s" 1 which dominates in the brightest filaments corresponds to the radial velocity of the galactic rotation in the shell's location. The CO emission distribution (Dame et.al, 1988) also indicates that the ring nebula is located in a cavity observed at Vhr ~ 0-5 kms" 1 . The fact that the same velocities are found out in the filaments and CO-cavity allows us to suppose that the bright filaments are formed in the dense molecular clouds by the ionizing radiation of the WR 134. The negative velocity component Viar ~ -40 km s" 1 can be connected with the stel-
Sitnik k Pravdikova: Cyg OBI supershell
FIGURE
21
2. The small-size shell around WR 134.
lar wind of WR 134. This velocity component cannot be associated with the Perseus arm, since similar negative velocity was found out in the interstellar absorption lines in the spectrum of WR 134 (St-Louis k Smith, 1991; Nichols-Bohlin k Fesen, 1993). So it is reasonable to suggest that the CO-lagoon has been formed by ionizing radiation and by the stellar wind of WR 134, and that the wind of WR 134 is responsible for the negative velocity component. The following arguments seem to confirm this suggestion. High negative velocities are seen to the north of the star. This is just the region where according to CO measurements and our Ha observations one can expect a higher wind-blown bubble expansion velocities because of the lower ambient density. There is another manifestation of the wind's action. The monochromatic observations of the shell show an [OIII] filament located further from the star WR 134 than the Ha filament (Lozinskaya, Pravdikova k Sitnik, 1993a). Such a stratification of optical emission is typical for the cooling gas behind a shock front formed here by stellar wind. The origin of the third velocity component (Vjsr ~ 12-14 km s"1) is not completely clear yet. It may be connected with the downstream gasflowingaway from the ionization front formed by WR 134. Indeed this velocity is also observed in the absorption spectrum of WR 134 (St-Louis k Smith, 1991; Nichols-Bohlin k Fesen, 1993) indicating that the gas motion takes place in front of the star and is directed toward the star. Moreover this velocity is equal to the sound velocity typical for a downstream flow. In a similar way we have studied the kinematics in the HH-regions near WR 141 and WR 142 and that of the possible supernova remnant G73.9 + 0.9. Thus, the results we have got at present for the supershell and especially for its southern part can be summarized as follows: (1) The Cyg OBI supershell is observed as a shell-cavity at optical, radio and IR wavelengths. The southern part of the supershell is formed by individual objects of a different nature, namely wind-blown shells and supernova remnants. (2) In molecular clouds we discovered small-size cavities around WR stars. They are observed at the same velocities as the bright filaments at the periphery of HII regions. The
22
Sitnik & Pravdikova: Cyg OBI supershell
cavities and bright filaments are traces of the interaction between the ionizing radiation, stellar wind and molecular cloud. (3) In the southern part of the supershell we have found gas motions of a different nature - galactic rotation, peculiar motions at a high negative velocity and at sound velocity. (4) High negative velocity motions up to 70 km s" 1 can be associated with SNR and winds of WR stars. The combined wind and ionizing radiation of these stars and supernova explosions must play some role in forming the supershell. Some manifestations of the stellar wind have been found in the southern part of the Cyg OBI supershell. (5) Gas motions at a positive velocity Visr ~ 12-14 km s" 1 might be caused by the downstream gas flowing away from the "combined" ionization front (like the ring nebula around WR 134). This velocity component is observed in different parts of the supershell as well as in the absorption spectra of the WR and 0 stars of Cyg OBI and Cyg OB3 (St-Louis & Smith, 1991; Nichols-Bohlin & Fesen, 1993). Thus these positive velocities cannot be associated with the background gas or with the receding side of the supershell. But this problem needs further investigation. We consider these results to be a step towards understanding the global and local interstellar structure and of the origin of the the supershell around Cyg OBI. Our investigations of the region are continuing.
REFERENCES DAME, T. M. ET AL. 1987 Astrophys. J. 322, 706.
T. LOZINSKAYA, T. LOZINSKAYA, T. LOZINSKAYA, T. LOZINSKAYA, T.
A. &c SITNIK, T. G. 1988 Soviet Astron. Lett. 14, 100. A. & REPIN, S. V. 1991 Astron. Zh. 67, 1152. A., PRAVDIKOVA, V. V. & SITNIK, T. G. 1993a Astron. Zh. 70, 24. A., PRAVDIKOVA, V. V. & SITNIK, T. G. 1993b Astron. Zh. 70, 469. A., PRAVDIKOVA, V. V. &; SITNIK, T. G. 1994 Astron. Zh. In press. NICHOLS-BOHLIN, J. & FESEN, R. A. 1993 Astron. J. 105, 672.
LOZINSKAYA,
SAKEN, J. M.,
SHULL, J. M.,
GARMANY C. D., NICHOLS-BOHLIN, J.,
Astrophys. J. 397, 537. ST-LOUIS, N. & SMITH, L. J. 1991 Astron. Astrophys. 252, 781.
FESEN, R. A.
1992
Asymmetry in the Vertical Distribution of Giant Molecular Clouds in the Carina Arm ByEMILIO J. ALFARO1'2, JESUS C ABRER A-C ANO 13 AND ANTONIO J. DELGADO1 1
Instituto de Astrofisica de Andaluda, (CSIC), P.O. Box 3004, Granada 18080, Spain 2
Astronomy Department, Boston University, MA02215, USA
3
Universidad de Sevilla, P.O. Box 1045, Sevilla 41080, Spain
The vertical distribution of molecular complexes along the Carina-Sagittarius arm has been studied on the basis of the giant molecular cloud (GMC) data compiled by Myers et al. (1986). The analysis indicates that the CO complexes are preferentially located below the formal Galactic plane. A separation of the sample into two groups: (a) GMCs associated with HII regions, and (b) GMCs without associated HII regions, establishes that group (a) shows, in average, larger z departure and mass than group (b). This result seems to suggest that the star formation activity in this major arm displays a vertical asymmetry which opens up interesting questions about the triggering mechanisms of star formation in spiral arms. The density and location of young stars in major spiral arms and the relation to their parent molecular clouds are important to the understanding of how molecular clouds evolve and form stars in our Galaxy. In previous work (Alfaro et al. 1992, 1993) we analyzed the vertical structure of young open clusters (YOCs) along the optical segment of the Carina-Sagittarius arm, and its connection with the density of YOCs as representative of star formation activity. The main conclusions of that work can be summarized as follows: 1. A clear correlation between YOC density and ^-departure from the formal Galactic plane is found when this density distribution is compared with the vertical structure. The cores of both supercomplexes are closely coincident with the two minima of the vertical profile, and the regions of lowest star-forming tracers appear associated with the relative maximum of z. 2. The z distribution of molecular complexes, aligned along 40 kpc in the arm, shows that the CO complexes are preferentially located below the Galactic plane. This distribution could be biased by: a) the warp of the Galactic disk, and b) the possible out-of-plane location of the Sun. In order to by-pass any possible bias produced by the warp phenomenon, we limited our analysis to the first Galactic quadrant (Dame et al. 1986) where one would expect the warp to be negligible or operate in the opposite direction. Nonetheless, the z distribution for this quadrant also shows the downward trend. 3. The asymmetry would only be apparent if the Sun turns out to be located above the actual midplane. The effect of this upon the result can however be avoided by comparing the z-location of GMCs with respect to each other and not with respect to an absolute reference plane. With this aim we separated the sample into two groups: one containing clouds with HII regions and another are formed by the GMCs with no HII regions. Our analysis led to the result that both groups display a different vertical behaviour, where the probability of finding a difference of means larger than, or equal to, the experimental value is less than 10%, assuming equal means as the null hypothesis. We try to assess it by performing the same analysis for the sample given by Myers et al. (1986). The catalogue of GMCs in the first Galactic quadrant listed by Myers et al. contains 56 objects, out of which 33 CO clouds present association with HII regions and 23 do 23
Alfaro et al.: Vertical Distribution of GMCs in the Carina Arm
24
lolaculor Clouds wlUi araoalatad HH region*
/
-T—v ! \ \
Itoleaular Claud* -without associated HH r*gl<ma
1 1
\
1. Histograms for GMC samples with and without associated HII regions. Histogram in mass (top), in z (middle), and in a linear combination of both variables (bottom).
FIGURE
not. The histogram in z for the two groups (Figure 1, middle), shows that the difference between both distributions, although present, is not so striking as displayed by the data of Dame et al. (1986). The quantitative two-sample analysis led to the result that the null hypothesis cannot be rejected, and that both samples should be considered as representative of the same parent population. On the other hand, a similar analysis for the mass led to the already known result that the two groups are probably different (Myers et al. 1986) in the sense that the more massive clouds are also those more frequently associated with HII regions of high radio continuum emission (Figure 1, top). We try however to go further in discerning the ability of the z variable to discriminate between both groups. With this aim we will test whether some combination of z and mass variables would improve the separation between groups, beyond those obtained with each single variable. The result of this analysis is shown in Figure 1 (bottom) where a linear z-log(mass) combination (Factor 2) represents the new variable. Visual comparison of this plot with the others shows that the combined effect produces a better separation than individual variables. This furthermore indicates that the z variable contains discriminating information, enhancing the one provided by the mass. REFERENCES ALFARO, E. J., CABRERA-CANO, J. & DELGADO, A. J. 1992 Astrophys. J. 399, 576. ALFARO, E. J., CABRERA-CANO, J. & DELGADO, A. J. 1993 In Star Formation, Galaxies and the Interstellar Medium (ed. J. Franco, F. Ferrini &; G. Tenorio-Tagle), p. 364. Cambridge Univ. Press. DAME, T. M., ELMEGREEN, B. G., COHEN, R. S. & THADDEUS, P. 1986 Astrophys. J. 305, 892. MYERS, P. C , DAME, T. M., THADDEUS, P., COHEN, R. S., SILVERBERG, R. F., DWEK, E. & HAUSER, M. G. 1986 Astrophys. J. 301, 398.
Supersonic Motions in Giant HII Regions ByCASIANA MUNOZ-TUNON 1
Instituto de Astrofisica de Canarias, E-38200 La Laguna, Tenerife, Spain
There are several questions in the field of the study of giant extragalactic HII regions (GEHRs) that remain controversial. The origin of the observed supersonic motions and the validity of the relationship between size and velocity dispersion obtained on the base of single-aperture observations are still debated. From a purely observational point of view, the spatial extent of kinematical features within the nebula, the sense in assuming Gaussian profiles as representative of the variety of emision lines found in GEHR, and finally how lawful the assignment of a single value (IT, .ft) as a defining kinematical parameter of a GEHR are also questioned. The energy input required to provide the supersonic motions observed in GEHRs has been the subject of several scenarios proposed in the literature and some of the points in favour and/or against them are mentioned in this contribution. Bidimensional spectroscopy with good spatial and spectral resolution, sampling of a particular emision line over the whole emitting area, is shown to be the most suitable observational technique for understanding the global kinematics of GEHRs. In NGC 604 for instance, features like loops and filaments although clearly seen in deep exposures and obviously resulting from stellar winds and SN explosions, do not dominate the bulk of the emission that comes from smaller but much brighter areas. Thus the total flux from split or non-Gaussian profiles that arise from shells and/or filaments is much lower than that obtained from bright knots where line profiles are well fitted by Gaussians. An evaluation and comparison of the emission according to line quality for each region is absolutely necessary in order to know the degree of disruption suffered by the remaining cloud under the action of stellar winds from massive stars. Here a new procedure is proposed for defining the size of each region based on its kinematics. The attained values (
1. Motivation: the state of the art Giant Extragalactic HII Regions have typical sizes ranging from 100 pc to more than 1 kpc. The large amount of ionized gas (MHII « 104 to 106 MQ) is embedded within an atomic and molecular cloud about ten times more massive. The high H<» luminosity (1039 erg s" 1 ) is atributible to several hundred massive (0 and B) stars. There is a dependence of the integrated luminosity on the morphological type of the host galaxy; GEHRs in early-type spirals and low-luminosity irregulars are 10 to 100 times less luminous than the above quoted value, whereas luminous Sc and Irregulars show regions 10 to 100 times brighter. Some of the physical properties of GEHR are reviewed by Shields (1990). The shapes and sizes of these large star-forming regions are diverse, and their internal structure is complex. Studies giving values for each region as a whole provide temperatures ranging between 5000 K to 20000 K. The mean electronic density (n e ) goes from 50 to a few hundred cm ~ 3 with typical r.m.s. values between 1 and 10 cm" 3 . The filling factor (nl/{nrm,)2), an observational parameter characterizing the degree of dumpiness of the region, is then / « 10~2. This confirms what on the other hand can be inferred from spatially resolved pictures: GEHRs present a very rich structure, with holes, tunnels, filaments and loops where a small fraction of the nebular volume seems to be ocupied by the high-density ionized gas. Following the advances in astronomical instrumentation together with the increase in 25
26
Munoz-Tunon: Supersonic Motions in Giant HII Regions
telescope diameters, the number of detailed studies of the structure of individual regions is growing. It has been found that there are significant changes in the physical parameters, i.e. temperature, or electronic density, depending on position on the nebulae. 1.1. Supersonic speed (a) Using photographic image plates, Kennicutt (1984) classified giant regions into different Morphological Categories, such as: Giant HII regions (C), Diffuse giant HII regions (D), Multi-core complexes (M) and Shell or ring-like regions (S). Long before, Sandage and Tammann (1974a) following a parallelism with other astronomical objects, gave a practical (observational) definition for these regions. GEHRs were posed as being systems with a Core/Halo structure based on which their core radius r c o r e was defined. Using single-aperture spectroscopy in 1970 Smith k Weedman reported the supersonic behaviour of the acquired velocity dispersion in some of these regions. This feature later on turned into a characteristic common to all of them and in fact should be part of the definition of GEHRs. The classification as supersonic, based on the sound velocity (c,) in a given medium, depends on the system temperature and for the above-mentioned range of temperatures implies velocities which lie between 9.0 km s" 1 and 17 km s" 1 . This physical characteristic of GEHRs left open an additional problem of how to maintain the observed supersonic motions. Several scenarios have been proposed for keeping the ionized gas in GEHRs moving at supersonic speeds. All of them basically use two physical resources, either the gravitational field or hydrodinamical flows. In the first case, the gravitational energy resulting from the potential well of the entire system (the gas-stars ensemble) would make gas clouds move at a supersonic velocity. The second mechanism, hydrodynamical flows, has been employed in two approaches. One of them uses the density discontinuities in the ISM that would favour the champagne phase of the HII region producing high velocity (supersonic) streams. The other one is based on the mechanical energy transferred to the medium by massive stars. The presence of large numbers of massive stars in GEHRs makes it plausible that their winds and SN explosions play a major role in the kinematics of star-forming regions. 1.2. GEHRs as distance indicators Using single-aperture results on the velocity width of the integrated spectrum and a photometric definition for the radius, Melnick (1977) found a correlation between line width and mean core-halo diameter for a group of giant HII regions. Using a larger sample and a more systematic data base, Terlevich k Melnick (1981) reported a a — Ardre empirical law followed by GEHRs. The similarity of this correlation (hereafter referred to as TSiM correlation) with that found in elliptical galaxies, bulges of spirals and globular clusters led them to suggest that giant HII regions are self-gravitating systems. In the same paper by it is also suggested that GEHRs satisfy the Faber k Jackson law (L oc a4) valid for the same virialized systems. The above results presented a new task to be sorted out within the field of star formation. We should understand not only the presence of supersonic motions lasting the whole life-time of the nebulae but also its physical relation between the observed velocity dispersion, luminosity and sizes of emitting objects. The molecular hydrogen in star-forming areas also follows a a — f{R) relationship similar to the T&M correlation (Sanders et a\. 1985). In Figure 1 the two observational laws are plotted together. Until very recently the only explanation proposed in the literature for the T&M correlation was given by the same authors. Following their argument, gravity would have to play a mayor role in big star-forming regions if they follow the same observational law
Munoz-Tunon: Supersonic Motions in Giant HII Regions • A a
Elbpticol o a l a i m GlotxJor clusters Giant HD regions Linear tit to Es Linear tit toEs ondGCs Linear tit toEs andGCs using core radii Linear til to HD regions
27
6166.
o
s
1
log velocity dispersion a (km/s)
FIGURE 1. Correlation between radius and velocity dispersion for globular clusters, elliptical galaxies and GEHRs (from Terlevich & Melnick 1981). Overlapped a similar plot for Galactic molecular clouds; from Sanders et al. (1985).
as globular clusters, bulges and elliptical galaxies. This model has often been referred to as the Gravitational Model (GM). Under the framework of the GM, gas clouds would move under the potential well of the gas-star system . Upon photoionization, the gas velocity dispersion would reflect the velocity distribution of a virialized system, defined by its total mass. The precise modelling of the GM should, however, run into problems. Clouds moving in such a way would frequently collide dissipating their kinetic energy on a relatively small time scale. Thus, unless clouds have an almost negligible cross section, the system would end up in a large cold cloud within the stellar cluster. Winds from massive stars have often been employed in the literature to explain the kinematics of GEHRs. Dyson (1979) developed the Coeval Winds Model (CWM in what follows) in which the kinetic energy observed in GEHRs is provided by a coeval burst of massive stars. A coeval burst of star formation is required to ionize the cloud. The large shell resulting from the combined action of these massive stars would produce an almost Gaussian integrated velocity profile when observed with coarse spatial resolution. The global gas velocity dispersion (cr) would be supersonic. There is an easy way to probe the CWM making use of spatially resolved observational techniques. The line profiles
28
Munoz-Tunon: Supersonic Motions in Giant Ell Regions
obtained when crossing the nebula should vary showing the maximum line-splitting at those positions near the centre of the bubble. This behaviour has been found for instance in Hubble III (Clayton 1987) and other ring-like nebulae but it is not observed in those GEHRs showing other (not ringed) morphologies. On the basis of the CWM, other scenarios based on stellar wind energy have been proposed in the literature, making only massive stars present in GEHRs reponsible for providing the kinetic energy of the gaseous system (see, for example, Kennicutt & Chu, this volume). In any case, the T&M correlation is not addressed by the wind models. A new idea has recently been proposed in the literature by Tenorio-Tagle et al. (1993. Stellar winds from low-mass stars moving under the potential well of the system (cloudstellar cluster) are the ingredients of the Cometary Stiring Model (CSM). For a given scenario of cluster cluster formation implying a critical temperature for star formation, the collection of low-mass stars with an increased cross-section given by their winds, and moving within the gravitational field of the system, would stir the remaining cloud. The cometary shocks, generated by the stellar wind sources ramming through the left-over cloud produce the cloud agitation. When at a subsequent stage massive stars form, the ionized cloud will show the supersonic sigma acquired from the cometary passage of the wind-driven sources. A detailed description of the model as well as the temporal evolution of the cloud is given by Tenorio-Tagle in this volume. The CSM accounts for the empirical correlation both between scale and velocity dispersion, cr, and between luminosity and a found in star-forming regions, molecular clouds and relaxed spheroidal stellar systems.
2. Observational set-up: a new approach The observational approach followed to obtain the above-mentioned results has been made with available techniques. Imaging (photometry) was used to get the core radii from surface brightness distributions. The dimensions (R) of GEHRs have then been obtained with photometric observations, with the assumption that it is possible to define a core/halo structure. Kinematics, mainly velocity dispersion (
Muiioz-Tunon: Supersonic Motions in Giant Ell Regions
29
depends on the chosen area and/or position in the nebulae. In sum, the issue is whether or not it is correct to associate a single (R, a) to a given GEHR. All this new information makes it necessary to define a new observational strategy and data analysis to match if we want to undertake, from the observational point of view, the task of defining the size and line width of GEHRs. Here we propose an alternative approach; spectroscopic 2-D mapping covering the whole region with good spatial resolution is a relevant present-day observational approach. Sampling of one emission line is enough for our purpose, provided it is measured with sufficiently good spectral resolution. The mapping of a region at seeing-limited spatial resolution in one emission line (e.g. H a ) with good spectral resolution should allow us: 1) to know the spatial trend of non-Gaussian asymmetric or split profiles, and 2) to estimate what the percentage of the total flux coming from different zones is, defined according to the shape of the emission lines. The use of the TAURUS2 Fabry-Perot imaging spectrograph (Taylor k Atherton 1980) at the 4.2-m William Hershel Telescope at the Observatorio del Roque de los Muchachos is most suitable for this sort of study. For a detailed description of this instrument see the Observers' Guide (1984) for the Isaac Newton Group of telescopes. The output from this set-up, after phase correction and calibration (Lewis Ic Unger 1991), is a 3-D data cube with x and y being the spatial plane and z the wavelength (etalon step) sampling. In Figure 2 a sketch of an output file is shown. Typical data-cube dimensions are (256, 256, 100) and the chosen etalon free spectral range is about 15 A, enough to sample H a (A = 6562.8 A). Expossure times of about one hour per run are required. The sum of all z planes would produce an image, similar to a narrow-band H a image, called a collapsed image. For a given pixel value (x, y) moving in z provides a spectrum. In summary, each TAURUS run contains typically a total of about 65536 spectra from an area of about 1.1 arcmin2. 2.1. Data Analysis: MATADOR The main problem with TAURUS data comes precisely from its high capability. It provides such an amount of data that often makes you want to run away before any useful information is in sight. After fleeing for quite a while we decided that the available software was not powerful enough to match the potential of the instrument, and therefore we decided to design software that could handle the problem. The outcome was named MATADORf, which was based on the IDL environment and was planned to resolve all present needs while keeping the option open of further extension in the future. In what follows all the data analysis referred to is done with MATADOR . As a first step, and in order to increase the S/N ratio, a global 3-D smoothing is performed. Gaussian functions are taken with widths chosen on the basis of seeing (in x and y) and calibration lamp-profile values (in z). It is important to remark that the line width, i.e. the a of a Gaussian profile, is defined as:
7 = /oexp(-|i).
