Dwarf Galaxies

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,::;,q'-:;' ..-",("• .:: 114 ML l •I ~ -." "." I_I The First ESO/ESA Workshop on the Need for Coordinated Space and Ground-based Observations -

DWARF GALAXIES

Geneva, 12-13 May 1980

Report Edited by M. Tarenghi and K. Kjar - iii -

INTRODUCTION

The Space Telescope as a joint undertaking between NASA and ESA will provide the European community of astronomers with the opportunity to be active partners in a venture that, properly planned and performed, will mean a great leap forward in the science of astronomy and cosmology in our understanding of the universe. The European share, however,.of at least 15% of the observing time with this instrumentation, if spread over all the European astrono mers, does not give a large amount of observing time to each individual scientist. Also, only well-planned co ordinated ground-based observations can guarantee success in interpreting the data and, indeed, in obtaining observ ing time on the Space Telescope. For these reasons, care ful planning and cooperation between different European groups in preparing Space Telescope observing proposals would be very essential.

For these reasons, ESO and ESA have initiated a series of workshops on "The Need for Coordinated Space and Ground based Observations", each of which will be centred on a specific subject. The present workshop is the first in this series and the subject we have chosen is "Dwarf Galaxies". It was our belief that the dwarf galaxies would be objects eminently suited for exploration with the Space Telescope, and I think this is amply confirmed in these proceedings of the workshop.

It is a pleasure for me to thank all the chairmen, who worked hard to prepare their sessions, and all the con tributing speakers. Per Olof Lindblad - v -

EDITORS' NOTE

Our appronctJ to the publication of this report has been similar tu that followed for previous ESO workshop and conference proceedings: the speakers were a;;ked to submit their manuscripts, which were then used directly for off set printjnff' We are grateful to the speakers for pro viding us with most of the papers shortly after the end of the workshop.

The discussions after the papers are only printed when they were confirmerl in written form.

We should like to express our grateful thanks to all those who helped us in the organization of this conference, in particular Renate van Doesburg.

M. Tarenghi K. Kjar - vii -

LIST OF PARTICIPANTS

VAN AGT, S.L. Sterrenkundig Institut Toernooiveld NL-Nijmegen

ALLOIN, D. ESO

ARDEBERG, A. ESO

ARP, H.C. ESO

AUDOUZE, J. Institut d'Astrophysique du CNRS 98 bis, boulevard Arago F-75014 Paris

BARBIERI, C. Istituto di Astronomia Vicolo Osservatorio 5 I-25100 Padova BERGERON, J. ESO

BERTOLA, F. Osservatorio Astronomico Vicolo Osservatorio I-35100 Padova

CASTELLANI, V. Istituto di Astrofisica Spaziale C. P. 67 I-00044 Frascati

CRANE, P. ESO

DANZIGER, J. ESO

DENNEFELD, M. ESO

DISNEY, M. Dept of Applied Mathematics and Astronomy University Colleget Cardiff P.O. Box 78 GB-Cardiff CF1 1XL

D'ODORICO, S. ESO

FEITZINGER, J. Astronomisches Institut Ruhr-Universitat Bochum Postfach 102148 D-4630 Bochum 1

GROSB~L, P. ESO - viii -

HARTQUIST, T. University College London Dept of Physics and Astronomy Gower Street GB-London WCIE 6BT

HUCHTMEIER, W.K. Max-Planck-Institut fur Radio astronomie Auf dem Hugel 69 D-5JOO Bonn

JONES, B. University of Cambridge Institute of Astronomy Madingley Road GB-Cambridge, CBJ OHA

KUNTH, D. ESO

LAUSTSEN, S. Institute of Astronomy University of Aarhus DK-8000 Aarhus C

LEQUEUX, J. Observatoire de Paris Section d'Astrophysique F-92190 Meudon

LINDBLAD, P.O. ESO

MACCHETTO, F. ESA/ESTEC Astronomy Division Domeinweg NL-Noordwijk

MATERNE, J. Institut fur Astronomie und Astrophysik der TU Berlin Ernst-Reuter-Platz 7 D-l000 B~rlin 10

MELNICK, J. ESO

NORMAN, C. Sterrewncht Leiden TIlly{~cns La-boratoriurn