(2.1)
In the literature, to characterize an emission line, one often finds expressions in terms of the full width at half maximum (fwhm) or the e-folding width. In order to avoid confusion we shall refer to line-width meaning a where, t MATADOR was developed by Vladimir Gavryusev and Casiana Munoz-Tuiion at the Instituto de Astrofisica de Canarias within the framework of the Grupo de Estudios de Forinacion Estelar (GEFE)
30
Munoz-Tunon: Supersonic Motions in Giant Ell Regions
X f (x, y, z ) = f ' ( x , y ) = "Collapsed Image"
256 plx
FIGURE
2. Sketch of phase-corrected TAURUS2 output file.
fwhm = 2>/2 In 2a - 2.35o-
(2.2)
e - folding = y/2
(2.3)
and Results presented here were taken with an IPCS-CCD, with typical seeing values of 0.7 arcsec and a spatial scale, oversampling the seeing, of 0.26 arcsec pixel" 1 . Spectral scanning AA/step = 0.16 A, equivalent to 7.3 km s" 1 in H a . Instrumental width, taken from a 2-D mapping of the calibration lamp, was ffinst = 1.75 ± 0.15 step. The width value, as directly measured on the observed line, will be called a. The result, after substracting thermal and instrumental widths will then be due to mass motions and is accordingly called
" 2 ="L» + ^ + "L,-
(2-4)
Once a data cube is smoothed two kinds of analysis can be undertaken using MATADOR, either global or detailed. Within the Global Analysis Procedure it is possible to do automatic 2-D fitting with a single Gaussian . A threshold, mask or criterion to set a limiting value to the pixels to be fitted can be defined. For instance, a signal threshold can be used as the starting point in order to avoid the fitting to noisy regions on the image. The output from the global fitting procedure are eight 256x256 maps showing the Gaussian fitted parameters: (1) peak intensity, (2) central wavelength, (3) sigma, (4) background, and (5-8) four more maps with the errors corresponding to fitted parameters. The Detailed Analysis Procedure allows one to choose the spectra to be studied by defining the position on the image and the beam size. It allows for the definition of the analytic function and number of them to be used if one wishes to fit the emission profile. These allow us to classify emission flux according to line quality and to redefine the size of the emitting region with a kinematical criterion.
Munoz-Tunon: Supersonic Motions in Giant Ell Regions
31
3. Some results for N G C 5461 and N G C 5471 Here we present some results on the two brightest GEHRs in M101, at a distance of 7.2 Mpc. The collapsed images in H a of NGC 5461 and NGC 5471 are presented in Figures 3 and 4. In Figure 3(b), for comparison with the TAURUS data, the narrow-band H a image from GEFE f collaboration is also shown. Nodes labelled A, B, C, D and E, as given by Skillman (1984), are marked on the image. The pixel size of TAURUS "images is 0.26 arcsec, equivalent to 9.05 pc linear scale in the galaxy, slightly different from that corresponding to the GEFE image. The comparison of both images in Figure 3 shows the global consistency in morphology as measured by two intrinsically different observational techniques. Despite the good quality of the GEFE image, the higher spatial resolution and seeing values of the TAURUS data allow us better to resolve some features in the emission structure. The collapsed image in H o of NGC 5471 is shown in Figure 4. On the image, knots A, B, C and D as, defined by Skillman (1984), are marked. This region, the brightest in M101, looks like an almost circular halo with several bright knots. For this reason it is considered a multi-core GEHR. The spatial resolution and seeing conditions of the observations presented here allow the identification of two knots, A and A', in what previously was classified as a single structure. 3.1. The kinematics of NGC 5461 From the sigma map of NGC 5461 (not presented here) the main results to be pointed out are: • Knots E and D are not connected with the main region. They present a much lower velocity dispersion within the range of normal (not supersonic) HII regions and are possibly physically separated from the rest of the region. This is also apparent in our Figure 3, although it is not until the velocity dispersion map is obtained that one can definitely establish that: • Knot C is also kinematically detached from the rest of the region. • By a kinematical criterion, NGC 5461 is not a region of lOOOpc size. The supersonic sigma values restrict NGC 5461 to knots A and B only. This reduces the size of this GEHR to about 500 pc in diameter. • "Well behaved Gaussian emission-line profiles extend over a diameter of about 500 pc and it is possible to assign a velocity dispersion value of amm s» 25 km s" 1 to the emission coming from this area. • On the edges of the 500 pc region there are poorly defined profiles. The intensity is much lower than that of the bulk of the region and asymmetries, line splitting and peculiar shapes are found. No detailed analyses have been performed due to S/N limitations. With our experimental approach it is then possible to assign a (D, a) value to NGC 5461. The size D — 500 pc was defined following the kinematics of the region and is smaller than previously reported, as measured using imaging techniques. Areas without welldefined line profiles are found bordering the region, but the bulk of the emission (in terms of intensity and covering area) comes from an area of 500 pc in diameter and can be represented by well-behaved Gaussian profiles with a
Munoz-Tuiion: Supersonic Motions in Giant HII Regions
32
- 1
200
250
500
450
400
350
300
150
200
250
300
350
FIGURE 3. (a) Collapsed image of NGC 5461 H a TAURUS data-cube. lpixel=0.26arcsec=9.05pc. (b) Narrow-band H a image of NGC 5461 from the GEFE collaboration.
Munoz-Tuiion: Supersonic Motions in Giant HII Regions
33
• All emission profiles from the sampled area are supersonic. • A large fraction (« 80%) of the emitting area presents a roughly constant emissionline width with Gaussian profiles characterized by a amm w 20 km s" 1 . • Split line profiles have also been detected in a small area on the border of the nebula. Their normalized peak intensity is « 1/20 of that of the bright regions. • There is a fraction of the emitting area (w 20%) with broader profiles. This includes the above mentioned split lines. A further study of the broader profiles over NGC 5471, results from the global analysis fitting procedure. The emission can be classified into two groups: 1) Broad unsplit lines. These arise from two small areas close to, but slightly displaced from, the two brightest knots (A and B). The broad profiles in knot A seem consistent with the molecular core disruption as proposed by Rodriguez-Gaspar & Tenorio-Tagle (this volume). Only knot B shows clear evidence of multiple components in the line profile. It is clearly non-Gaussian and seems consistent with the SN detected in the same area. 2) Lines not well represented by single Gaussians, coming from low surface brightness zones and identified with the split profiles reported above. Several limitations exist also in 2-D Fabry-Perot spectroscopy. One is the small spectral range which may hide the presence of low-intensity very broad components. In NGC 5471 (Castaneda et al. 1990; Mas-Hesse et al. 1994) have found the very broad (1000 km s"1) low (less than 20% of the peak) intensity component in node C, that we fail to detect in our present data. To sum up, it could be said that the a value measured in NGC 5471 using single aperture includes areas (20% of the total) with wider profiles. These localized broader profiles may be explained by winds and/or champagne flow mechanisms. Most of the emitting area however presents a (rmm(region) ss 20 km s" 1 and covers a radius of > 200 — 300 pc. An explanation other than winds from massive stars is necessary to account for the width of the emission measured in most nebular areas. The features from massive star effects, although detected, are only located on a small fraction of the nebula. 4. The case of NGC 604 NGC 604 in M33 at 720 kpc, ten times closer than M101 has been used extensively as a laboratory to study the ISM. The richness of its morphology, which can be resolved with our spatial sampling, makes this GEHR most suitable for the sort of kinematic study proposed here. Generally speaking, it has a core/halo structure. The halo as seen in H a can be described as a low-intensity background against which various loops and filaments can clearly be detected. The core is far brighter than the structures in the halo and can be resolved into two bright knots about 50 pc apart. Detailed analysis of the spectra arising from several areas on the image provides an important result; emission profiles from the core and the halo are different in shape. Line splitting and poorly defined Gaussian profiles are found in the halo, whereas the core mainly the two bright knots - has well defined Gaussian emission profiles. Figures 5 and 6 present, together with two collapsed images of NGC 604, spectra corresponding to several positions in the nebula. The four spectra shown in Figure 5 correspond to the four boxes drawn on the image. These emission profiles are typical of the halo region. As remarked above, line profiles are far from being Gaussian, displaying line-splitting and strong asymmetries. However, the two spectra in Figure 6, corresponding to the two emission peak maxima, are very similar and absolutely Gaussian in shape.
34
Munoz-Tunon: Supersonic Motions in Giant HII Regions 120 140
100
120
100
80 80
60
40
40
20 20
20
40
60
80
100
120
140
FIGURE 4. Collapsed image of TAURUS data-cube in HQ for NGC 5471. 1 pixel = 0.26 arcsec = 9.05 pc.
In both figures numbers in the grey level intensity scale represent the total number of counts (in detector units) measured within the FSR. In the spectra, "array index" means the z plane in etalon steps and "averaged intensity" is the normalized intensity for each given box. The most important result in NGC 604 is that typical line profiles arising from the core Ipeak(core) are more than 10 times brighter than halo emission lines / pea t(/ia/o); and core emission lines and halo emission lines differ in shape. Single-aperture spectra display an observed velocity dispersion
80%) presents asymmetries and line-splitting with a large variety of line profiles. • "Well-behaved" Gaussian profiles are associated with an area of D « 60 pc and o-mm = 18 km s" 1 .
Munoz-Tufion: Supersonic Motions in Giant HII Regions
35
Spactrum ot x-113 y-226 WnX- 3 bJnV- 5 MnZ» 1
250 '
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:'• - '''''••}>''''.*'•
'
150
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o I
.
.
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150
200
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FIGURE 5. Collapsed image of NGC 604 in HQ; 1 pixel = 0.26 arcsec = 0.86 pc. Four spectra, taken from several positions in the halo, are also plotted corresponding to boxes marked on the image. Intensity is given in arbitrary units and is normalized to the bin size.
5. Final Remarks The story started with the results obtained after single-aperture observations of GEHRs. The measured supersonic line widths led people to look for a physical mechanism to fuel the supersonic motions in giant star-forming regions. Time went by and we were able to see not only further out but also better. Things that seemed to be hidden before began to appear in our CCD images. Loops, filaments, rings and many other nice structures were disclosed from what before was a boring background. That sorted out part of the problem. We know, or at least we think we know, the necessary physics to be able to understand the interaction of massive stars with the interstellar medium. We are almost able to model the precise way in which these loops, filaments and rings are built. It is possible to recognize the features left by stellar winds and SN explosions in the ISM. Thus, the mechanisms that provide the energy for supersonic motions in quite a large fraction of the emitting area of some GEHRs is already known. Meanwhile, the T&M correlation was found. This relates the size of the brigthest zones in each GEHR and the measured velocity dispersion. Moreover, and despite the increase in spatial resolution, most GEHRs present areas - coincident with the brightest zones - where we fail to recognize the signature features of the massive stars. These areas, that could be named the kinematic cores, still present a supersonic velocity dispersion. Therefore, the TScM correlation and the origin of the observed supersonic motions in the kinematic cores of GEHRs have become the present observational challenge in the field.
36
Mufioz-Tunon: Supersonic Motions in Giant Ell Regions Spectrum ot x« 98 y» 155 X S binY5 binZ- 1 binX
250
150
200
150
100-5 Spectrum at x> 149 y-152 binX= 5 binY= tinZ= 1 2
100
'
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50 6
50
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50
100
150
200
250
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FIGURE 6. Collapsed image of NGC 604 in H a . Two spectra of the bright knots in the core are also plotted, corresponding to the boxes marked on the image. Intensity is given in arbitrary units and is normalized to the bin size.
Two-dimensional spectroscopic mapping covering the whole region with good spatial resolution is a very adequate observational approach to advancing our understanding of the kinematics of GEHRs. Such observations and the use of the data analysis procedure proposed here allow us to classify emission flux according to line quality and to re-define the size of the emitting region with a kinematical criterion. The kinematics of NGC 5461, NGC 5471 and NGC 604 have been studied and presented here as an example of our proposed approach. We are developing a similar study over a large sample of GEHRs and although clearly we do not yet have statistically significant results here we present a tentative list of conclusions: • After a spatially resolved study, it is permitted to assign a value of (
Muiioz-Tunon: Supersonic Motions in Giant Ell Regions
37
20
* %
/-fyi
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-
^'"jr'kTi
i
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r
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FIGURE
7. Narrow-band Ha image of NGC 588 in M33 from the GEFE collaboration.
the energetics of massive stars. But we are no longer observing the kinematical cores of the regions and thus the measured (
REFERENCES H. O., VILCHEZ, J. M. & COPETTI, M. V. F. 1990 Astrophys. J. 365, 164. CHU Y. H. & KENNICUTT, R. C. 1994 Astrophys. J. In press. CLAYTON, C. A. 1987 Mon. Not. R. astron. Soc. 226, 493. DYSON, J. E. 1979 Astron. Astrophys. 73, 132. CASTANEDA,
38
Munoz-Tunon: Supersonic Motions in Giant HII Regions
KENNICUTT, R. C. 1984 Astrophys. J. 287, 116.
R. C. & CHU, Y. H., 1994 This volume. J.& UNGER, S., 1991 William Hershel Telescope. TAURUS data and how to reduce it. MELNICK, J. 1977 Astrophys. J. 213, 15. KENNICUTT,
LEWIS,
MAS-HESSE, M.,
MUNOZ-TUNON, C ,
VILCHEZ, J. M., CASTANEDA, H. O. & CARTER,
D,,
1994 This volume. G., 1994, This volume. N., TENORIO-TAGLE, G., CASTANEDA, H. O. & MUNOZ-TUNON, C , 1994 This volume. SABALISK, N., TENORIO-TAGLE, G., CASTANEDA, H. O. & MUNOZ-TUNON, C , 1994 Astrophys. J. Lett. Submmitted. SANDERS, D. B., SCOVILLE, N. Z., SOLOMAON, P. M. 1985 Astrophys. J. 289, 373. SANDAGE, A. & TAMMANN, G. 1974 Astrophys. J. 190, 525. SANDAGE, A. & TAMMANN, G. 1974 Astrophys. J. 194, 223. SHIELDS, G. A., Ann. Rev. Astron. Astrophys. 28, 525. SKILLMAN, E. 1984 PhD Thesis. TAYLOR, K. & ATHERTON, P. D. 1980 Mon. Not. Roy. Astron. Soc. 191, 675. TENORIO-TAGLE, G. 1994 This volume. TENORIO-TAGLE, G., MUNOZ-TUNON, C. & Cox, D. P. 1993 Astrophys. J. 418, 767. TERLEVICH, R. & MELNICK, J. 1981 Mon. Not. R. astron. Soc. 195, 839. UNGER, S., BRINKS, E., LAING, R. A., TRITTON, K. P. & GRAY, P. M. 1988 Isaac Newton Group, La Palma. Observers' Guide. RODRIGUEZ-GASPAR, TENORIO-TAGLE,
SABALISK,
A Kinematic Study of NGC 604 By N. S. P. SABALISCK, G. TENORIO-T AGLE, H. O. CASTANEDA and C. MUNO Z-TUNON Instituto de Astrofisica de Can arias, E-38200 La Laguna, Tenerife, Spain We present a kinematical study of NGC 604 in the emission lines of Ha and [OIII] A5007 A, based on data obtained with TAURUS-2, a Fabry-Perot imaging spectrograph. The main result is that the global supersonic velocity value is dominated by emission coming from the two maximum brightness peaks covering an area of about 12 arcsec2. NGC 604, located in M33, is one of the most prominent giant extragalactic HII regions (GEHR). It was observed with TAURUS-2 at the 4.2-m William Herschel Telescope (La Palma), with the Image Photon Counting System (IPCS) as detector. Two data cubes, each with an integration time of 3600 sec, were obtained in the emission line of Ha and [OIII] A5007. The 125-/im etalon was used, which gave a free spectral range for Ha of 17.6 A (806.6 km s"1) and 10.0 A (596.1 km s - 1 ) for [OIII]. The velocity resolution (a of the calibration line) was 23.1 km s" 1 for Ha and 19.6 km s" 1 for [OIII]. In order to determine the origin of the supersonic mass motion in NGC 604, the kinematics of the region were analyzed by means of individual spectra, applying automatic fitting routines with single Gaussian fits. The measured velocity dispersions were corrected for instrumental and thermal broadening. The thermal velocity dispersion correction was calculated assuming Te = 10,000 K. The integrated spectrum of the whole region shows a global heliocentric radial velocity of —244 km s" 1 and a velocity dispersion (
Sabalisck et al.: A kinematic Study of NGC 604
40
NCC604 - Ha
:
•
I
n.
V
. .r 0
10
20
30
40
VELOCITY DISPERSION
SO
SO
0
(km/s)
i 10
NGC604 - [0llI]5007k
20
30
40
VELOCITY DISPERSION
SO
SO
(km/s)
1. Histogram of velocity dispersion versus the sum of the integrated flux for (a) Ha and (b) [OIII] A5007 A. The sum of the integrated flux is taken at 2 km s" 1 intervals of a. Velocity dispersion values are corrected for thermal and instrumental broadening. FIGURE
NGC 604 - Ha emission line (BO B )
100
NGC 604 - Ho emission line (15.2")
200 -
77
a 100 13.3 E a = J7 5 k n/i gtoDai a
pixel
160
200
km/s 100 pixel
200
12
FIGURE 2. (a) - Ha flux maps of NGC 604, used to divide the region into 49 squares of 60.8 arcsec2 each (marked with a thin dotted line). We obtained the flux and a of each part to calculate the global a in different flux ranges, identified with different kinds of lines. In (b) we have plotted the most intense area, divided into squares of 15.2 arcsec2. Only the zones with more than 70% of the maximum flux values are shown here. Their corresponding global a is 16.5 km s" 1 . The X-Y axis of these maps are in units of pixels (1 pixel = 0"26).
The result implies an intrinsic mechanism agitating the high-density gas left over from the violent star forming process. We gratefully acknowledge partial support from DGICYT Grant PB91-0531(GEFE) and the NATO Grant CRG920198 for collaborative research.
REFERENCES HIPPELEIN, H. & FRIED, J. W. 1984 A. & A. 141, 49. MELNICK, J. 1977 Ap. J. 213, 15. ROSA, M. & SOLF, J. 1984 A. k A. 130, 29.
UV Spectroscopy of Giant Extragalactic HII Regions: the Case of NGC 604f ByLOURDES SANZ FERNANDEZ DE CORDOBA Lab. Astroffsica Espacial y Fi'sica Fundamental, Instituto Nacional de Tecnica Aeroespacial, Apartado 50727, 28080-Madrid, Spain The UV spectrum of a giant extragalactic HII region is compared here with the synthetic spectrum generated by a fitting procedure which uses a combination of stellar spectra to match the observed spectrum.
1. Observations The ultraviolet observations of giant extragalactic HII regions provide very valuable information on the stellar population embedded in these complexes. The UV spectra of HII regions in the IUE short- wavelength range, SWP (1152-1950A), show a stellar continuum which rises steeply towards the shorter wavelengths. Superposed over this continuum, emission lines and P-Cygni profiles are often observed, as an indication of the presence of hot luminous early-type stars with mass loss. NGC 604, the brightest HII complex in the nearby galaxy M33, is, because of its size and distance, a good candidate for a detailed study of a giant extragalactic HII region in the UV based on WE data. NGC 604 has been observed with IUE on different occasions in low resolution (6 A mm"1) with the large aperture slit (10x20 arcsec2). A total of 16 SWP and 8 LWR, or LWP (1950-3200 A), has been obtained. The spectra correspond to observations of different zones within the region. The general characteristics of the IUE spectra of NGC 604 in the short wavelength range are similar to the ones mentioned above. In the near UV, the observed continuum is practically constant, no emission lines are detected, there is no 2200-A absorption feature, and only interstellar absorption lines are present. 2. Analysis of the data For the detailed qualitative study of the stellar population of this giant extragalactic HII region in the UV, a program for reproducing the UV observed spectrum with a certain combination of individual stellar spectra was developed. This fitting program is based on a least squares fit to a non-linear function. An iterative process minimizes x2 using a combined method of gradient search plus an analytical function. Input data are • the UV spectrum of NGC 604, corrected only for foreground galactic extinction, E(B - V) = 0.03, and • UV spectra of luminous stars of different spectral types (between O3-B5 with UV luminosity classes s and g, and WR stars, N and C types), selected from the IUE-Archive and corrected for distance and reddening. The output result is a composite spectrum, the optimized fit for the set of standard f Based on observations from the International Ultraviolet Explorer satellite. 41
42
Sanz Fernandez de Cordoba: UV Spectroscopy of Giant HII Regions
20.0
0.0
1400.0
1600.0
1B0O.O
Wavelength |A]
FIGURE 1. Comparison between an observed UV spectrum of NGC 604 (solid line) and the optimized composite spectrum obtained by the fitting process (dotted line). stars being used as input. The quality of the resultant fit is a function of the selected set of standard stars.
3. Results Figure 1 is a graphical comparison between a UV spectrum of NGC 604 and the composite spectrum obtained at the end of the fitting process. In this case, the synthetic spectrum results from a combination of O5d, 05s, O7d, O9d, WN7 and WC7 stars selected as input. As shown by this figure the observed spectra of NGC 604 can be satisfactorily fitted to the synthetic spectra, not only in the continuum but also in the observed stellar emission lines and P-Cygni profiles. With this systematic fitting procedure the IUE UV spectra of NGC 604 were analized and compared in detail in order to obtain the luminous and young stellar population content in the complex (Sanz Fernandez de Cordoba 1986). A more detailed presentation and discussion of the results obtained is to be published.
SANZ FERNANDEZ DE CORDOBA,
REFERENCES L. 1986, PhD Thesis, Univ. Granada, Spain.
Evolution of GEHRs: The Effects Caused by Champagne Flows By J. A. RODRIGUEZ-GASPAR AND G. TENORIO-TAGLE Instituto de Astrofi'sica de Can arias, E- 38200 La Laguna, Tenerife, Spain We analyzed the possible evolution of a high surface brightness GEHR into an extended diffuse region due to champagne-produced outflows. We calculate the typical time-scales and the final radius and density for a variety of initial parameters. In addition, a numerical calculation which allows visualization of the evolution is shown. The Giant extragalactic HII regions (GEHRs) in nearby spiral, irregular and blue compact galaxies have widely been studied during the last years due to their remarkable properties, such as their large luminosities, diameters and supersonic gas motions (see review by Shields 1990). Kennicutt (1984) analyzed the structural properties of GEHRs and classified them morphologically as "classical", "high-surface-brightness", "diffuse", "multiple core", "shell" and "supergiant". Here, we suggest a time sequence from the "high-surface-brightness" to the "diffuse" class by the effect of champagne flows. Let us assume a spherical cloud that includes an extended region with radius rc and constant density n c and a central core with a power-law density stratification, '-\niO(r/rior
for ri0 < r < ric
(UA
>
where rj c = r,o (»iio/>»c) , is the external radius of the core. Outside the cloud there is the ISM with a constant density noThe cloud is ionized by a powerful cluster of O5V stars located at the center and with a Lyman continuum photon luminosity of Ni(O5 V) = 5 • 1049 sec" 1 each one. Eventually, the whole cloud becomes ionized and the gas in the core and at edge of the cloud is set into motion because of the pressure gradients. The inner core is disrupted, producing a strong shock whose evolution is given by Franco, Tenorio-Tagle k. Bodenheimer (1990) as r,(<) = r
=
- u;)]
(0.2)
At the edge of the cloud the usual conditions for producing a champagne flow are present. The pressure discontinuity develops into a strong shock that moves into the ionized ISM and a rarefaction wave (RW) that moves inwards into the cloud. The shock strength or Mach number M, is given by Bedjin k Tenorio-Tagle (1981) as »c/»io = Mj exp \{Mj — 1)/M,]. The RW causes a continuous streaming of material out of the cloud and the high-density cloud expands into a more extended and "diffuse" region. The time necessary for this is set by the speed of the RW, tc — rc/c, during which the radius of the HII region grows to rcj = r c ( l + Ma), and the average density falls to n c no/nc + M, . = (0 3) 3
"•'
MT(I
+
MS) -
'
The time necessary to dissipate small condensations inside GEHRs is smaller than the 43
Rodriguez-Gaspar & Tenorio-Tagle: Evolution of GEHRs
44 CORE
Core Disruption
.\
i-
•I
Cloud expansion
.^ !
CLOUD
r
i
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/ '
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1 1
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8.0X103 RodH (pc)
Rodii (pc)
FIGURE 1. Hydrodynamic models and surfaces brightness for H(v in erg cm" sec ' sr ' . The initial parameters are: u,o = 106 cm" 3 , nc = 10J cm" 3 , no = 1 cm" 3 , r,o = 2 • 1016 cm , Tc = 3.86 • 1019 cm and NL = 1051 sec"1 (w 20 05 V stars). The evolutionary times are: 0 (initial condition which is represented by a solid line), 3.8 • 104 (dashed line), 1.8 • 105 (dotted line) and 6.3 • 105 (dash-dotted line) yrs.
GEHRs' evolution time. Thus, the "high-brightness" regions having small and denser condensations, unresolved observationally, could evolve into "diffuse" regions with a nearly uniform density. For example, for 30 Doradus, tc = 8.5 • 106 yr (we assume nc « 10 cm" 3 and rc w 100 pc; Kennicutt 1984) while the time neccesary to dissolve a dense core (n,o = 103 cm" 3 and TJO = 1 pc) is 3.7-105 yr, about a 30% smaller. However, this value is only a lower limit to the disruption time of the core. A hydrodynamical simulation using the procedure described by Tenorio-Tagle et ai. (1986) showed excellent agreement with previous estimates. Figure 1 shows the plots of density, velocity, temperature and Hcv brightness at selected evolutionary times. This work is supported by DG1CYT (grant PB91-0531GEFE) and NATO (grant CRG920198 for collaborative research).
REFERENCES BEDJIN, FRANCO,
P. J. &; TENORIO-TAGLE, G. 1981 Astron. Astrophys. 98, 85. J., TENORIO-TAGLE, G. & BODENHEIMER, P. 1990 Astrophys. J. 349, 126.
KENNICUTT, R. C. 1984 Astrophys. J. 287, 116. SHIELDS, G. A. 1990 Ann. Rev. Astron. Astrophys. 28, 525. TENORIO-TAGLE, G., BODENHEIMER, P., LIN, D. N. C. & NORIEGA-CRESPO, A. 1986
Not. R. Astron. Soc. 221, 635.
Mon.
Bursts of Star Formation in Central Disks of Galaxies By VLADIMIR SURDIN Sternberg Astronomical Institute, Universitetskii Prospect 13, 119899 Moscow, Russia Central gaseous disks around the nuclei of flat galaxies continually increase their mass due to spiralling giant molecular clouds (GMCs) under the action of dynamical friction. The radius of the disk depends on a tidal condition in the central parts of a galaxy equal to GMC tidal disruption distance. A central part of the disk can become molecular and be able to undergo a subsequent spontaneous burst of star formation when the mean surface density of the disk becomes larger than the critical UV-opacity column density. In the outer parts of the disk, formation of gaseous clumps and stellar aggregates can be self-consistent with a characteristic clump mass of 107 MQ and new-born stellar masses of 106 MQ in each clump. This scenario is a good approximation to the observable characteristics of central molecular disks of normal galaxies like ours. However, the interaction of galaxies must modify the maximum cloud and stellar aggregate mass up to ~ 108 M Q and lead to stimulated bursts of star formation.
1. Introduction Bursts of star formation are frequently observed in the nuclear regions of galaxies. Observations have revealed compact nuclei surrounded by an extended disk, or disk-like HII regions with radii of about 200-500 pc (Wilson et al. 1991). These regions are, as a rule, the brightest HII regions in such galaxies. One of the main reasons for such intense starbursts is the high star formation efficiency (SFE) in galactic nuclear regions (Planesas et al. 1989). Given the high SFE in the elementary cells of star formation there are dense clumps (n ~ 104 cm"3) in GMCs - it is not surprising if huge analogs of the clumps, the central molecular disks, are located in the central region of some galaxies and have masses of a few 108 MQ. We know examples of this kind of disk in our Galaxy (Bally et al. 1988), M82 (Sofue et al. 1992; Wild et al. 1992), and several more. External parts of the molecular disk are consistent due to the presence of cold neutral gas. In some cases these regions share a few cloudy complexes with active star formation, and mimic multi-nuclei galaxies. In this paper I will try to find answers to the following questions: - What physical mechanisms lead to the particular parameters and structure of the central molecular disk (CMD) and the central gaseous disk (CGD)? - Why does the process of star formation near galactic cores occur as a large-scale outburst? - Why is the SFE of the process so high? 2. Formation of central disks There are some ideas towards the exploration of star formation bursts in isolated galaxies such as, for example, NGC 253 (Rieke et al. 1980). It may have been an externally stimulated process through the accretion of intergalactic gas clouds fueling the nuclear region and probably triggering star formation bursts (Sofue &; Wakamatsu 1991). But in isolated galaxies one should consider the angular momentum problem, and the main cause of a starburst must be inner dynamical processes, such as gas viscosity in a differentially rotating disk (Lesch et al. 1990), dynamical friction of GMCs (Surdin 45
46
Surdin: Bursts of Star Formation
1980a, 1980b, 1981; Stark et al. 1991), and angular momentum loss in the potential of a central bar (Hernquist 1989). Of the three mechanisms dynamical friction is the only one that works without any requirements regarding the particular form of the galactic potential. This mechanism provides a natural explanation for the absence of HI and CO emission between R = 3 — 4 kpc and R = 1.5 kpc in the Galactic disk. When GMCs spiral into the galactic center under the action of dynamical friction, they are destroyed by the tidal gravitational field of the galaxy at some distance from the galactic center. The condition for tidal stability requires that the mean density of a gravitationally bound cloud (pc) exceeds the galactic mean density inside the cloud's orbit (pG) (Spitzer 1987): Pc>ApG,
(2.1)
where A is a constant depending on the galaxy mass distribution: A = 3 for a Keplerian rotation curve, and A = 2 for constant rotational velocity (Rastorguev k Surdin 1978). If equation (2.1) is used for the galactocentric distance at which a GMC is disrupted, one can obtain the radius of the central gaseous disk (RCGD) in the simple case of constant circular velocity Vo in the galactic disk: RCGD = 1 ( . . . , ° _ t ) ( i n " p
M C
_3 )
kpc,
(2.2)
V200 kms V \100 H2 cm A J where UGMC is the mean GMC density. For our Galaxy Vo — 200 km s" 1 and UGMC — (50 — 200) cm" 3 , therefore RCGD = (0.7— 1.4) kpc, in good agreement with observations of the HI circumnuclear disk of our Galaxy. Intensive spontaneous formation of molecular clumps inside the CGD can develop only if the mean surface density of the CGD is larger than the critical high-opacity column density required for molecule formation:
300 (^
J MQ pc" 2 ,
(2.3)
where Z is the metallicity of the gas. From (2.2) and (2.3) we can obtain the maximum mass of the CGD just before starburst:
" (f
10-
The local SFE value depends first of all on the mean density of a cloud. It is because the star formation process is interrupted by the energy output of massive stars. If the cloud's free-fall time, tjj ~ (Gpc)" 1 / 2 , is greater than the pre-main sequence evolution time of massive stars (tpms ~ 106yr), then the stars of a first generation would have enough time to destroy the cloud by their energy output. That is why GMCs as a whole with tff « 7 • 106 yr have SFEs as low as 0.1 - 1% (Surdin 1989). But the cores of GMCs with tjf « 4 • 105 yr have SFEs as high as 20 - 60%. A possible explanation of nuclear starbursts is the delay of clump formation under the tidal action of a galactic nucleus. If the condition of high SFE is taken into account: tff < tpms
(2.5)
and tidal condition (2.1), we can obtain the radius of the central molecular disk where a starburst can occur: RCMD = 0.4 (
^
\200 km s
_,)
kpc
/
(2.6)
Surdin: Bursts of Star Formation
47
and estimate the mass of the starburst disk:
(f)«0.
(2.7)
We may also evaluate the CMDs mean star formation rate (SFR): SFR = SFE- McMDT~\
(2.8)
and luminosity (L): L = io^.(^iA
T
\
MQ
(2.9)
)
where r is the duration of the starburst. If we assume that all massive stars are born at the same time, then r is the main-sequence lifetime of the stars: r m ,(30 MQ) = 5x 106 yr. With a value of SFE ~ 100% we obtain 4Q yr" 1 \zuu Km s
/
(2.10)
\ ^ /
and
L~W42( ^ T) f^ergs" 1 . \ 2 0 0 k m s - V \ZJ&
(2.11) '
v
These formulae give us a good approximation to the observed characteristics of CGDs, CMDs and starbursts in the circumnuclear regions of disk galaxies. 3. Self-induced formation of gaseous clumps and stellar aggregates In principle, the star formation process in the central part of a galaxy may be selfstimulated. It is well known that in molecular rings almost each GMC would create a young stellar aggregate which will evolve into an OB association and/or an open cluster. Threse newborn stars inject sufficient radiative and mechanical energy (stellar wind) to destroy their parent molecular cloud and create a rapidly expanding interstellar shell. As a result of the shell expansion a dense banana-shaped envelope forms in differentially rotating galactic disks (Tenorio-Tagle & Palous 1987). A large part of the interstellar matter swept up during the envelope evolution accumulates at the tips of the "banana remnant" and leads to the formation of GMC, which can begin a new cycle of star formation. In this picture, star formation and the formation of GMC are a self-connected process. It can take place not only in the molecular ring of a galaxy but in the central gaseous disk too. We can evaluate typical masses of clumps (C) and stellar aggregates (S) in a process of consequential formation. The evolution of supershells embedded in differentially rotating disks was analyzed by Tenorio-Tagle & Palous (1987) and Palous et al. (1990). For disks with flat rotation curves Palous et al. (1990) obtained a simple formula for the mass of a molecular cloud formed on the tip of banana-like envelope: M c = 6 x
lcm
where h.4 is the half-thickness of the disk, n0 is the density of interstellar matter, and E is the total input of energy. On the other hand, let us consider a gravitationally bound cloud of mass Me, and radius Re- Its binding energy is \Uc\ = GMQ/RCHOW many stars can be created in the cloud before it is destroyed by the stellar energy input? We can write the energy
48
Surdin: Bursts of Star Formation
input as Es=n(3E0Ms, (3.13) where EQ is the output energy of one massive star, \i is the fraction of this energy conserved in the hot bubble of gas, /? is the number of massive stars per unit mass of young stars, and Ms is the total star mass. Star formation will cease under the natural condition (Surdin 1989): \UC\ = ES. (3.14) From this equation, for EQ = 10 51 erg, fi = 0.1, and /? = 10~ 2 (Larson 1974), we have:
J^^OOHTcm-3)173'
<3-15>
where nc denotes the mean density of the clump. In CGD we can take nc ~ 103 — —104 cm"3, and finally we have:
Taking into account the equality fiE — Es, one can construct, from (3.12), (3.13), and (3.16), a system of two equations with two unknowns Ms and Me-
Ms
\
( Mc
\5/3 (
Mc
0.44 /
\ 0.54 /
F
\ 0.66
where agas is the total column density of interstellar matter perpendicular to the disk. The solution of the system is: Mc = 7 x 103 Mo (r^—)
( ,/ga'_2) 3.4
(J.lOj
Because the described proccess of cloud formation is stimulated, the value of Me is the maximum mass of a cloud in their population. Let us consider, for example, the central part of our Galaxy (R < 1 kpc). For this case agas ~ 150 MQ pc~ 2 and ha ~ 40 pc (Sanders et al. 1984). From (3.18) we can predict the maximum mass of a cloud: Me — 2 x 107 MQ. It is only a few times bigger then the mass of the Sgr B cloud. The solution (3.18) also predicts a rather high SFE (= Ms/Mc ct 15%) in the CGD of the Galaxy.
4. Discussion As it has been shown that the main characteristics of central gaseous disks of galaxies can be explained by the tidal instability of spiralling GMCs, the critical value of the UV-opacity column density, and the delay of clump formation in the disk under the tidal action of a galactic nucleus. This last mechanism provides the formation of clouds with a mean density of the order of 104 cm" 3 only. The free-fall time of such clouds is less than the pre-main sequence evolution time of massive stars. This may be the main reason for a high value of SFE during nuclear starbursts. We have also considered a self-consistent stimulated process of cloud and stellar aggregate formation. In a stable
Surdin: Bursts of Star Formation
49
regime the process produces low-mass clouds (Me ~ 107 MQ in our Galaxy) and not very abundant stellar aggregates (Ms ~ 106 MQ). Finally, I would like to mention the dependencd of Ms and SFE from the vertical scale of the disk: Ms oc hil
(4.19)
and SFEoc/i 16 . (4.20) The gravitational interaction of galaxies may lead to an increase in velocity dispersion and the vertical scale of gaseous disks by large factors; for example, in, numerical experiments of Elmegreen (1993), up to 2 — 3 times. It can also lead to an increase in the value of Ms of up to ~ 108 MQ and in SFE up to 80%. In other words, it can trigger a burst of active star formation. Participation at the Conference was made possible in part by Travel Grant No 1005/4 from the International Science Foundation.
REFERENCES A. A., WILSON, R. W. & HENKEL, C. 1988 Astrophys. J. 324, 223. B. 1993 This conference. HERNQUIST, L. 1989 Nature 340, 687. LARSON, R. B. 1974 Mon. Not. Roy. Astron. Soc. 169, 229. BALLY,
J.,
STARK,
ELMEGREEN,
LESCH, H., BIERMANN, P. L., CRUSIUS, A., REUTER, H. P., DAHLEM, M., BARTELDREES, A.
& WiELEBlNSKl, R. 1990 Mon. Not. Roy. Astron. Soc. 242, 194. J., FRANCO, J. & TENORIO-TAGLE, G. 1990 Astron. Astrophys. 227, 175. PLANESAS, P., GOMEZ-GONZALEZ, J. &; MARTIN-PINTADO, J. 1989 Astron. Astrophys. 216, 1. RASTORGUEV, A. S. & SURDIN, V. G. 1978 Astron. Tsirk. 1016, 3. PALOUS,
RIEKE, G. H., LEBOFSKI, M. J., THOMPSOM, R. I., Low,
F. J. & TOKUNAGA, A. T.
1980
Astrophys. J. 238, 24. SANDERS, D. B., SOLOMON, P. M. & SCOVILLE, N. Z. 1984 Astrophys. J. 276, 182. SOFUE, Y., REUTER, H.-P., KRAUSE, M., WIELEBINSKI, R. &: NAKAI, N. 1992 Astrophys. J. 395, 126. SOFUE, Y. & WAKAMATSU, K. 1991 Publ. Astron. Soc. Jap. 43, L57. SPITZER, L. 1987 Dynamical Evolution of Globular Clusters. Princeton Univ. Press. STARK, A. A., GERHARD, O. E., BINNEY, J. & BALLY, J. 1991 Mon. Not. Roy. Astron. Soc. 248, 14P. SURDIN, V. G. 1980a Astron. Tsirk. 1113, 3. SURDIN, V. G. 1980b Astron. Tsirk. 1126, 1. SURDIN, V. G. 1981 Astron. Tsirk. 1178, 6. SURDIN, V. G. 1989 Astron. Nadir. 310, 381. TENORIO-TAGLE, G. & PALOUS, J. 1987 Astron. Astrophys. 186, 287. WILD, W., HARRIS, A. I., ECKART, A., GENZEL, R., GRAF, U. U., JACKSON, J. M., RUSSELL, A. P. G. & STUTZKI, J. 1992 Astron. Astrophys. 265, 447. WILSON,
A. S.,
HELFER,
T. T.,
HANIFF,
G. A. &
WARD,
M. J. 1991 Astrophys. J. 381, 79.
Violent Star Formation ByGUILLERMO
TENORIO-TAGLE
Instituto de Astron'sica de Canarias, E-38200 La Laguna, Tenerife, Spain The quasi-static gravitational collapse of clouds should lead to fragmentation into small Jeansunstable cloudlets with a column density similar to that of the clouds as a whole. This criterion defines the mass and size of newly-formed fragments as the collapse proceeds. One could in principle expect that all collapsing matter should end up in a massive singularity at the centre of the collapsing configurations. However, the onset of star formation is shown here to stop the collapse. This should happen if the newly-formed stars produce winds while ramming through the left-over cloud and in this way cause the stirring required to stabilize the collapse. An inmediate consequence of this is the velocity dispersion generated in the star-forming region, which is supersonic in the case of massive clusters and detectable upon the appearance of massive stars in giant HII regions. The stirring cause by the supersonically moving wind-driven sources is also shown to cause a distinct cloud structure, or filling factor, in excellent agreement with recent observations of regions of violent star formation. The two effects, i.e. the acquired velocity dispersion and the filling factor caused in the parent cloud, allow us to differentiate between high- and low-mass clusters. In both cases, however, the disruptive energy from massive stars ends up erasing the clues stored in the gas during cluster formation.
1. Introduction There are several trivial but definitive conclusions regarding the history of spheroidal stellar systems that one can infer from the present agglomeration of stars in globular clusters and galactic bulges. In both cases for example, it is quite obvious that star formation stopped ages ago. It is also evident that star formation took place over a sufficiently long period of time, or that there must have been such strong and sudden bursts as to convert large fractions of the original cloud masses into stars, allowing the clusters to appear today as bound stellar systems. However, most important of all, is the fact that there must have been a stabilizing agent. Something stopped the collapse of the clouds as a whole and did not allow the formation of supermassive objects at the centre of the configurations. Instead, what seems the standard answer of Nature to such a situation is what actually happened: plain, simple fragmentation! This led to innumerable smaller and smaller self-gravitating cloudlets out of which eventually resulted the ensemble of stars that today seem to have relaxed their spatial and velocity distributions into a virialized state. If one could follow the history of stellar clusters far back in time, one would eventually find the short time interval over which their massive stars were evolving along the main sequence. The left-over gas would then have appeared as giant HII regions while the heating caused by the ultraviolet light would have inhibited any further stellar formation, thereby defining the efficiency of the process. In other words, by then, all matter that had to go into stars, had gone into stars already. The little remaining gas exposed to the activity of the massive stars, to their winds and supernova explosions, was then dispersed, or blasted away, leaving all stellar sources (massive and low-mass stars) moving under the gravitational potential of the newly-formed cluster. It would have been then, during 50
Tenorio-Tagle: Violent Star Formation
51
the lifetime of the massive stars, that the gas velocity dispersion (o-ga,) had been found to be correlated with the size of the HII regions in the same manner as one finds today that the stellar velocity dispersion (cr») correlates with the size (r) of spheroidal stellar systems (a* ~.Ar 1 / 2 ). Going even further back in time, before the appearance of massive stars, the clouds must have been infested by a large number of low-mass stars, the stars that we see today, all moving within the gravitational potential of the collapsing clouds. Many of these sources undergoing winds while ramming through the leftover cloud, caused a continuous stirring of the leftover gas, inhibiting the further collapse of the clouds as a whole, and at the same time providing their distinctive structure, or cloud filling factor. This however, could have happened only once the cloudlets resultant from fragmentation had acquired stellar sizes (M* ~ MQ). Surely before then, fragments were too large and too massive to form stars and thus the collapse of the clouds and their fragmentation unavoidably continued. It has proven difficult to trace the history of clusters. However, during the last few years a major effort has been channelled into understanding all possible clues, particularly those offered to us during one of the most obvious stages, during the giant HII region phase. These, if rightly interpreted, may offer a definite path towards understanding the cluster formation sequence. Giant HII regions have been studied as large-scale events showing us large sections of galactic disks - the threshold column density - capable of undergoing major bursts of stellar formation. Also, as entities demanding a large number of OB stars to account for their large-scale ionization and excitation, and they are indeed the best laboratories to study the mechanical energy deposition from massive stars. Several studies have also centred their attention on the detected supersonic turbulence. Without exception, the line profiles of giant HII regions imply a supersonic velocity dispersion (cr) in the range 15-40 km s"1 , and the nature of these supersonic motions, also termed supersonic turbulence, has received several interpretations during the last twenty years. These have attempted to point to the energy source, the essential ingredient in all models, capable of providing continuous replenishment. This comes from the fact that supersonic motions cause shock waves that rapidly dissipate energy; thus, a continuous injection of energy seems unavoidable if one is to maintain the shock waves and supersonic turbulence, at least during the lifetime of the HII region. Some authors have suggested the action of multiple strong stellar winds from massive stars as the driving energy source (e.g. Rosa & Solf 1984; Dyson 1979; and more recently Kennicutt & Chu 1994). Others have invoked the champagne expansion of the HII regions as responsible for the supersonic flow (Skillman & Ballick 1984; Hunter & Gallagher 1985). There is no doubt that both mechanisms do lead and contribute to supersonic motions within the ionized gas. However, none of these models or interpretations has attempted to offer an explanation of the observed velocity dispersion vs. size correlation. On the other hand, Terlevich & Melnick (1981) have suggested that supersonic turbulence in giant HII regions ultimately ought to have a gravitational origin. This conclusion seems unavoidable if one looks at the positions occupied by giant HII regions in the log r vs. logcr diagram, with globular clusters and the central regions of elliptical galaxies at opposite ends of the correlation. This fact was originally interpreted as a collection of ionized clouds moving supersonically under the gravitational potential of the newlyformed stellar cluster, a model often referred to as the "gravitational model". This is, however, totally ill-founded as ionized clouds should be completely disrupted through champagne-likeflowsin a few sound crossing times, and hence the model cannot explain the persistence of supersonic turbulence throughout the lifetime of giant HII regions. The only way out is to generate these supersonically moving clouds continuously during
52
Tenorio-Tagle: Violent Star Formation
the life time of the exciting stars. The good news however, is that perhaps everybody was right! According to the cometary stirring model (see Tenorio-Tagle et a\. 1993, hereafter referred to as paper I), gravity causes the stellar velocity dispersion, and the stellar winds from the many low-mass stars to agitate the left-over cloud leading to supersonic turbulence and density enhancements and thus to the observed correletation. This possibility surpasses the obvious boundaries of energy deposition by massive stars. It goes beyond ionization and excitation into the true "realm of the nebulae". It goes beyond the search of a threshold column density of matter required to form stars into the physics of gravitational collapse. On the whole, it provides us with a definite insight into the physics of cluster formation.
2. The quasi-steady gravitational collapse Perhaps a simple but valid approach to the problem is to follow the collapse of the original cloud as if it happened through a sequence of quasi-steady stages, all of which could be described by a simple analytical formulation. In such a case, one can define a cloud of radius r c and mass Mc whose gravitational potential (U) is given by
The cloud will have an internal pressure (Pc) given by 217/2 p
~
3 ~
A fragment within the cloud with temperature Tj and primarily thermal support has Pc as an external pressure and a central pressure: nfkTf
^Pc+
GMjfL,
(2 .3)
where nj = pj /m. A fragment is gravitationally unstable only if rj
2mG
K
'
or if Pc > \nfkTf,
(2.5)
and from these the well known result follows that self-gravitating fragments must have column densities comparable to those of the cloud as a whole: r}pj ~ rcPc.
(2.6)
The behaviour of a cloud during its quasi-static collapse includes processes involving the release of magnetic flux and angular momentum, generation of supersonic motions, dissipation of energy by those motions, shrinking of the cloud as a whole, and fragmentation into smaller gravitationally-bound objects. But, on the whole, the quasi-static collapse is supposed to fulfil at all times the virial theorem, and thus the characteristic speed of fragments in the potential well of their parent cloud will be
Tenorio-Tagle: Violent Star Formation
53
This, should be lost upon encountering a similar amount of matter, and thus, given the comparable cloud and fragment column densities, it should be lost in a cloud-crossing time (txing = 2 r c / v / ) . To sense the unavoidable fate of the matter involved, one should realise that as the quasi-static collapse proceeds,
the cloud would shrink further as the potential well deepens and new fragments should move faster despite their losses. The column density would increase and smaller fragments should still form with the same column
density.
Following the above equations (see paper I), one could show that the fragment radii and masses will follow
and
and that, thus, rj, and Mj decrease rapidly as the collapse proceeds and r c gets smaller. 2.1. The day the collapse stopped
One could then in principle continue to cut smaller and smaller fragments out of the cloud almost forever, or for as long as the collapse continues. However there comes a moment, the magic moment, when this is no longer possible, a moment when the collapse of the cloud as a whole is stabilized. We have inferred this from the observations, because we know that the collapsing matter did not disappear into a massive black hole at the centre of the configurations, but that instead thousands of stars formed and clusters present today a core and a tidal radius. Two conditions are needed to halt, or stabilize, the collapse. The first is that fragments should eventually acquire stellar sizes and turn into stars. At that moment, given their negligible cross-section, fragments can fully detach from the collapsing flow and thus are able to continue to move for the first time undecelerated through the left-over gas. The second condition is that the newly-formed suns undergo stellar winds. This latter event drastically modifies the flow as the left-over gas would feel for the first time an extra pressure, that of the winds ramming through it, and thus instead of causing the deceleration of the fragments, now, the newly formed stars sweep, condense and accelerate the gas up to the speeds acquired from the gravitational potential.
54
Tenorio-Tagle: Violent Star Formation
Given the above result for fragment masses, star formation at mass Mf = Af» becomes possible only when the cloud reaches a radius 2mG .,, c
, , .1 /•)
[M Mc)
'* ~ ZkTj * ' at which point both it and the proto-stellar fragments have column densities o
/ o i oi
\
(2J0)
2
The cloud pressure is Pc ~ ^{Pcrcf
(2.12)
and the characteristic velocity of the forming stars in the gravitational well is
Given that the cloud has stabilized, stars will form over a relatively long period of time, during which the characteristic value of v* remains as above. During this same period, fragments within the stabilized cloud will have similar velocities. A consequence of this is that one can identify the stellar and cloud gas dispersion velocities with v*, which, from the above relationships, implies either a critical pressure, or a critical column density, and/or a critical temperature as the cause of the observed
1/2
This implies a value of A ~ 1.1(T//10/<)(M»/MQ) 1 ^ 2 , while the observed correlation leads to A ~1.7. 2.2. The stabilizing agent The stabilization of the cloud collapse implies that the dispersion velocity of the gas and of the stars forming from it will be characterized by the depth of the gravitational potential, a ~ ( G M c / r c ) 1 / 2 . The gas is compressed and accelerated by the collection of bow shocks caused by the (supersonic) motion of multiple stellar wind sources. TTauri winds have already been propossed as a mechanism to account for the energetics, dynamical structure and lifetime of dark molecular clouds (Norman & Silk 1980; Franco 1984). In this view, the interaction of shells of swept up matter leads to a hierarchical clumpy structure with a mass spectrum that defines the evolution of the cloud and out of which an IMF could be inferred for long-lived clouds. However, the supersonic velocity widths cannot be accounted for with the interaction and collision of the wind-driven shells, as dissipative processes would soon erase any trace of the potential well from the system. Thus, newly-formed stars in such scenarios, would not have the velocity
Tenorio-Tagle: Violent Star Formation
55
dispersion required to stabilized the collapse of the cloud. The bow shock hypothesis, or cometary stirring model, could, on the other hand adequately maintain stability if the ensemble of bow shocks reached all parts of the cloud in a cloud-crossing time. If bow shocks have a cross section KR\0W and move with a velocity v*, then TV such shocks will encounter a volume
in one crossing time. In paper I, the TV sources required to maintain the cloud from collapsing any further were estimated and shown to lead to a cluster bolometric luminosity:
The result is that the cluster luminosity is approximately co~A /G, independently of the stellar wind velocity or luminosity, so long as the stars with radiatively driven winds dominate both the bow shock ensemble and the luminosity. The numerical result is = ^ = 0-5 x 10 8 L o ( ° _ t) . (2.17) 9G \ 1 0 km s / In conjunction with our previous result that a ~Arcll2 and equation (14) lead to Lcluster
and, as shown in paper I, this result is in excellent agreement with the observations.
3. The structure of star-forming clouds The cloud filling factor in star-forming regions is one of the long standing and largely ignored puzzles in Astronomy. It is a quantity originally thought to bring into an agreement the root mean square (rms) densities as measured in radio continuum with those determined at optical wavelengths, such as from Oil and/or SII lines. It is supposed to provide us with an idea about the structure of clouds, or at least indicate their degree of dumpiness. The factor of 10 difference between the two densities, typical in giant extragalactic HII regions (Kennicutt 1984), and the even larger values measured in galactic photoionized nebulae as in Orion (Osterbrock & Flather 1959), remain however unexplained. The controversy results, once again, from the well-known fact that photoionized condensations are supposed to be dispersed, or diluted, into the background, in a time scale which is much smaller than the lifetime of the ionizing stars (t* ~ 5 x 106 yr), or the evolutionary time of an HII region. The question is therefore how to explain the measured HII region filling factors, or the higher densities detected at optical frequencies, or otherwise, what replenishes the high-density condensations within the ionized gas throughout the history of the nebula. 3.1. The filling factor of star-forming clouds. The filling factor has been defined as the ratio of the square of the constant density < ne >r m » over the n2e measured at optical wavelengths, and it leads to a narrow range of values, of the order of 10~2 in giant HII regions (Kennicutt 1984). This expression is derived from the Stromgren relationship, which equates the number of ionizing photons
56
Tenorio-Tagle: Violent Star Formation
(Ft) to the number of recombinations in a constant-density (< ne > r m s ) ionized volume; i.e.
F* = 1*^/3 2rms,
(3.19)
which, when compared to the number (N) of cloudlets evenly spaced within the same ionized volume and capable of absorbing the same ionizing photon flux if their higher density is ne, becomes
l
(3.20)
and leads to the volume filling factor (/„) jvr-,3
.-
" 'cloudlet r
_ f
- h -
3
v. 2
_ ^ " e -^rms
HII
—2
/ Q <m
•
{6.ZL)
"e
If the stabilization of star-forming regions against further collapse is derived from the bow shocks and wakes caused by the winds from moving stars, the supersonic passage of the cometary sources must at the same time cause a sudden density enhancement of the swept up medium; i.e. right at the head of the bow shocks the density, if the shocks are isothermal, should reach a compression factor (3.22)
where Mjnach is the Mach number (the shock velocity over the sound speed of the ambient gas). The value of the density will be lower as one measures it away from the symmetry axis, as there only the velocity component perpendicular to the shock will be thermalized while the parallel component will be unaffected by the bow shock. However, even there, away from the bow shock symmetry axis, it will still be larger that the background density nnnThe compression across the cometary shocks would lead to wakes, filaments and/or channels and holes in the cloud, with a density maximum (ne) right behind the heads of the cometary shocks. As soon as this is produced, it would notice its overpressure with respect to the background medium and thus it would begin to expand. Its complete dispersal would takes place along the wakes caused by every comet, which would become wider and less dense as matter acquires values of density similar to those of the background gas. However, at all times the high-density gas would be replenished at the head of the wind-driven sources. If one identifies the background density with the < ne > r m , , then clearly one can also express the volume filling factor as fv
=
774 M
( 3 - 23 )
Mach
and thus conclude that the degree of dumpiness in star-forming regions is directly related to the supersonic passage of stellar wind sources ramming agaist the left-over cloud, and therefore the filling factor, or the structure of star-forming clouds, is directly related to the mass of the newly-formed cluster. A detailed comparison with the observations can be made (see Figure 1) if one sets limits to the temperature of giant HII regions. This, depending on metallicity, usually lies within the range 5000 K < T < 15000 K. Representative examples are the giant HII regions of M51 (Diaz et al. 1991) and low-metallicity HII galaxies (Terlevich et al.
Tenorio-Tagle: Violent Star Formation
57
FIGURE 1. The volumefillingfactor of star-forming clouds. The two solid lines encompass the narrow range of predicted filling factor values derived assuming THII — 5000 and 15000 K. Several well-known giant HII regions are placed (suqares) according to their measured velocity dispersion and their /„ = < ne >j mj /n,, while the filled circles indicate their predicted value according to the Tkn value given in the literature
1991), at opposite ends of this range of values. The temperature range implies sound speeds, (kT(\ + x)/m)ll2, in the range 9 to 15.7 km s" 1 , while the measured gas velocity dispersion for giant HII regions spans from 15 to 40 km s" 1 (see Table 1 of Hippelein 1986 or Arsenault & Roy 1988). Figure 1 shows the corresponding narrow range of predicted filling factor values (/„ = 1/Mj^,,,.,,) across the a range of applicability of these sources. Also shown are some of the classical giant HII regions for which accurate determinations of temperature, and thus sound speed, as well as gas velocity dispersion are available, together with filling factor values derived from the ratio < ne >r m «/ n eNote that prior to the birth of massive stars the compression factor and the dispersal time-scale of the shocked gas are both large, due to the low sound speed (CHQ ~0.3 km s"1) in the cloud. However, the birth of the ionizing sources marks the ensuing evolution (see Figure 2) and determines the structure of the parental cloud depending only on the total mass of the newly-formed stellar cluster. Clouds ionized by a low-mass cluster (Mciuater < 105 M©), with a correspondingly small c (< CHII) value, will cause the cometary sources to become subsonic upon photoionization. In this case the cloud structure ceases to be
58
Tenorio-Tagle: Violent Star
Formation
restored and thus that revealed to an observer is a consequence of the earlier evolution. In other words, photoionization would rapidly wash away, through local champagne flows, any sign of the cloud structure within the ionized volume. However, the original cloud matter (stirred prior to massive star formation) will continuously become apparent as the ionization fronts proceed into the cloud, showing its original structure (see Figure 2). T h e situation is drastically different in the case of massive stellar clusters, for which cometary sources remain supersonic (a* > CHII), even after complete photoionization of the parental cloud. In this case, the earlier cloud structure is rapidly dispersed, but the cometary sources will continuously replenish a higher-density medium causing, until complete cloud disruption, a characteristic cloud filling factor. This happens despite the continuous rapid disruption, through well-localized champagne flows, that restore the cloud low-density medium as the condensations are washed away to approach pressure equilibrium. The unavoidable fate of violent star-forming regions is cloud disruption. It is the action of massive stars through photoionization, stellar winds and their final supernova explosion that end up dispersing the parent clouds. Clearly, during the evolution of the HII region all sections of the cloud affected by the strong energy deposition will stop showing the effects produced by the presence of a stellar cluster. The effective clearance of matter will erase the supersonic velocity dispersion as well as the cloud structure. Therefore, as a function of time, only the best shielded and/or densest sectors of the parental cloud (which one should realise are not necessarily the central parts of the cloud), will be able to show through true Gaussian line profiles their supersonic kinematic nature, the signature of massive cluster formation. A good example of this is NGC 604 whose central and densest condensations (see the contribution by Munoz Tuiion in this volume and Sabalisck et a/. 1994), although separated by about 50 pc, present almost identical Gaussian profiles with the same value of a. Regions between the densest knots show strongly asymmetric line profiles, a sure signature of the many stellar winds presently disrupting the cloud. These lines can be fitted with broader Gaussians, however, of a much lower intensity than the true Gaussian profiles emanating from the densest knots. Therefore, for as long as this continues to be the case, for as long as some sections of the parent cloud survive the action of the massive stars, single aperture observations including the densest region of NGC 604 will show the supersonic crgas signature of the cometary stirring in what is still left from the parent cloud. 4. Conclusions In the cometary stirring model the cloud agitation is caused by the continuous supersonic passage of isothermal bow shocks, or "cometary" shocks, generated by the stellar wind sources ramming through the left-over cloud. These maintain the cloud stable against further collapse while causing supersonic turbulence and at the same time generating a distinct structure of the remaining cloud. All this becomes possible in the initial phases of stellar formation once the fragments resultant from cloud collapse become small enough, acquire stellar dimensions and form stars. This event detaches them from the collapsingflow,while their wind activity leads to the formation of cometary shocks. Such sources, given in sufficient numbers as to overrun the left-over cloud in a crossing time, cause a cluster luminosity shown to be correlated with cr4 and furthermore lead to a cluster size correlated with 105MG) the still supersonic passage of the numerous wind-driven sources will continuously restore the gas
Tenorio-Tagle: Violent Star Formation
59
Halo
^
FIGURE 2. Schematic view of star forming clouds. On the left, and as a function of time, is the sequence promoted by low mass clusters (MC|U,«er < 105AfQ ) able to cause the structure or filling factor of their parent clouds only before the birth of massive stars. Upon photoionization the cometary sources become subsonic and stop generating the cloud filling factor. Therefore, the structure observed in these clouds is that revealed to us as the ionization fronts progress into the leftover HI or H2 gas. On the right-hand side are the equivalent stages for massive cluster cases. Here despite photoionization the supersonic passage of the wind driven sources continuously replenish a cloud structure.
60
Tenorio-Tagle: Violent Star Formation
velocity dispersion and the cloud filling factor. It is then that the correlations between size and luminosity vs. a could be recognized. The input of mechanical energy from massive stars, however, will soon begin to upset the balanced situation by disrupting larger and larger sections of the parent cloud, erasing the clues stored in the gas phase during cluster formation. In this way, the observed gas turbulent motions are related to the total gravitational energy of the star-forming system, and in completed stellar systems the relationships are a relic of the former gas phase correlation. I would like to acknowledge many pleasent discussions on these topics with Casiana Muiioz Tunon, and Guido Munch. I also would like to thank the DGICYT (grant PB910531GEFE), the EEC (grant CI1*-CT91-O935) and NATO (grant CRG920198 for collaborative research) for partial finantial support to carry out this project. REFERENCES ARSENAULT, R. & ROY, J.-R. 1988 Astron. Astrophys. 201, 199. DIAZ, A. I., TERLEVICH, E., VILCHEZ, J. M., PAGEL, B. E. J. & EDMUNDS, M. G. 1991
M.
N. R. A. S. 253, 245. DYSON, J. E. 1979 Astron. Astrophys. 73, 132. FRANCO, J. 1984 Astron. Astrophys. 137, 85. HlPPELElN, H. H. 1986 Astron. Astrophys. 160, 374. HUNTER, D. A. & GALLAGHER, J. S. 1985 Astron. J. 90, 1457. KENNICUTT, R. C. JR. 1984 Astrophys. J. 287, 116. KENNICUTT, R. C. JR. & CHU, Y.-H. 1994 Astrophys. J. Submitted. MUNOZ TUNON, C. 1994 This volume. NORMAN, C. & SILK, J. 1980 Astrophys. J. 238, 158. OSTERBROCK, D. E. & FLATHER, E. 1959 Astrophys. J. 129, 26. ROSA, M. & SOLF, J. 1984 Astron. Astrophys. 130, 29. SABALISCK, N. S. P.,
Astrophys.
TENORIO-TAGLE, G., CASTANEDA, H. O., MUNOZ TUNON, C.
1994
J. In press.
SKILLMAN, E. D. &: BALICK, B. 1984 Astrophys.
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TERLEVICH, R. J., MELNICK, J., MASEGOSA, J., MOLES, M. &; C O P E T T I , M. V. F . 1991
Astron. Astrophys. Suppl. 91, 285. TERLEVICH, R. & MELNICK, J. 1981 M. N. R. A. S. 195, TENORIO-TAGLE,
G.,
MUNOZ TUNON,
839.
C. & Cox, D. P. 1993 Astrophys. J. 418, 767.
Super-Associations as Star Complexes with Violent Star Formation By YURI N. EFREMOV Sternberg Astronomical Institute, Universitetskii Prospect 13, 119899, Moscow, Russia
1. Super-associations and star complexes The concept of a "superassociation" was first introduced by Baade (1963) in his Harvard lectures in 1958. He gave this name to a region about 500 pc across around the giant HII region 30 Dor in the LMC which is full of OB-associations. The same region was the first example of a super-association given by Ambartsumian (1964). Altogether, 19 OB-associations and young clusters here form a morphological unit of 1 kpc in diameter with evident hierarchical structure (Efremov 1988, 1989). One may say that a "super-association" is the counterpart of a hydrogen emission nebula (HII region) in B, V etc. broad bands. Wray and de Vaucouleur (1980) have shown that in the B bandpass the continuum-to-emission ratio is always greater than 10:1. Thus in this band one deals mainly with the star population of a super-association. Nevertheless, the diameter of the 30 Dor HII region is only 250 pc and it occupies less than 0.1 of the total area of the super-association, the remaining HII regions here being much smaller. This may well also be the case for extragalactic super-associations - giant HII regions. In bright star cloud NGC 206 in M31, named by Baade (1963) as a real super-association, there are only a few small HII regions, not seen at all on the B plates. In many respects super-associations (SAs) are similar to common star complexes (Efremov 1978, 1993), the main difference being the richness of an SA in HII gas and OB stars that causes the high total luminosity. Whether there exists a clear borderline between these complexes and SA, either in size, luminosity or mass, remains still to be investigated, though it is likely that the average SA is larger than a common complex. As yet, the luminosity function of star complexes is still unknown. Data for the total sample of star complexes in a number of galaxies are necessary in order to attain this goal . Those objects with lower luminosity have usually not been considered at all and only those isolated patches in galaxies were searched and named as SA whose luminosity is above B = —14 (Ambartsumian 1964; Petrosian, Saakyan & Khachikian 1984). In our opinion SAs have the same nature in principle as common star complexes but with an unusually high rate of (massive) star formation over the total area of the complex. Within a super-association there are a dozen or two OB associations and young massive clusters, whereas within a common complex of similar size the number can range from zero to three or four associations (Efremov 1988, 1989). In what follows we accept the point of view that complexes are formed from giant superclouds forming in a gaseous-stellar galactic disk owing mainly to gravitational instability (Elmegreen k Elmegreen 1983; Elmegreen 1987; Larson 1988; Efremov 1988, 1989, 1993; Elmegreen et al. 1994). According to Kennicutt (1989) the threshold gas density for massive star formation corresponds to just the Jeans mass and wavelength in the gaseous disk, which are in fact observable parameters of superclouds (Rand 1993). 61
62
Efremov: Star Complexes and Super-Associations
2. Reasons for violent star formation in super-associations 2.1. Increased velocity dispersion
There must be special conditions for violent star formation to take place over a whole complex. We have in total no more than two or three SAs within the Local Group (OB 78 = NGC 206 in M31, 30 Dor in LMC and NGC 604 in M33), while there are about 270 star complexes in these three galaxies. This ratio means that SAs are not simply young complexes, because the mean ages of SAs and complexes relate as 1 to 10. (One should note that only within the Local Group of galaxies may complete samples of complexes, including those without active star formation, be picked out, and that even for these galaxies the distributions of complexes are known in size only, not in luminosity.) We suggested that one probable condition for violent star formation in a complex is high-velocity dispersion in a parent supercloud. According to Elmegreen (1992), and Elmegreen, Kaufman k. Thomasson (1993), a cloud with higher internal dispersion should produce stars with greater efficiency because of the cloud's greater resistance to selfdestruction under energy input from new-born O stars and SNe. Such clouds should also produce a larger fraction of massive stars owing to more heating of star-bearing gas because of higher star formation efficiency. The former authors gave a number of examples of superclouds with violent star formation in interacting galaxies. Some disturbance of the cloud- bearing interstellar medium should exist that leads to higher velocity dispersion in superclouds resulting from the ISM (Elmegreen 1992). These superclouds are larger than those usually found in a normal undisturbed galactic disk. Super-associations are common within interacting galaxies but in normal galaxies they usually located at some special position in a spiral arm, such as at the beginning of a spur (as in OB 78 in M31) or at the end of an arm (Petrosyan et al. 1984). Within irregular Magellanic-type galaxies SAs are often located at the end of a bar (30 Dor in the LMC, NGC 2363 in NGC 2366 etc). In such locations some kinematical disturbance of the interstellar medium and higher velocity dispersion in superclouds is also possible. Indeed the star velocity dispersion within the 30 Dor complex is 14 km s" 1 , whereas in other concentrations of supergiants in the LMC it is only 6 km s" 1 . This large dispersion in 30 Dor is explained by three overlapping velocity fields (Martin et al. 1984). There are some data on the high dispersion of HII velocities inside super-associations, though whether it is connected with turbulence or gravitational motions is still a controversial issue. This is reviewed, for example, in Efremov (1989). The arguments of Terlevich & Melnick (1981) and Melnick et al. (1987) in favour of the virial nature of these motions seem attractive and we consider their data as one more piece of evidence that star complexes are born as physically connected entities from a self-gravitating supercloud and are not random agglomerations of clusters, stars and clouds. 2.2. Density and pressure enhancement
It is quite possible that other initial conditions might lead also to more or less violent star formation. Density and temperature within the parent superclouds of SAs are probably similar to those in a spiral density wave and a preference for high-mass star formation was argued for by many authors (e.g. Mezger 1987). A drastic increase in the gas density, owing to cloud collisions or to external pressure (such as observed inside galactic spiral density waves) may lead to a burst of star formation (Larson 1987), and in a galaxy with a strong spiral wave (such as M51, M83 and NGC 6946) some star complexes situated along arms may be well called SAs. It is quite probable that in every situation the transition between SAs and common complexes is a continuous one. The best examples of intermediate entities are provided by the Andromeda galaxy (Battinelli et al. 1993).
Efremov: Star Complexes and Super-Associations 2.3.
63
Triggering by shock-shock collision
Shock-shock collision within a supercloud leads to enhanced reflected shock waves that propagate with high velocity and trigger additional star formation (Chernin k Efremov 1993). There must exist within a parent supercloud at least two simultaneously acting star formation centres producing shock waves in the interstellar medium to produce violent star formation under the action of this mechanism. If there are a number of such centres, the resulting pattern of age gradients in a superassociation must be very complicated and much work still needs to be done before a reliable comparison with observations is possible. Nevertheless, within small parts of the 30 Dor super-association, age progression is observed from the star-forming molecular cloud near the HII region NGC 2074 to the older dispersed association NGC 2081 (Lortet & Testor 1988). Farther away the still older blue globular cluster NGC 2100 is located in the same line. Probably there are some signs of an age gradient in a certain direction, compatible with Chernin k. Efremov's (1993) expectations. This mechanism naturally explains the large velocity dispersion but probably not the preferred positions of super-associations. 2.4. Abundance peculiarities inside superassociations? Another important issue is the possible influence of abundance on the star formation rates, efficiencies and especially the initial mass functions inside different complexes. There exists a strong concentration of WN to young star complexes both in our Galaxy (Alfaro, Delgado k Cabrera-Cafio 1992) and in the LMC (to the 30 Dor complex, Moffat et al. 1987), whereas stars of WC subtype do not show such behavior. It is improbable that WC stars are simply older because there are indications that they are more massive than WN stars (Vazquez k Feinstein 1990). The ratio of WN/WC stars is larger in galaxies and in regions of a galaxy with lower metallicity (Maeder 1987; Arnault et al. 1989) and one suspects that brighter (i.e. with higher SFR) complexes may have lower metallicity compared with the bulk of star complexes. This difference in abundance may have something to do with the initial luminosity function and lead to higher resultant star formation efficiency as well, for lower-metallicity protoclusters should be more successful in resisting the disruptive forces from O-stars and supernovae (Elmegreen 1983). Here the same arguments act as in the case of large velocity dispersion. Lower metallicity might also lead (via influence on mass loss or even initial mas function) to a preponderance of WN stars in complexes with more active star formation. Then the more frequent occurrence of SAs at the periphery of a spiral galaxy may be connected with the lower abundance there. By the way, the blue (not the red or WR star) progenitor of SN1987A in the LMC is also more compatible with lower abundance in the LMC and probably especially in the 30 Dor super-association. Moreover, the star is suggested to be a member of a small cluster and it is quite possible that stars in clusters generally have lower abundances than surrounding field stars just because the latter were born in already dissolved - owing to larger metallicity? - associations (Efremov 1991). There remains of course the difficult problem of why some complexes of 1-kpc scale have lower abundances than that of the general field - or those of other complexes. One may guess that the mixing of the interstellar medium within a galaxy disk is not so effective as one usually thinks, and that some relics of primordial or, more plausibly, subsequent merging are alive now. There may be collisions of extragalactic clouds with the gas of the galactic disk leading to star formation accompanied by the deviation of the galactic disk from the mean plane (Alfaro et al. 1991). These clouds should have generally lower abundances than galactic disk gas, because there has been little or no preceding star formation there.
64
Efremov: Star Complexes and Super-Associations
Surely these considerations do not relate to starbursts in galactic centres where the most interesting phenomena probably connected with violent star formation in high metallicity medium have been observed (Terlevich 1993). Unfortunately we have too many explanations for violent star formation phenomena. Is any one of these correct?
REFERENCES ALFARO, E. J., CABRERA-CANO, J. & DELGADO, A. J. 1991 Ap. J. 378, 106. ALFARO, E. J., DELGADO, A. J. & CABRERA-CANO, J. 1992 Ap. J. Lett. 386, L47.
AMBARTSUMIAN, V. A. 1964 In IAU Symp. No. 20: Galaxy and Magellanic Clouds, p. 122.
Canberra, Austral. Acad. Sci.. ARNAULT, P H . H., KUNTH, D. & SCHILD, H. 1989 A. & A. 224, 73
W. 1963 Evolution of Stars and Galaxies. Harvard Univ. Press. P., EFREMOV, YU. N. & MAGNIER, E., 1993 This conference. CHERNIN, A. D. &; EFREMOV YU. N. 1993 This conference. BAADE,
BATTINELLI,
EFREMOV, YU.N. 1978 Sov. Astr. Lett. 4, 66. EFREMOV, YU.
N. 1988 Soviet. Sci. Rev. E: Astrophys. Space Phys. Rev. 7(2), 105. Harwood,
London. N. 1989 Ochagi zvezdoobrazovania v galactikakh (Origins of Stars Formation in Galaxies) Nauka, Moscow.
EFREMOV, YU.
EFREMOV, YU. N. 1991 Pisma Astr. Zh. 17, 404.
EFREMOV, YU. N. 1993 Star Formation, Galaxies and Interstellar Medium (ed. J. Franco et al.), p. 360. Cambridge Univ. Press. ELMEGREEN, B. G. 1983 M. N. R. A. S. 203, 1011. ELMEGREEN, B. G. 1987 Ap. J. 312, 626.
ELMEGREEN, B. G. 1992 In Star formation in Stellar systems, III Canarian Winter School, (ed. G. Tenorio-Tagle et al.),p 381. Cambridge Univ. Press. ELMEGREEN, B. G. & ELMEGREEN D. M. 1983 M. N. R. A. S. 203, 31. ELMEGREEN, B. G., KAUFMAN, M. & THOMASSON, M. 1993 Ap. J. 412, 90.
ELMEGREEN, D. M. ET AL. 1994 Ap. J. 425, in press. KENNICUTT, R. C. 1989 Ap. J. 344, 685.
LARSON, R. B. 1987 In Starbursts and galaxy evolution (ed. Trinh Xuan Thuan et al.), p. 467. R.B 1988 In Galactic and extragalactic star formation (ed. R.Pudritz and M.Fich), p. 459.
LARSON,
LORTET, M.-C. & TESTOR, G. 1988 A. & A. 194, 11.
MAEDER, A. 1987 In Starburstst and Galaxy Evolution (ed. Trinh Xuan Thuan et al.), p. 107. MARTIN, N. ET AL. 1984 In IAU Symp. 108: Magellanic Clouds, p. 137. MELNICK, J., MOLES, M., TERLEVICH, R. & GARCIA-PELAYO, J.-M. 1987 M. N. R. A. S. 226,
849. MEZGER, P.C. 1987 Publ. Astr. Just. Czech. Acad. Sci. 69, 91. MOFFAT, A. F. J., NIEMELA, V. S., PHILLIPS, W., CHU, Y.-H.& SEGGEWISS, W. 1987 Ap. J. 312, 612. PETROSIAN, A. R., SAAKVAN, K. A. & KHACHIKIAN, E. E. 1984 Astrofisika. 21, 57. RAND, R. J. 1993 Ap. J. 404, 593. TERLEVICH,
R. 1993 This conference.
TERLEVICH, R. & MELNICK, J. 1981 M. N. R. A. S. 195, 839.
VAZQUEZ, R. A. & FEINSTEIN, A. 1990 Rev. Mexicana Asti. Astr. 21, 346. WRAY, J. D. & DE VAUCOULEURS, J. 1980 A. J. 85, 1.
Violent Star Formation Driven by Shock-Shock Collisions By ARTHUR D. CHERNIN AND YURY N. EFREMOV Sternberg Astronomical Institute, Moscow University, Moscow 119899, Russia Regions of violent star formation such as supergiant HII regions or superassociations often reveal a binary space structure: they contain two separate components within which very intensive star formation proceeds more or less simultaneously. This observational fact suggests an evolutionary scenario for the phenomenon, in which the key role is played by a strong collision of shock fronts produced by the energy release of the previous generation of massive stars in the region.
1. Binary stuctures and nonlinear gas dynamics There are well observed giant regions of intensive star formation - superassociations that consist of two (or three) components: Per-Cas, Sco-Cen and Car in the Milky Way, OB 78 in M31, Region IV = 30 Dor E + 30 Dor W in the LMC (Efremov 1988, 1989). A dust lane is observed in some cases between the two parts of the region that makes this composite structure especially obvious. In regions like OB 21 in M31, HII clouds give an even more contrasting picture when they concentrate at the two opposite sides of the dust lane. Can this binary spatial structure be a clue to the physical nature of the violent star formation phenomenon? We assume that the answer to this question is positive and present it here in the form of an evolutionary scenario in which shock-shock collisions in the interstellar gas play a key role. In a simple version, the scenario may consist of three basic stages: 1. Two "ordinary" star formation bursts occur in two separate but neighbouring areas of interstellar matter and produce two expanding shocks; 2. The shocks collide and two reflected shocks appear moving in the opposite directions; 3. Extremely intensive formation of massive stars proceeds in the very dense gas (shocked twice!) behind the reflected fronts due to gravitational and hydrodynamic instabilities.
2. The physics of shock-shock collision The initial stage of this picture is similar to triggered star formation as suggested by Elmegreen and Lada 1977 (see also recent reviews by Tenorio-Tagle & Bodenheimer 1988; Elmegreen 1992). The two other stages are characteristic of the shock-shock scenario. The shocks from the neighbouring star formation events can come into effective interaction, if the distance between the centres of star formation is not too large and the time difference in the output of energy from the stars is not too great. This unity of place and time leads to the unity of action. Our study of the formation and evolution of the shock-shock interaction is based on the classic work by Courant & Fridrichs (1948) and also on a set of computer simulations of the process performed at Ioffe Physical-Technical Institute (Barausov, Voinovich & Chernin 1988, 1992; Chernin 1993). 65
66
Chernin k. Efremov: Violent Star Formation Driven by Shock-Shock Collisions
According to these results, the shock-shock interaction gives rise to the formation of a regular dynamical pattern in the zone of the contact collision of the fronts. This structure includes, besides the reflected shocks that are mentioned already, a Mach ring-shaped front and two contact jumps (for the geometry of the pattern see the figures in the references above). The dynamics of the reflected shock which is of special interest here is critically affected by the character of the matter distribution in the area where the shock propagates. This density distribution is prepared by the initial shocks that compress the major part of the gas into two spherical shells behind their fronts. For the reflected shock, it is the medium of decreasing density along the direction of the shock propagation. The shock in the medium of decreasing density is known to move with an increasing velocity (Chisnell 1955; Whitham 1958). In addition, the accelerating shock in this gas distribution proves to be self-collimating (Chernin 1994): the central area of the front moves along the direction of the steepest density decrease, and so its velocity increases in the most rapid way. As a result, the surface of the front becomes more and more convex to the direction of its propagation. As a result, the gas flow behind the front gets stream-like or even jet-like with time.
3. Back to astronomy The reflected fronts drive the gas back to the areas of the initial star formation bursts. This opens a way to two new, more or less simultaneous, and - most importantly - much larger bursts of star formation in the same two areas. In the regions of violent star formation that might form and evolve in this manner, we can expect: A) Two populations of stars in each component of the binary star forming region with an age difference of about 10-30 Myr that is related to the time interval between the initial star formation event and the secondary one behind the reflected shocks; B) Different space distribution and kinematics of the stars of the two populations, reflecting the initial condition in the gas from which they formed in each of the two components: more or less isotropic distribution of the first generation stars and a streamlike distribution of the second generation stars; C) The zone of the shock-shock interaction bounded by the Mach ring and the surfaces of the tangent jumps accumulates dust particles of the face ridge of the initial shocks; there are no outflows here along the line connected the centres of the two initial starbursts, but there are outflows in the directions perpendicular to this line due to the expansion of the Mach ring front; because of this the dust forms a shell here which is observed in the projection as a strip between the two components of the region. Features that are fairly similar to A and C are definitely or probably observed in some supergiant HII regions and superassociations (see Efremov 1982, 1988, 1989, 1993). Features like B may provide an observational test for the mechanism of violent star formation assumed in the scenario. The physics of shock-shock collisions may suggest, perhaps, a basis for a better understanding of other observed violent processes associated with star formation. One example is small-scale supersonic bipolar structures (a few parsecs in size) around young stellar objects. Highly collimated and accelerating jets observed in these objects may be due to reflected shocks produced by shock-shock collisions (Chernin 1994). Another example may be related to the larger space scales. As was first argued by Shklovsky (1984), many active galactic nuclei are most probably the regions of violent star formation (see also Terlevich & Melnick, 1985). It is instructive from the point of view
Chernin & Efremov: Violent Star Formation Driven by Shock-Shock Collisions 67 of the present discussion (for more details see Efremov & Chernin 1994) to consider that active galactic nuclei often or even always display a binary (multiple) spatial structure when the resolution is high enough in the observations (e.g. Heap et al. 1993; Vacca et al. 1993).
REFERENCES D. I., VOINOVICH, P. A. & CHERNIN, A. D. 1988 Preprint No. 1274, Ioffe PhysicalTechnical Institute , Leningrad. BARAUSOV, D. I., VOINOVICH, P. A. & CHERNIN, A. D. 1992 Sov. Astron. Lett. 18(12). BARAUSOV,
CHERNIN, A. D. 1993 A. k A. 267, 315.
A. D. 1994 In press. R. F. 1955 Proc. Roy. Soc. A223, 250. COURANT, R. li FRIDRICHS, K. O. 1948 Supersonic Flow and Shock Waves. Interscience. EFREMOV, YU. N. 1982 Sov. Astron. Lett. 8, 663. EFREMOV, YU. N. 1988 Stellar complexes (Soviet. Sci. Rev. E: Astroph. Space Phys. Rev. 7(2), 105). Harwood. EFREMOV, YU. N. 1989 Ochagi zvezdoobrazovania v galactikakh (Origins of Star Formation in Galaxies). Nauka. EFREMOV, YU. N 1993 This conference. EFREMOV, YU. N. & CHERNIN, A. D. 1994 In press. ELMEGREEN, B. G. 1992 In Star Formation in Stellar Systems. HI Canarian Winter School (ed. G .Tenorio-Tagle et al.), p. 381. Cambridge University Press. CHERNIN,
CHISNELL,
ELMEGREEN, B. G. & LADA, C. J. 1977 Astrophys. J. 214, 725.
S. R. ET AL. 1993 This conference. SHKLOVSKY, I.S. 1984 Sov. Astron. Lett. 11, 163. TENORIO-TAGLE, G. & BODENHEIMER, P. 1988 Ann. .Rev. Astton. Astrophys. 26, 146. HEAP,
TERLEVICH, R. & MELNICK, J. 1985 M. N. R. A. S. 213, 841.
W. D., CONTI, P. S. & LEITHERER, C. 1993 This conference. WHITHAM, G. B. 1958 J. Fluid Mech. 4, 84. VACCA,
The Search for Hierarchical Structure inside M31 Superassociations ByPAOLO BATTINELLI 1 , YURI N. EFREMOV2 AND EUGENE A. MAGNIER3 1 2
Osservatorio Astronomico di Roma, Viale del Parco Mellini 84, 1-00136 Roma, Italy
Sternberg Astronomical Institute, Universitetskii Prospect 13, 119899, Moscow, Russia
3
Astronomical Institute "Anton Pannekoek" and Center for High Energy Astrophysics, Kruislaan 403, 1098 SJ Amsterdam, The Netherlands and Center for Space Research and Departament of Physics, Massachusetts Institute of Technology, USA
The identification and classification of young star groups in other galaxies is still a controversial topic (Battinelli 1991a). Very different estimates of OB association sizes in different galaxies were explained by Hodge (1986) by a difference in linear resolution and limiting magnitude, but the existence of two kinds of resolved star groups is also essential in this issue. The bulk of OB associations are members of larger groups, star complexes (Efremov 1978), of diameter 400-1000 pc that also include individual fainter and older stars, such as Cepheids. There exists a hierarchical, embedded sequence of young star groups: there are associations, aggregates, complexes and supercomplexes (regions). Associations and complexes are the more common ones (Efremov 1988, 1989, 1993). Given sufficient resolution and limiting magnitude, one can see both complexes as well as brighter and smaller associations mainly inside complexes, and such is the case for M31, where Efremov, Ivanov & Nikolov (1987; hereafter EIN) found, by eye, OB associations of typically 80 pc diameter and star complexes of typically 600 pc. The latter complexes are mainly the same large groups of blue stars that were identified by van den Bergh (1964) under the name of OB associations. The issue arises whether one has simply OB associations with a large range of sizes, the smaller ones being younger as suggested by van den Bergh 1964, or if there exist two kinds of star groups of different hierarchical level, with younger ones as constituent parts of larger and older ones, as suggested by Efremov (1978, 1989) and EIN. To tackle this matter, the application of an objective method of isolation of resolved star groups is of primary importance (Battinelli 1991a, 1992). This task was carried out recently by Magnier et al. (1993; hereafter MBL) for M31 stars. They applied a variant of a cluster analysis method (the Path Linkage Criterion elaborated by Battinelli 1991b) to a sample of about 7000 blue stars extracted from a set of about 300000 stars for which BVRI CCD photometry was obtained by Magnier et al. (1992). MBL found 174 associations with contamination of lower than 5% by random clumps, the total number of associations (within central 1 deg2) being 280 and average diameter 90 pc. Thus MBL confirm the value found by EIN for the size of M31 associations, very similar to the sizes of associations in the Milky Way and Magellanic Clouds. In MBL's list of M31 associations, there are also half a dozen larger groups such as NGC 206 (van den Bergh's OB 78) with large diameters. These groups, having larger densities of blue stars, are intermediate between common star complexes and superassociations. Within OB 78 (= OB 13 in MBL, with a size of 580 pc - a genuine superassociation), MBL identified 8 clumps in a second hierarchical level with diameters of about 100 pc. In the present contribution we apply the second level PLC to six other clumps with 68
Battinelli et al.: Superassociations in MSI
69
diameters larger than 250 pc from the MBL list. Internal structure was found in each of them, with mean diameters close to 100 pc and with the following numbers of "real" clumps (i.e. those with a probability under 10% of being random): OB 92 (5), OB 136 (2), OB 118 (3), OB 142 (1), OB 155 (3). One may conclude that hierarchical structure is a typical property of the larger clumps of blue stars in M31, which corroborates the suggestion that these large clumps are association complexes. The complexes isolated by MBL using the distribution of blue stars are evidently younger (or with higher star formation rate) than the bulk of complexes in M31, i. e. those older complexes which are not detected by MBL but seen by eye both by van den Bergh and EIN. Fainter stars and/or a larger scale length than that used by MBL should be used to search for star complexes in the Andromeda galaxy with the PLC. Thus, only a few of the brightest complexes are easily seen in the blue star distribution, and the case of M31 may well be typical. Only in the nearest galaxies is the total sample of complexes generally detectable. The implied ubiquity of star complexes which encompass the bulk of associations in disk galaxies (Efremov 1988, 1989, 1993) is in agreement with top-down scenario of star formation according to which the vast superclouds that produce star complexes are the initial structures (Elmegreen 1979, 1987; Elmegreen & Elmegreen 1983).
REFERENCES BATTINELLI, P. 1991a Mem. Soc.Astron. It. 62, 959.
P. 1991b Astron. Astrophys. 244, 69. BATTINELLI, P. 1992 Astron. Astrophys. 258, 269. BATTINELLI,
ELMEGREEN, B. G. 1979 Astrophys. J. 231, 372. ELMEGREEN, B. G. 1987 Astrophys. J. 312, 626. ELMEGREEN, B. G. & ELMEGREEN, D. M. 1983 Mon. Not. R. Astron. Soc. 203, 31.
N. 1978 Sov. Astron. Lett. 4, 66. N. 1988 Astrophys. Space Phys. Rev. 7(2), 105. EFREMOV, YU.N. 1989 Origins of Star Formation in Galaxies: Star complexes and Spiral Arms (in Russian). Nauka, Moscow. EFREMOV, YU. N. 1993 In Star Formation, Galaxies and the Interstellar Medium (ed. J. Franco, F. Ferrini & G.Tenorio-Tagle), p. 360. Cambridge Univ. Press, Cambridge. EFREMOV, YU. N., IVANOV, G. R. & NIKOLOV, N. S. 1987 Astrophys. Sp. Sci. 135, 119. HODGE, P. W. 1986 In Proc. IA U Symp. 132: Luminous stars and associations in galaxies (ed. C. H. W. de Loore et al.), p. 369. Reidel, Dordrecht. MAGNIER, E. A. ET AL. 1993 Astron. Astrophys. 278, 36. MAGNIER, E. A. ET AL. 1992 Astron. Astrophys. Suppl. 96, 379. VAN DEN BERGH, S. 1964 Astrophys. J. Suppl. 9, 65. EFREMOV, YU. EFREMOV, YU.
A Stochastic PSF Model: Smooth Spirals in Differentially Rotating Disks ByJAN PALOUS !AND BRUNO JUNGWIERT 2 'Astronomical Institute, Academy of Sciences, Bocni II 1401, 141 31 Prague 4, Czech Republic 2
Center for Theoretical Study, Charles University, Celetna 20, 110 00 Prague 1, Czech Republic
The propagating star formation model with anisotropic probability distribution is investigated. In each star-forming site we define the probability ellipse and show that its two parameters, the excentricity and the orientation relative to the galactic rotation, are closely related to the thickness and inclination of the resulting spiral arms. The relative size of a star-forming region with respect to the whole galaxy is also discussed. Simulations are compared to the observed galactic morphologies and we mimic the differences between the two groups of galaxies of types M101 and NGC 7217.
1. Propagating star formation The idea that star formation at one place in a galaxy can initiate star formation in its neighbourhood was first suggested by Opik (1953) and Oort (1954). Since then, a possible chain of physical processes that joins two regions of subsequent star formation has been proposed in which ionizing radiation from massive stars in a cluster leads to the disruption of the parental molecular cloud via supersonic champagne flows halting further star formation. The mechanical energy input from stellar winds and supernova explosions causes the agglomeration of gas in expanding shells. Their fragmentation, the building of molecules in high opacity areas, and large-scale gravitational instabilities may produce molecular clouds, where the next generation of star formation occurs. The star-forming cycle described above is the basis of deterministic PSF models (Palous et al. 1994). However, the physical parameters such as density, metallicity and cooling times of the ISM, are only partly known. At the same time, the detailed model of star formation inside the cores of a molecular cloud is not generally accepted. The role of magnetic fields is also poorly understood and the interaction of expanding shells with the ambient medium is still under discussion. The undetermined ingredients in the star-forming cycle can be approached with probabilistic arguments in in terms of stochastic self-propagating star formation (SSPSF) models, in which star formation propagates from place to place with a certain probability. Such a description was introduced by Gerola k Seiden (1978). The SSPSF is a percolation model: it describes large-scale galactic properties resulting from the temporal and spatial percolation of small-scale events (Seiden k Schulman 1990). The galaxy is globally parametrized with the finite probability Po, which decreases locally over the recovery time r after the last star formation in the region's history is taken into account. Apart from Po and r, other two parameters are involved: the velocity of rotation and the relative size of a star forming region with respect to the whole galaxy. The SFR depends on Po and the model shows the non-linear behaviour typical for a phase transition: for Po smaller than a critical value P c the SFR is zero. For Po immediately higher than P c the SFR rises steeply. Spiral shapes similar to those of observed galaxies are formed if Po is in a narrow interval near Pc. 70
J. Palous & B. Jungwiert: Stochastic Propagating Star Formation
71
2. SSPSF model with an elliptical probability distribution In differentially rotating galactic disks, two major deformations of expanding shells are predicted and observed (Palous et al. 1990): a. The originally round shell acquires an elliptical shape. The major axis of the shell becomes larger and rotates with respect to the galactic center direction from 45° to 90°. b. The mass collected by the shell continuously slips behind the overtaking shock, and accumulates near the tips, causing an uneven column-density profile. For SSPSF models, the uneven ISM distribution along the shell periphery suggests an anisotropic probability, which is highest along the shell's major axis. This idea was first proposed by Chiang & Elmegreen (1982). In our model (Jungwiert & Palous 1994) the anisotropy is introduced by a probability ellipse connected to any star forming site. The effective propagation probability in a certain direction is proportional to the length of the chord in the probability ellipse, which is defined by two parameters: 1. the axial ratio 6 : a, and 2. the inclination / of the major axis relative to the galactic center. Thus our SSPSF scheme involves the above two parameters plus the following ones, which were already present in the models of Seiden & Schulman (1990): 3. the propagation probability Po, 4. the rotation curve of the galaxy, 5. the time step At, 6. the recovery time r, and 7. the relative size of a star-forming region with respect to the whole galaxy. The relation of these quantities to the spiral features should be discussed and compared to the observed galactic morphologies. The role of 6 : a and I, as well as their connection to the ISM properties and to the rotation curves is analyzed by Jungwiert & Palous (1994) and is only briefly summarized here. In this paper the influence of the relative size of a star-forming region is also mentioned.
3. Results 3.1. The rotation curve, I and b : a The rotation curve determines the amount of differential rotation at a certain galactocentric distance, which affects the global shape of spiral arms. With an isotropic probability distribution the SSPSF produces hyperbolic spirals with the pitch angle depending on the rotation speed Vrot- Slow rotators create more open spirals than do fast rotators. The ellipticity of the probability distribution also generates spiral-like structures. In the absence of differential rotation it leads to logarithmic spirals, if / is kept constant throughout the galaxy. With both ellipticity and differential rotation included, the global spiral shapes are neither purely hyperbolic nor purely logarithmic. The effect of I is shown in Figure 1. With I increasing from 45° through 65° to 85" the winding up of the spiral arms increases: low / produces more open and shorter spirals than high /. The effect of b : a is demonstrated in Figure 2. The anisotropy of the probability distribution plays an important role when it is sufficiently large b : a < 1 : 2. Higher excentricities (b : a = 1 : 10 or 6 : a = 0, where the latter means that the ellipse degenerates to the abscissa) produce remarkably organized spirals, which are smoother than for isotropic SSPSF.
72
J. Palous k B. Jungwiert: Stochastic Propagating Star Formation
7 = 65°
7 = 45°
•
• • • • • • • • •
10
- 1 0
-10
10
10
7 = 85°
-10-10
10
1. A sequence of I = 45°, 65°, and 85°. Vtot = 150 k m s " 1 , and i : a = 1 : 10 in all three cases. Star clusters of ten different ages are shown, the symbol radius decreases linearly with age. Axes are labelled in kpc. FIGURE
10
-10 -10
b:a=l:l
10
io
10
10
b:a=0
-10
-10
10
FIGURE 2. A sequence of b : a = 1 : 1, 1 : 10, and 0. Vr<>t = 150 km s""1 in all three cases. / = 50° when 6 : a < 1.
3.2. The relative size of a star-forming region
In Figures 1 and 2 the size of a star-forming region is close to 200 pc and the radius of the galactic disk is 10 kpc. For comparison, we present in Figures 3 an 4 the same experiments for a star-forming region twice as small. The spiral features are thinner but more numerous. This was already discussed by Seiden and Schulman (1990) in relation to the difference between flocculent and grand-design spirals. They argue that galaxies with relatively small HII regions tend to have flocculent structure. We can confirm this point in the case of isotropic probability distribution (Figure 4, left). In simulations with anisotropy, the long and smooth spiral arms are present even when star-forming regions are small.
J. Palous & B. Jungwiert: Stochastic Propagating Star Formation
10
7 = 45°
-10 -10
10
10
7 = 65°
-10 -10
73
10
to -^10
10
FIGURE 3. The same as Figure 1, except for the relative size of a star forming region, which is twice as small.
10
-10 -10
FIGURE
b:a=l:l
10
10
-10
-10
b:a=l:10
10
••• ' • •
10
-10
-10
b:a=0
10
4. The same as Figure 2 except for the relative size of a star-forming region, which is twice as small.
4. MlOl versus NGC 7217 We try to compare the output from SSPSF models with anisotropic probability distribution with pictures of observed galaxies. To mimic robust, open, grand-design spirals of MlOl type, moderately low 6 : a (~ 1 : 10) and low 7 (~ 45°) are needed. A simulation with such a probability ellipse and with the rotation curve of MlOl is shown in Figure 5. Rather conspicuous, long and smooth arms (unattainable with the isotropic SSPSF) can be seen as well as more patchy areas (at T = 6 Gyr note the existence of two similar and symetrically located arms, extending over 270° in galactic longitude). Of course, the resemblance to MlOl can be only qualitative since the model is stochastic and the right choice of b : a and I is unknown. On the other hand, to reproduce the tightly wound and broken spiral arms of NGC 7217 type, extremely low 6 : a (~ 0) and high I (~ 75°) values are plausible. A more detailed discussion with figures of both galactic groups is given by Jungwiert and Palous (1994).
J. Palous & B. Jungwiert: Stochastic Propagating Star Formation
74
= 2Gyr
T = 4 Gyr
10
10
10
T = 6 Gt/r
. v. -10 -10
10
10
-10
10
FIGURE 5 . T i m e evolution of a galaxy w i t h 6 : o = l : 1 0 , 7 = 45° a n d t h e r o t a t i o n c u r v e of M101.
5. Conclusions We conclude from Figures 1 and 2 that the anisotropic probability distribution produces long and smooth spiral arms or ring-likefilamentsdepending on 6 : a and /: higher anisotropy leads to more organized and thin spirals, whereas higher I yields less steep spirals or ring-like filaments. Broader description of these results are given in Jungwiert k Palous (1994). We also conclude that with small sized star-forming regions the spiral arms are thin and smooth. They keep the overall shape, length and regularity provided the anisotropy is strong. This is not in agreement with Seiden and Schulman (1990), who relate the small size of HII regions to the feathery (flocculent) appearance of galactic disks. The project was supported by the grant No. 205/93/0090 of the Grant Agency of the Czech Republic. JP thanks organizers for the support at the Conference in Puerto Naos.
REFERENCES W. & ELMEGREEN, B. G. 1982 IBM Research Report RC 9726. H. & SEIDEN, P. E. 1978 Astrophys. J. 223, 129. JUNGWIERT, B. & PALOUS, J. 1994 Astron. Astrophys. Submitted. OORT, J. H. 1954 Bull. Astron. Inst. Nether. 12, 177. OPIK, E. J. 1953 Irish Astron. J. 2, 219. PALOUS, J., FRANCO, J. &; TENORIO-TAGLE, G. 1990 Astron. Astrophys. 227, 175. PALOUS, J., TENORIO-TAGLE, G. & FRANCO, J. 1994 Mon. Not. R. astron. Soc. Submitted. SEIDEN, P. E. & SCHULMAN, L. S. 1990 Advances in Physics 39, 1.
CHIANG,
GEROLA,
Spatiotemporal Pattern Driven by a Self-Regulating Mechanism of Star Formation By ANTONIO PARRAVANO Programa de Postgrado en Astronomia y Astrofi'sica, Facultad de Ciencias, Universidad de Los Andes, Merida, Venezuela.
The spatiotemporal pattern of the regions of active star formation in galaxies is the result of cooperative effects that in many cases produce very organized patterns such as spiral arms. The spiral wave theory has been commonly evoked to explain the spiral pattern, where the enhanced density in the arms promotes the coalescence and condensation of clouds and provokes a strong increment in the star formation rate. Nevertheless, the presence of a self-regulating mechanism of the star formation acting on a large scale can also produce an ordered spatial pattern. Here, we include a self-regulating mechanism based on the sensitivity of the condensation and evaporation of small cool clouds upon the radiation density in the 912-1100 A band (Parravano 1987, 1988) produced by massive stars. Since the UV radiation, produced by short living stars, can affect the condensation and evaporation of distant cool clouds, active star formation in adjacent regions tend to occur out of phase (Parravano, Rosenzweig & Teran 1990). This non-local mechanism of self regulation is applied here in order to study its effect on the spatiotemporal pattern in a one-dimensional array of cells which represents a galactic disk ring. The model also includes several local mechanisms of mass exchange between stars, clouds, warm gas, and hot gas. In each cell, the mass-balance of the self-regulated system is represented by: •
>
.
=
•
>
+ «1 >
•
i
/
+ AW ,
•
»
= >
—*0 >
1.
=(l-or)—
+ Ac,
where £_ = a£s £ c +b£'., and £IT = PEw /(«..)- r E,»(»..)' The variables E a , EC) E w , E/,, E r , and £; are, respectively, the mass surface density of massive stars, clouds, warm gas, hot gas, remnants, and long living stars in the cell. Additionally, t* is the average lifetime of massive stars, a is the fraction of mass ejected by the massive stars during its evolution, a and b denote the coefficients for the induced and the spontaneous star formation rates, and (3 and F are the rate of condensation and evaporation of small clouds per unit mass of warm gas and cloud mass, respectively. The term E^(t —
Parravano: Spatiotemporal pattern
76
0 FIGURE
time
1500
Myr
1. Spatiotemporal pattern of the surface density of massive stars.
respectively. The regulating functions f(uuv) and g(uuv) control the mass transition between the warm gas phase and the cloud phase (Parravano et al. 1990). In order to study the effect of the regulating mechanisms in a multizone model, we assume an array of TV cells in a ring. The cells do not exchange mass but are coupled radiatively via the regulating functions f(uuv) and g(uuv) (Parravano et al. 1990). The radiation density uuv in a cell depends not only on the number of massive stars in the cell but also on those in other cells. Figure 1 shows the spatiotemporal pattern of the surface density of massive stars, E, (»,.?), for an array of N = 50 cells (vertical direction; each 1 kpc long), and 500 equally spaced time moments (horizontal direction; each instant 3 x 106 yr apart). At the initial time, all the cells, except for eight, have the state E, = E* = E h = 0 and £ c = 2 x 106 M 0 kpc" 2 . The cells 1, 21, 24, and 40 have the state E9 = S c = Eu, = 0, and EA = 2 x 106 M 0 kpc~2, and the cells 2, 22, 25, and 41 have the state E, = Ec = EA = 0, and E«, = 2 x 106 M© kpc~2. The different grey tones in the spatiotemporal pattern linearly correspond to black for E s = 0 and to white for £j = 3 x 104 MQ kpc" 2 . Even here, where the model excludes dynamical effects such as density waves, the radiative interaction between the cells tends to force the star formation rate in the cells to oscillate out of phase. The resulting spatiotemporal pattern looks like that of a spiral pattern on a ring. Moreover, when a schematic spiral pressure wave is added to the model (not shown here), the expected pattern only appears when the "dynamic" and the "self-regulated" characteristic times are similar. Therefore, the processes included in the present study could cooperate with the dynamical processes to reinforce and maintain the spiral pattern in models of disk galaxies.
REFERENCES PARRAVANO, A. 1987 Astion. Astrophys. 172, 280. PARRAVANO, A. 1988 Astron. Astrophys. 205, 71. PARRAVANO, A., ROSENZWEIG, P. & TERAN, M. 1990 Astrophys. J. 356, 100.
Detection of an Age Gradient along the z-Axis in a Star-Forming Region B y E M I L I O J. ALFARO 1 ' 2 , JOSE FRANCO 3 , E D M U N D O MORENO 3 AND J E S U S C A B R E R A-C ANO 1 - 4 'Instituto de Astrofisica de Andaluci'a, (CSIC), PO Box 3004, Granada 18080, Spain 2
Astronomy Department, Boston University, MA02215, USA 3
Instituto de Astronomi'a, UNAM, Mexico D.F., Mexico
4
Universidad de Sevilla, PO Box 1045, Sevilla 41080, Spain
The age and location of stellar clusters and Wolf-Rayet stars in the third Galactic quadrant are analyzed. The cluster sample has been divided into three age groups: 1) younger than 107 yr, 2) between 10r and 3 x 107 yr and 3) between 3 x 107 and 10* yr. The mean z-locations of these samples in the central region of the Big Dent display a well defined z-age stratification. The existence of such an age gradient seems to corroborate previous hypotheses suggesting that the star formation activity was, probably, triggered by the same strong perturbation which originated the depression. A model in which the Big Dent originated by the collision of a high-velocity cloud with the Galactic disk is able to reproduce the observed gradient. The analysis of the three-dimensional spatial distribution of a sample of young open clusters within 3 kpc around the Sun led Alfaro et al. (1991) to the discovery that the four nearest supercomplexes, previously detected by Efremov & Sitnik (1988), appear to be located below the formal Galactic plane. The one placed in the third Galactic quadrant (labelled III in ES88) shows the largest and deepest ^-displacement and has been called the Big Dent. The mechanisms able to explain the observed large z-departure as well as the different stages of star formation in the Big Dent need a source of mechanical energy able to inject in the Galactic disk a considerable amount of energy and momentum along the z-axis. The halo-disk interaction, via high velocity clouds (HVCs), has been envisaged as a firm candidate to this power source (Tenorio-Tagle et al. 1986; Franco et al. 1988; Alfaro et al. 1991; Comeron k Torra 1992). Obviously, realistic modelling of HVC impacts relies on the physical constraints derived from the analysis of the structures suspected to originate by this mechanism. The time evolution of the star-forming complexes then becomes a notable clue to the understanding of how this process should work. Here we deal with the history of the star formation in the Big Dent. We analyze the existence and possible origin of a z-age gradient in this supercomplex and compare it with that predicted by a simple analytical model, including the effects of interstellar drag and radiation pressure, of the evolution of a cloud formed in a HVC-disk collision. The model provides the locations, as a function of time, of this star-forming cloud and their newly formed stars. The stars are not affected by the interstellar drag and can detach from their parent cloud, generating the z-age correlation through a series of different episodes of star formation. The sample of open clusters has been split into three groups of age: 1) younger than 107 yr, 2) between 107 and 3 x 107 yr and, 3) between 3 x 107 and 108 yr. The mean age for each group is l ) 4 . 3 ± 3 . 5 x 106 yr, 2) 22.9± 5.9 x 106 yr and 3) 60.7 ± 19.4 x 106 yr. The WN stars have been considered to be coeval with the youngest open clusters. The subsamples were further analyzed to get averaged z values for two different positions inside the area covered by the Big Dent. Our analysis yields the results shown in Figure 77
78
E. J. Alfaro et at.: Age Gradient along the z-Axis
o
i
o
o o
_, ^ — %
1
t—1 /
•—€ J-J'
'_
o i
/
• —
1
/
'
c)
I
/ ¥— /
1
1
±I
/ / -/->
WN stars
;
-
WN stars
>
/ 1
.
20
1
1
40 Age (106 yr)
60
.
_
1
80
FIGURE 1. z—Age diagram for the four object samples. Black squares and open circles indicate the average 2-value at two different positions inside the area covered by the Big Dent. The model solution for a gradient originated by a HVC-disk collision is shown as a dashed line.
1. The signal-to-noise ratio estimated for each set allows us to verify the case for a z-age gradient in this star-forming region with a sound statistical significance. We have modelled the kinematics of a cloud generated by a HVC-disk collision, including both the interstellar drag and radiative pressure. The effective radiative force on the cloud is considered to be produced by newly formed stars. These stars are assumed to be formed at a series of different episodes within the cloud. The solution shown in Figure 1 represents the estimated 2-location for the different burst of star formation, assuming that the present time corresponds to 4 x 107 yr after the previous passage through the plane. The model provides an age gradient which is in reasonable agreement with the observed values, and indicates that an HVC-disk collision is certainly able to explain the vertical structure of the Big Dent. Nevertheless, this result does not discard other possible mechanisms for generating this structure but adds new support to previous hypotheses, suggesting that the star formation activity was probably triggered by the same strong perturbation which originated the depression.
REFERENCES E. J., CABRERA-CANO, J. & DELGADO, A. J. 1991 Astrophys. J. 378, 106. COMERON, F. & TORRA, J. 1992 Astron. Astrophys. 261, 94. EFREMOV, YU. N. & SITNIK, T. G. 1988 Soviet Astron. Lett. 14, 347.
ALFARO,
FRANCO, J., TENORIO-TAGLE, G., BODENHEIMER, P., ROZYCZKA, M. & MIRABEL, I. F.
Astrophys. J. 333, 826. TENORIO-TAGLE, G., BODENHEIMER, P., phys. 170, 107.
ROZYCZKA,
M. &
FRANCO,
1988
J. 1986 Astron. Astro-
Abundances of HII Regions and the Chemical Evolution of Galaxies By MANUEL PEIMBERT, PEDRO COLIN AND ANTONIO SARMIENTO Instituto de Astronomia, UNAM, Apdo. Postal 70-264 Mexico 04510 D. F., Mexico We discuss the observational constraints to the chemical evolution of the interstellar medium, ISM, provided by the abundances of HII regions. We present a review of the results derived from these constraints for the solar vicinity and for metal-poor galaxies. It is found that, contrary to previous results, black holes do not play an important role in the chemical evolution of galaxies. Chemical evolution considerations indicate that substellar objects (M < 0.08 MQ) have a mass density smaller than 0.02 MQ pc~ 3 (2
1. Introduction HII regions are excellent probes of the chemical composition of the ISM of the Galaxy and of other galaxies. From the study of bright HII regions it is possible to derive accurate abundances of H, He, C, N, 0 , Ne, S, and Ar for galaxies that are located at many megaparsecs from us. Two important sources of error in the abundance determinations of bright HII regions are the temperature structure and the fraction of heavy elements that is embedded in dust. The observed HII region abundances can be compared with those predicted by models. To produce a chemical evolution model for a given galaxy the following ingredients are needed: a) a set of stellar evolution models for all masses, b) an initial mass function, IMF, c) a star formation history, and d) a mass flow history. In § 2 we will discuss chemical evolution models, presenting some of the most popular, in § 3 we will compare chemical evolution models with observational constraints for the solar vicinity, and in § 4 we will review some of the ideas that have been proposed to explain the chemical evolution of irregular and blue compact galaxies. 2. Chemical evolution models 2.1. Net yield The heavy-element yield is defined by y(Z) = M(Z)/Mt,
(2.1)
where M(Z) is the mass that a generation of stars ejects as newly-formed heavy elements to the interstellar medium and M* is the mass of the same generation that remains locked into stellar remnants and long-lived stars, where we include the low-mass end of the IMF which might comprise objects that do not become stars. Similarly the oxygen yield is defined by 2/(0) = M(0)/M,. (2.2) The yields in equations (2.1) and (2.2), hereafter net yields, have been also defined as true yields or real yields (e.g. Peimbert 1985, Pagel 1987 and references therein). 79
80
Peimbert et al.: Chemical Evolution of Galaxies
To derive the net yield it is necessary to adopt: a) an IMF, and b) a set of stellar evolution models including the whole mass range expected. The net yields can be compared with "observed yields". The observed yields can be computed by means of: a) the determination of the abundances by mass in a star or in an HII region, b) a model of galactic chemical evolution that should include the initial gaseous abundances by mass and the history of gas flows in and out of the representative volume under consideration, and c) the mass that participates in the chemical evolution process, Mj, and the mass that does not participate in the chemical evolution process, Mni,] Mi, includes: the interstellar mass gas, Mg, the stellar mass, and the low-mass end of the IMF; M n j includes the non-baryonic mass and the baryonic mass that does not participate in the chemical evolution process like primordial black holes. Mj and M n j should be evaluated in a representative volume where the stars formed or the HII region is located. 2.2. Basic equations The basic equations of a chemical evolution model in which the instantaneous recycling approximation, IRA, is assumed are the following (Tinsley 1980, Maeder 1992): dMb
f
(oi\
f
-£- = fi-fo, (2-3) where / / and fo are the inflow and outflow rates. Similarly Mg changes according to: ^
fi-fo,
(2.4)
where R is the mass fraction returned to the interstellar medium by a generation of stars, and ip is the star formation rate. On the other hand, the heavy-element mass fraction in the ISM, Z, evolves via: Mg^ = y(Z)(l-R)il> + {Zi-Z)fI,
(2.5)
where y(Z) denotes the "observed" yield and Zi is the Z value of the accreting material. Equation (2.4) can be rewritten in terms of the gas mass fraction, (J, — Mg/Mb, as: ^
(//-/o)(l-/i).
(2.6)
Time can be eliminated by combining equations (2.5) and (2.6) as follows:
dZ
y{Z){lR)^ + {ZIZ)fI
-(1 For most of the heavy elements the IRA is very good because they are produced by stars more massive than 8 MQ. This is the case for 0 , Ne, Ar, and S, elements that are easily observed in bright HII regions. For C, N, Fe, and He, deviations from IRA have to be taken into account. 2.3. Closed-box model When there are no gas flows into or out of the system, / / = / o = 0, the solution to equation (2.7) is given by
v(z) = ^
(2.8)
Peimbert et al.\ Chemical Evolution of Galaxies
81
Notice that frequently equation (2.8) has been used with the total mass of the system, Mr, instead of Mj, a correct procedure only if M n j is negligible. The closed-box model has also been called the simple model. The yield given by equation (2.8) under the assumption that Mn\, — 0 is called the effective yield (e.g. Peimbert 1985, Pagel 1987 and references therein). The yield of agiven heavy element, yi, is obtained by substituting Z for the mass concentration of element i, X{ in equations (2.5), (2.7), and (2.8). 2.4. Infall Under the assumption that the inflow rate is given by / / = a ( l — R)ip, where a is a free parameter, and integrating equation (2.7) for the case where infall and stellar mass loss are balanced by star formation (Zj = 0,/o = 0, and a = 1), we obtain (Larson 1972, Tinsley 1980)
while for other values of a the yield is given by (Matteucci & Chiosi 1983, Peimbert, Sarmiento k Colin 1994)
(210)
Equation (2.10) is valid if a < 1, otherwise it only applies for /i > / i m i n , where Umin = (a - l ) / a ; when ^ -> nmin, Mb -> oo. 2.5. Outflow of well-mixed material Under the assumption that the outflow rate is given by fo = A(l — R)ip, where A is a free parameter and integrating equation (2.7) for / / = 0 we obtain (Matteucci k Chiosi 1983)
In general when inflow and this type of outflow are occurring and a linear variation of the yield with Z is permitted, the solution to equation (2.7) is given by (Peimbert et al. 1994)
-
a Z l t
(2.12)
where we have assumed that y(Z) = j/o + °-Z. Equation (2.12) is valid if (a — A) < 1, otherwise it only applies for fi > fimin = [(a - A) — l]/(a — A). In deriving equations (2.11) and (2.12) it has been assumed that the fresh Z material ejected by the stars to the interstellar medium is completely mixed with the gas before it is ejected to the interstellar medium, the wind associated with this outflow has been called normal or ordinary. Notice that this outflow reduces the yields of the elements but not their ratios; in particular, it does not affect the AY/AZ ratio (under the assumption of IRA). 2.6. Outflow of Z-rich material It has been proposed that outflow of Z-rich material is present, at least in some galaxies, to explain: a) the large helium abundances derived in some irregular and blue compact galaxies, b) the large AY/AZ ratios derived from samples of galaxies, c) the
82
Peimbert et al.\ Chemical Evolution of Galaxies
high Fe/O ratio in the Magellanic Clouds, and d) the Z-Mass relation present in irregular galaxies (e.g. Aparicio, Garcia-Pelayo & Moles 1988; Russell, Bessell & Dopita 1988a,b; Lequeux 1989; Moles, Aparicio, t Masegosa 1990; De Young & Gallagher 1990; Pilyugin 1993, Tosi 1994). Pilyugin has presented quantitative estimates of this effect on the AY/AZ(O) ratio. By assuming that a fraction 7 of the freshly made Z(O) is ejected to the intergalactic medium without mixing with the interstellar gas, and that the decrease of Mg due to this outflow is negligible, it follows that 2/(0) = (1 - j)y(O)SN + UNWIND .
(2-13)
where y{O)SN is the net yield due to SN explosions and y(O) WIND is the net yield due to the mass loss due to stellar winds before the SN explosion of the massive stars. About half of the freshly made He is ejected by stars in the 8-120 MQ range and about half by stars in the 1-8 MQ range, while almost all the 0 is produced by stars in the 8-120 MQ range and most of it is ejected to the interstellar medium by supernova explosions. Stars with main sequence masses in the 1-8 MQ range end their lives as white dwarfs and eject their excess mass in normal stellar winds in the red giant phase and as superwinds in the preplanetary nebula phase, the expanding shells have typical velocities of a few tens of km s - 1 a n d the material ejected is not expected to escape the parental galaxy. Alternatively, stars in the 8-120 MQ range lose part of their mass in stellar winds and afterwards explode as supernovae. The SN ejecta have typical initial velocities of a few thousands of km s" 1 , and a fraction of the SN ejecta is expected to leave the parental galaxy. Under these considerations it follows that AY 1 Jj,
_ y(Yj- 8 ) + (1 - l)y{y%-i2o)sN + y{y%-\2o)wiND ~ (1 - 7)y(O)SN + y(O)WIND - ^
(2-14)
where b is the fraction of the net He yield due to stars in the 1-8 MQ range and by winds of stars in the 8-120 MQ range; (1 — 6) is the fraction of the net He yield produced by SN explosions due to stars in the 8-120 MQ range; c is the fraction of the net 0 yield due to winds by stars in the 1-120 MQ range, and (1 — c) is the fraction of the 0 yield produced by SN explosions in the 8-120 MQ range. From the computations by Maeder (1992) for Z = 0.001 and the high mass-loss rate case, together with the IMF by Kroupa, Tout, & Gilmore (1993) it follows that 6 = 0.56 and c = 0.0, while for the Salpeter (1955) IMF it follows that 6 = 0.42 and c = 0.0. Consequently for low Z values, the c term can be neglected. Under the assumption that the two types of outflow are present, the [AY/AZ(O)]obs will still be given by equation (2.14) but the net oxygen yield would be given by
[c
+ (i _ c ) ( i _ 7 ) ] l n [ ( A + ! ) „ - ! - A] •
(2 15)
-
3. Solar vicinity 3.1. Observational constraints There are four observational constraints that can be used to study the chemical evolution of the solar vicinity: Z(O), AY/AZ, AY/AZ(O), and Z(O)/[Z - Z(O)]. In this discus-
Peimbert et al.: Chemical Evolution of Galaxies
83
sion we will use for Z(0) the solar value, for AY/AZ the M17 value, for AY/AZ(O) and Z(O)/[Z — Z(O)] an average of the solar value and that of HII regions of the solar neighborhood. The adopted values are: Z(0) = 0.0096 ± 0.0009, AY/AZ = 2.5 ± 0.5, AY/AZ(0) = 5.2 ± 1.4, and Z(O)/[Z - Z(O)] = 0.92 ± 0.10 (see Peimbert et al. 1994 and references therein). The Z(O) value for HII regions depends on their temperature structure; it is often assumed that the mean square temperature fluctuation, i 2 , is zero, and in this case the derived Z(O) is two or three times smaller than ZQ(O). Recent results indicate that t2 RS 0.04 for HII regions of the solar neighborhood (Peimbert, Torres-Peimbert k Ruiz 1992; Peimbert, Storey k Torres-Peimbert 1993a; Peimbert, Torres-Peimbert k Dufour 1993b) which increases the Z(O) of HII regions to values very similar to Z © (0). 3.2. IMF A recent determination of the IMF for the solar neighborhood is that by Kroupa, Tout, k Gilmore (1993), hereafter KTG. The KTG IMF is given by r0.035m-13 £ = I 0.019 m - 2 2 [0.019m-27
if0.08<m<0.5, if 0.5 < m < 1.0, if 1.0 < m < 120,
(3.16)
where m is the stellar mass in solar units and £,{m)dm is the number of stars in the mass interval from m to m + dm. The 0.08 < m < 0.5 exponent for the 95 per cent confidence interval is in the —0.70 to -1.85 range and the center is given by —1.3. We will define three IMFs, KTGa, KTGb and KTGc, that are obtained by extrapolating the KTG IMF to zero mass and adopting the center and the two extremes of the 95% confidence interval; the corresponding IMFs are given by £ a (m) = 0.051m- 07 ,
(3.17)
13
6(m) = 0.035m- , £c(m) = 0.025m-
185
(3.18) ,
(3.19)
for 0 < m < 0.5, while for m > 0.5 we have adopted equation (3.16). The KTG IMFs have the same slope for m > 1 as the Scalo (1986) IMF and the slope only differs for rn < 1 values. 3.3. Stellar models and chemical evolution models of the solar vicinity It is possible from stellar evolution models and a given IMF to compute the following net yields and net yield ratios: y(O), y(Y)/y(O), y(Y)/y(Z), and y(O)/y(Z - O). These values will depend on M(BH), the mass above which the stars end their evolution as black holes without enriching the interstellar medium with heavy elements. Black holes have been postulated by many authors to explain low values of the oxygen yield derived for the Galaxy and for other galaxies (e.g. Twarog k Wheeler 1982, 1987; Mallik k Mallik 1985; Maeder 1992, 1993). For different IMFs y(Y)/y(Z), y(Y)/y(O), and y(O)/y(Z - O) can be directly compared with the observational constraints AY/AZ, AY/AZ(O), and Z(O)/[Z - Z(O)], respectively; this comparison is independent of the infall rate if the infall is due to gas with primordial abundances. Alternatively, j/(O) depends on the specific chemical evolution model adopted for the solar vicinity. From four recent models of the solar vicinity it is found that y(O) - (0.7 ± 0.1) Z(O) (Twarog 1980a,b; Pagel 1989; Meusinger k Stecklum 1992; Malinie et al. 1993).
84
Peimbert et al.: Chemical Evolution of Galaxies
0.008-
A.85
0.006-
• \60 i j \ \ 40
Q.
\ 0.004-
\2 7 -5
O.OOfr Kc
- •
onnoi
8
10
12
14
16
y(Y)/y(O)
FIGURE 1. Oxygen yield, y(O), versus helium to oxygen yield ratio, y(Y)/y(O), based on the KTG IMF and the computations by Maeder (1992). The numbers inside the figure denote M(BH) values. The square is the permitted region for the Z(O) and AV/AZ(O) values for the solar vicinity. The narrow column is the permitted region for the Z(O)/[Z-Z(O)] value. The small rectangle bounded by solid lines is the only permitted region that satisfies the three observational constraints.
Peimbert et al. (1994) have found that the Z(O), AY/AZ, and Z(O)/[Z - Z(O)] values for the solar vicinity presented in § 3.1 can only be fitted by the KTGb IMF together with stellar evolution models by Maeder (1992); this fit is made without the presence of black holes. Furthermore, they also find that, regardless of the value for M(BH), the Salpeter IMF cannot satisfy the three restrictions at the la level. In what follows we will consider AY/AZ(0) instead of AY/AZ; AY/AZ(O) is more convenient because the models by Maeder (1992) do not consider the production of C, N, and Fe by intermediate mass stars, IMS, those in the 1 < M/MQ < 8 range; moreover for these stars, deviations from IRA can occur and should be taken into account. From the Z(O), AY/AZ(0), and Z{O)/[Z - Z{O)} constraints presented in § 3.1, the KTG IMF's and the stellar evolution models by Maeder (1992), we have developed Figure 1. From this figure it is found that at the 1
Peimbert et al.: Chemical Evolution of Galaxies
85
Maeder (1992); she also considers the enrichment of C, 0 , and Z by single IMS based on the computations by Renzini & Voli (1981) and of He based on the computations by Maeder (1992); moreover, she considers the enrichment due to Type la and Ib supernovae produced by binary evolution of IMS. Carigi does not reproduce the observed abundance pattern in the solar vicinity for any value of M(BH) by using the Salpeter IMF. The best agreement by Carigi with the observational constraints is obtained with the Scalo IMF and without the presence of black holes. From the KTGb IMF it is possible to derive a mass density of 8.5 x 10~3M© pc~ 3 in substellar objects by integrating equation (3.18). This value is considerably smaller than the 0.1 MQ pc~ 3 estimate of unobserved mass in the solar vicinity by Bahcall (1984) and Bahcall, Flynn k Gould (1992), and implies that if the missing mass amounts to 0.1 MQ pc~ 3 it is not due to substellar objects and should be of the type that does not participate in the chemical evolution process. 3.4. Conclusions Carigi (1994) and Peimbert et al. (1994) find, from independent considerations, that the production of black holes by massive stars is not needed to explain the chemical evolution of the solar vicinity. This result contradicts previous results by other authors who find that objects more massive than 20-28 MQ end their lives as black holes without enriching the interstellar medium with heavy elements. Carigi (1994) finds that while the Scalo IMF is in agreement with the observational constraints, the Salpeter IMF is not; similarly, Peimbert et al. (1994) find that while the KTG IMF is in agreement with the observational constraints, the Salpeter IMF is not. Peimbert et al. (1994) find, from chemical evolution considerations, that the mass of substellar objects in the solar vicinity is smaller than 0.02 M Q pc~ 3 (2
4. Metal-poor galaxies 4.1. Observational constraints A recent discussion of the problems involved with the chemical evolution of metal-poor galaxies has been presented by Lequeux (1993). One of the most important constraints for the study of the chemical evolution of metal-poor galaxies is provided by the Z-/J, diagram. In Figure 3 we present a Z(O)-fi diagram from the Mg and My data by Staveley-Smith, Davies k Kinman (1992), where Z(O) was derived from the (O/H)-L B relationship presented by Skillman, Kennicutt k Hodge (1989). Lequeux et al. (1979) presented the Z-/J, diagram for a group of irregular and blue compact galaxies. They found that y(Z) — 0.004 together with a closed-box model could explain the observations. Notice that this value of the yield is a factor 3.5 times smaller than that derived from the KTG IMF and the Maeder (1992) stellar evolution models (see below). Peimbert & Serrano (1982) tried to explain the distribution of irregular galaxies, blue compact galaxies, the Galaxy, and M83 by suggesting that the yield increases with metallicity as y(Z) = 0.002 + 0.6Z. All the objects but one considered by them fall in the region limited by the closed-box model and an infall model with a = 1. They gave arguments in favor of a change of the low-mass end of the IMF to explain the increase of y with Z. Matteucci & Chiosi (1983) and Chiosi k Matteucci (1984) suggested that in addition
86
Peimbert et al.: Chemical Evolution of Galaxies
to infall, outflow models of well-mixed material could also explain the position of irregular and blue compact galaxies in the Z-/J, diagram. Kumai & Tosa (1992) suggested that the Z-fi diagram could be explained with a closed-box model and a varying fraction of Mni,/Mb. Another important observational constraint is the AY/AZ. As mentioned in § 2 this ratio is affected by outflow of metal-rich material and is not affected by the gas flow behavior of other models. The presence of black holes could also be important in explaining the Z(O)-(J, diagram because they reduce the y(Z) value. Therefore we will also consider the Z(O)/[Z — £(O)] ratio which is very sensitive to M(BH). In what follows we will discuss the different suggestions that have been proposed to explain the Z(O)-fi diagram. We will consider the Z(O)-fi diagram instead of Z-fi diagram because: a) for most objects we only have a Z(0) value, and b) there are some elements, like C and Fe, for which the IRA is not as good as for 0 . The two main characteristics that need explanation are: a) the large range of /z values for a given Z(0) value, and b) the low (Z(0)) associated to the (/z) value. 4.2. Stellar models and black holes
For our discussion we will make use of the Maeder (1992) stellar evolution computations for Z — 0.001 and Z = 0.02 with the high mass-loss rate. From these models the Z production considered is due to stars with M > 8 MQ for which the IRA is excellent, moreover the Z(O)/Z ratio is a strong function of the stellar mass. In Figure 2 we present the y(O)/y(Z — O) ratio based on Maeder's computations and the KTG and Salpeter IMFs, notice that this ratio is independent of the behavior of the IMF for M < 8 M©; this ratio is compared with the observed Z(O)/[Z - Z(O)] value determined for the best observed metal poor HII region, NGC 2363; from this HII region it is obtained that Z(O)/[Z - Z(O)] = 1.48 ± 0.15 (Peimbert, Pena & Torres-Peimbert 1986). This observational restriction is presented in Figure 2. From this figure it follows that there is an excellent agreement between the KTG IMF and Maeder's model without the need to invoke black holes. Alternatively, with the Salpeter IMF, M(BH) = (83 ± 15) M 0 ; by considering the production of C, Fe, and N by stars with M < 8 MQ, the y(O)/y(Z — O) ratio becomes even smaller, strengthening the result that black holes do not play a role in the chemical evolution of metal-poor galaxies. 4.3. Closed-box model
From the computations by Maeder (1992) and the KTG IMF we find that y(Z) = 0.014 and 2/(0) = 0.0083 for Z = 0.001. In Figure 3a-3b we present the closed-box solution for y(O) = 0.0085. It is clear that a closed-box solution, under the assumption that Mnb
Peimbert et al.: Chemical Evolution of Galaxies
20
40
60
80
100
87
120
FIGURE 2. y(O)/y(Z - O) versus M(BH) (in solar mass units) diagram for the IMFs by KTG and Salpeter based on the computations by Maeder (1992). The yields ratio depends only on the IMF for m > 1, consequently there are no differences among cases a, b, and c. The permitted band corresponds to the Z(O)/[Z-Z(O)\ observed value of 1.48 ± 0.15.
and He abundances might be affected by the presence of Wolf Rayet stars (Pagel et al. 1992), but the O/H ratios derived from different HII regions in the SMC and the LMC (e.g. Dufour 1984 and references therein) and for different HII regions in I Zw 18 (Skillman & Kennicutt 1993) indicate that the self-enrichment in O is not very important. Furthermore, the presence of O self-enrichment makes it more difficult to explain the Z(O)-fx diagram, since we need to increase 2/(0) not to lower it. The problems associated with the determination of Mr could be greater. A line profile obtained with a single radio dish cannot be converted into a reliable total mass estimate because the inclination and radius derived from optical images can be misleading (Skillman 1992); moreover, the assumption that the velocity field is due to rotation might not be correct in some cases (e.g. Skillman k Kennicutt 1993), nor the assumption that rotation is completely absent in others. Reliable mass estimates must be based on interferometric observations and only a few objects have reliable total mass determinations. 4.5. Variations in the IMF To explain the Z(O)-(i diagram for low Z(O) values, smaller values of y(O) are needed. This can be done by reducing the fraction of stars in the high-mass end of the IMF (M > 8 MQ) or by increasing the fraction of objects in the low-mass end of the IMF (M < 1 — 2 M Q ), including substellar objects. Massey et al. (1989a,b), and Parker et al. (1992) find that most of the OB associa-
88
Peimbert et al.: Chemical Evolution of Galaxies
tions in the Magellanic Clouds have IMFs with slopes in the —1.6 < T < —1.9 range. Parker k Garmany (1993) find T = - 1 . 5 ± 0.2 for the 30 Doradus region. These results are in agreement with the KTG slope of - 1 . 7 and do not present evidence for substantial changes in the high-mass end of the IMF with metallicity. Moreover, from the study of many extragalactic HII regions in irregular and blue compact galaxies of different Z(O) content, McGaugh (1991) finds that a fixed IMF adequately explains the observed trends in the spectra, so no variation in the high-mass end of the IMF is inferred. Notice that a Salpeter slope for the high-mass end of the IMF predicts a higher heavy element yield, making it more difficult to explain the Z(O)-fi diagram. There is some evidence in favor of an increase in the fraction of low-mass objects with decreasing Z (e.g. Peimbert k Serrano 1982). Furthermore, Elmegreen (1994) has suggested that smaller-mass galaxies could have a lower low-mass end of their IMF due to their smaller velocity dispersion. In Figure 3c-3d we present a series of models with different outflow rates and with 2/(0) = 0.0012 + 0.67Z, or y(Z) = 0.002 + L I Z . This type of yield reduces the need for strong effects due to other mechanisms to explain the Z(O)-fi diagram.
4.6. AY/AZ(0) The best AY/AZ(0) ratio for irregular and blue compact galaxies is that by Pagel et al. (1992) and is equal to 10.3 ± 3.5, this value was derived without considering the fraction of O tied down in dust grains or the presence of temperature variations inside the observed HII regions. From the results by Skillmanet al. (1994) it is found that about 15 % of O is tied down in dust grains, therefore the AY/AZ(O) by Pagel et al. (1992) should be lowered to 8.8. By considering a modest amount of correction due to t2 the value should be lowered to about 6.7. This value is a factor of 2.4 higher than that computed from the KTG IMF and the stellar evolution models by Maeder (1992) and probably implies the presence of Z-rich outflows. This value is also a factor of 3.3 higher than that computed from the Salpeter IMF. Unfortunately not all the objects in the sample by Pagel et al. (1992) have Mr and Mg determinations. On the other hand, not all the objects in the sample by Matteucci k Chiosi (1983) and by Staveley-Smith et al. (1992) have Y and Z determinations. Colin et al. (1994) have used a subsample of the sample by Pagel et al. (1992) with 10 objects and find AY/AZ(0) = 10.3 ± 1.8 without considering the amount of O tied down in dust grains, or the effect off2 in the O determinations. The (Z(O)} and (log/z) values for the Colin et al. (1994) sample are 0.0016 and —0.52, respectively. 4.7. Inflow Inflow of material with pregalactic abundances, Y = Yp and Z — 0, does not change AY/AZ(O), though it is important in reducing Z for values in the 0 < a < 1.0 range. For larger values of a there is a minimum fi value, see Figures 3a-3b and equation (2.10), and this process is not able to cover the regions of low Z(O) and low fi values. The solutions presented by Matteucci k Chiosi (1983) that cover all the Z-\i diagram for a > 1.0 are incorrect because equation (2.10) does not apply for fi < (a — 1)/Q. 4.8. Outflow of well-mixed material It is possible to cover the Z-fi diagram, to the right of the solution for the closed-box model, by varying A. There are two problems with this solution: a) it does not explain the large AY/AZ(O) observed values, and b) for some galaxies, with low Z(O) and low
89
Peimbert et al.: Chemical Evolution of Galaxies
I/ 1 ; 1 ' ' I ' ' ' '
1 - LSB / :
"
^4i
0
i
i
(a) 4 - BCD/.'
i
-0.5
i
•
,-"• i
i
i
-1
i
(b)
i
i
-1.5
0
-0.5
-1 log
-1.5
FIGURE 3. Z(O)-fi diagram for the sample by Staveley-Smith et al. (1992), where BCD and LSB denote blue compact and low surface brightness galaxies, repectively. Figures a-b present a closed-box model for y(O) — 0.0085 and infall models for different values of the infall parameter a (see eq. 2.10). Figures c-d present well-mixed outflow models for a yield increasing with metallicity given by J/(O) = 0.0012 + 0.67 Z(O) and for different values of the outflow parameter A (see eq. 2.12); the 0.0 curve corresponds to the closed-box model.
90
Peimbert et al.: Chemical Evolution of Galaxies
fi, values of A of about 50 are needed, which would imply the presence of haloes with about fifty times more mass in gas than the amount of mass in the form of stars in the main body of the galaxy, we consider this possibility unlikely. This type of outflow does make the prediction that galaxies should be surrounded by massive gas haloes, a prediction that could be tested. In Figures 3c-3d we present models with a yield increasing with metallicity and various values of A. 4.9. Outflow of Z-rich material
To explain an increase in the AY/AZ(O) ratio due to the outflow of Z(O)rich material we can make use of equation (2.14) and the stellar evolution models by Maeder (1992). For an increase of a factor of 2 in the AY/AZ(0) ratio, for an object with Z = 0.001, a value of 7 = 0.7 is needed for the Salpeter IMF and a value of 7 = 0.64 for the KTG IMF, this would raise the yields derived from Z-fi diagram by a factor 1/(1 — 7) which is of the order of three, reducing the importance of the other suggested mechanisms to explain the Z-fi diagram. Moreover, the objects with small \i and small Z might yield even higher AY/AZ(O) ratios, and this process by itself could explain all the points in the Z-fi diagram. In particular, for the sample of Colin et al. (1994), (Z) = 0.0027, (Z(O)) = 0.0016 and (log^i) = —0.52, which, together with equation (2.8), give y(O)ej = 0.0013, and y(Z)ej = 0.00225. Based on the computations by Maeder (1992) interpolated for Z = 0.0027 and adopting the KTG IMF, we obtain b = 0.58 and c = 0.021; under the assumption that AY/AZ(O) has been increased by a factor of 2.4 due to Z-rich outflow we obtain 7 = 0.73, which, together with equation (2.15 for A = 0), give y(O) = 0.0057 and y(Z) = 0.01, still a factor of 1.4 smaller than the predicted yields from Maeder and the KTG IMF that amount to 0.0083 and 0.014, respectively. This result reduces considerably the importance of other mechanisms under consideration, nevertheless the errors in the AY/AZ(O) determination are large and better determinations are needed to evaluate the importance of this effect. 4.10. Dark matter that does not participate in the chemical evolution process
Kumai & Tosa (1992) have suggested that the Z-fi diagram can be explained with a closed-box model with a constant yield combined with Mnj 7^ 0, where each object would have a different Mnt/Mi ratio. Based on the sample by Matteucci k Chiosi (1983), Kumai & Tosa find a good correlation between the dark matter fraction, fo, derived by two independent methods. The methods to determine fo are based on: a) fitting the Z-fi diagram with a constant yield, and b) determining the difference between the dynamical mass and the sum of the gaseous mass plus the stellar mass derived from an adopted M«/L ratio. Of the effects that affect the Z-fi diagram discussed before there is one that is particularly relevant to this suggestion: errors in the determination of Mr could be responsible for part of the correlation derived by Kumai & Tosa (1992) because they will affect in the same manner the fo values derived by the two methods. Nevertheless, from the sample of Matteucci k Chiosi (1983) there is no correlation between M n j/Mj (or fa) and the mass of the galaxy. We would expect fD to increase for less massive objects (e.g. White L Frenk 1991; Persic &; Salucci 1988; Persic, Salucci & Ashman 1993). Persic et al. find for a sample of 56 galaxies, most of them spiral, that the fractional visible mass at the optical radius, Ropt, decreases for less luminous galaxies. From the sample by Staveley-Smith et al. (1992), Colin et al. (1994) find that there is an inverse correlation of Mni/Mb with MT for the subsample made up of the low surface
Peimbert et al.: Chemical Evolution of Galaxies
91
brightness galaxies of the sample, but there is no correlation between Mnb/Mb and for the subsample of blue compact galaxies. This result might imply that Mnj plays a role in the explanation of the Z-fi diagram for the low surface brightness galaxies but not for the blue compact ones. 4.11. Conclusions The closed-box model with the KTG or the Salpeter IMF and the stellar evolution models by Maeder (1992), under the assumption that Mnf, — 0, cannot explain all the points in the Z(O)-fi diagram. The distribution of points in the Z(O)-fi diagram is not mainly due to: a) errors in the Z(O) determinations, b) the presence of black holes, c) inflow, and d) a varying IMF at the high-mass end. Alternatively, the distribution of points in the Z(O)-/i diagram could be due to: a) errors in My, b) a varying IMF at the low-mass end, c) outflow of well-mixed material, d) the presence of dark matter that does not participate in the chemical evolution process, and e) outflow of Z(O) rich material. The first four possibilities cannot explain the high AY/AZ(O) observed values and one or more of them should be combined with the last to reproduce the Z(O)-fi diagram. Alternatively, if the AY/AZ(O) observed value is considerably higher than the value discussed above, then the derived 7 will become even higher and the last possibility alone might be responsible for the distribution of points in the Z(O)-fi diagram. Determinations of Y, Z(O), Mg, and MT of high qualtity are needed to advance in this problem, particularly of objects with low Z(O) and low fi values.
Manuel Peimbert wishes to acknowledge helpful discussions with many of the conference participants, in particular with Evan D. Skillman, moreover he also wishes to thank the members of the Instituto de Astrofisica de Canarias for the invitation to visit the Institute and participate in this conference and for their warm hospitality. The authors acknowledge Luis A. Martinez for his assistance with the handling of the postscript figures.
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Galaxy Properties in Different Environments: Star Formation in Bulges of Late-Type Spirals 1 2 By M. G. PASTORIZA !, E. BICA 1 M. MAIA 1 1 C. BONATTO AND H. DOTTORI 1
Departamento de Astronomia, IF-UFRGS, CP 15051, CEP 91501-970,Porto Alegre, RS, Brazil 2
Departamento de Astronomia, Observatorio Nacional, Rua Gal. Jose Cristino 77, Rio de Janeiro, 20921-RJ, Brazil
The star formation history in the nuclei of late-type spiral galaxies is compared between a sample in a high galaxy density medium (HDS) and a control sample (CS) of isolated galaxies. We have observed 20 HDS and 18 CS galaxies selected from a larger list generated by the application of a group-finding algorithm to the SSRS survey. Using equivalent widths of absorption lines and the continuum distribution, we determined the nuclear stellar population types, from those dominated by old populations to those containing star formation bursts of different ages and intensities. The HDS and CS stellar population type histograms are similar, suggesting that environmental influences, at least for the present samples, do not substantially affect the nuclear stellar population. However, the nuclear emission lines indicate that, in the BPT diagnostic diagrams, there is an excess of HDS galaxies located within or close to the AGN loci. For 6 HDS and 2 CS galaxies it was possible to determine Oxygen (O/H) and Nitrogen (N/H) abundances. The samples present similar (O/H) values, but in the CS galaxies the (N/O) ratio is lower at equal galaxy luminosity.
1. Introduction Evidence of several environmental effects that affect galaxy properties have been reported recently: (a) the morphology-density relation (fractional increase of early-type galaxies towards regions of high concentration, Dressier 1980; Postman & Geller 1984; Giovanelli, Haynes & Chincarini 1986; Maia & da Costa 1990); (b) the morphologyclustercentric radius relation (Whitmore, Gilmore & Jones 1993). However, the impact of the local density of galaxies on the star formation rate (SFR) is not well understood. Two effects have been suggested to influence the SFR: (1) the SFR could be enhanced by tidal interaction (Bushouse 1986; Kennicutt et al. 1987); (2) in the very high density regions in the cores of clusters, close encounters leave preferentially anemic spirals (Dressier 1984). The W(Ha) of galaxies within groups tends to be larger than those outside groups (Maia & da Costa 1990). In order to understand the dependence of the SFR on the galaxy environment, we have studied the stellar component and properties of the emitting gas for two galaxy samples: High Density Sample (HDS): p > 18galMpc- 3 , and Control Sample (CS): p < 0.0004gal Mpc~ 3 . t Visiting Astronomer at the Cerro Tololo Interamerican Observatory, operated by the Association of Universities for Research in Astronomy, Inc. under contract with the National Science Foundation. 94
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2. The samples The southern Sky Redshift Survey Catalog (da Costa et al. 1988, 1989, 1991) was taken as data base for the sample selection. This SSRS catalog consists of 2028 galaxies selected from the ESO/Uppsala survey, satisfying the conditions: log[D0] > 0.1, 6" < -30°, 6 < -17.5°. The High Density Sample is formed by galaxies that are in groups of 3 or more members generated by the group finding algorithm described by Maia, da Costa & Latham (1989). The groups have a surrounding density contrast (8p/p) > 500, equivalent to p > 18 gal Mpc~3. Groups are formed by accumulation of galaxy pairs with a member in common having a projected separation D12'. D12 = 2{V/H0)am(912/2)(Vy^3, Vo = 400 km s" 1 , and (V) is the selection function for the galaxy Catalog normalized to VF = 1000 km s" 1 . The Control Sample is made up of galaxies that are not assigned to any group in the search procedure for neighbors. This search was performed at (Sp/p) = 0.01, equivalent to p< 0.0004 gal Mpc- 3 . All the selected objects have radial velocity (after correction for the Virgo infall) VR < 8000 km s" 1 . HDS contains 151 galaxies, and the CS contains 179 galaxies. 3. Subsamples of late-type spirals (Sb to Sc) We have selected HDS and CS late-type spirals brighter than m? = 14 mag, and avoided edge-on galaxies, since our spectral analysis deals with the central region within 1.2 kpc. The HDS late spiral sample contains 50 galaxies and the CS, 79. 4. The observations Observations were carried out at the 1.0-m telescope of CTIO (Chile), with the Cassegrain Spectrograph and 2D-Frutti detector; spectral range: AA3600 - 7000 A, resolution 5 A pixel"1 at A5000 A. Reductions were made with IRAF at our University. 5. Absorption line spectra The stellar population analysis is based on the equivalent width of the absorption features. Continuum tracing and spectral windows are those defined by Bica & Alloin (1986). A typical error is a w 0.5 A. The equivalent widths were used to classify the galaxies in terms of templates which span the same properties of the stellar population in normal galaxies. Characteristics of the templates: SI to S3 — red population, metallicity decreasing from 4 to 1 Z©; S4 to S7 — increasing contribution from blue components, younger than lGyr: from 10% to 70% of flux at A5580 A. Most of the HDS and CS galaxies match these templates (see example in Figure 1). No substantial differences are found between the stellar population types for the two samples as shown in Figure 2.
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U
4000
5000
6000
MA) FIGURE
1. Top panel: example of a red stellar population; bottom panel: blue stellar population.
6. Emission lines and nuclear activity We have subtracted the stellar population contribution from the calibrated spectra normalized at A5870 A. We found all types of nuclear activity among the galaxies of the samples. The diagnostic diagrams using emission-line ratios show that most of the HDS galaxies are located in the region occupied by active galaxies or transition between HII regions and AGNs, while the CS galaxies have HII region-like spectra.
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1O HDS
6
--
4
-
2
-
8
-
N
cs
6 N 4
-
2
-
-
S1
S2 S3 E7 S4 S5 S6 Stellar Population Type
S7
FIGURE 2. The distribution of stellar population does not show any remarkable differences between CS and HDS galaxies.
7. Oxygen and nitrogen abundances Evidence that spiral galaxies in the Virgo Cluster have systematically higher interestellar abundance than the field ones have been presented by Shields, Skillman & Kennicutt (1991). However, Henry et al. (1992) found that Virgo spiral and field galaxies are indistinguishable in terms of abundance properties. In Figure 3 are plotted log(N/O) and 12 + log(O/H) vs. absolute magnitude. Wefindthat independently of galaxy luminosity, HDS and CS galaxies present similar values of (0/H), and, for equal luminosity, CS galaxies present a tendency of lower (N/0) than HDS galaxies. These findings suggest that the environment does not affect (0/H) but may have some influence on (N/O).
8. Conclusions The main conclusions of this paper are: (a) the environment does not play a significant role in the stellar population type of the bulges of the late-type spirals; (b) Seyfert-like nuclear activity is more frequently found in HDS galaxies; (c) the environment does not have much effect on the oxygen abundance, but may have some influence on the nitrogen for HDS galaxies.
98
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Environments
0.0
ocs
r
-0.5
_.
-1.0
HDS
N4303
p CD
-1.5 A N3310
A LMC
-2.0
-2.5
N43O3
9.5
A A MS1
9.O
N3310A LMC
o
IS
AM81
£ 8.5 8.0
7.5
-22
-20
-18
-16
FIGURE 3. Chemical abundance for the 2 samples plotted against Absolute Magnitude.
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DRESSLER, A. 1980 Astiophys. J. 236, 351. A. 1984 Astion. Astrophys. 22, 185. GIOVANELLI, R., HAYNES, M. P. & CHINCARINI, G. L. 1986 Astiophys. J. 300, 77. DRESSLER,
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Star Formation in Galaxies in the Bootes Void ByDONNA WEISTROP Department of Physics, University of Nevada, Las Vegas, NV 89154, USA Observations are reported for 16 of the 25 galaxies found in the Bootes void. At least five of the galaxies are spirals, and another five are disk systems. Two interacting galaxy pairs have been definitely identified, and there are additional candidate pairs. Several of the galaxies are luminous in Ha, due in most cases to significant amounts of star formation. The observed galaxies do not resemble the galaxies predicted to inhabit voids. There is evidence for structure in the spatial distribution of the galaxies, in particular half the galaxies are located in a single plane 10 Mpc wide.
1. Introduction The Bootes void, a spherical region with radius 62 Mpc at a distance of 310 Mpc (Ho = 50 km s" 1 Mpc" 1 ), is one of the largest known low-density regions in the largescale distribution of galaxies (Kirshner et al. 1981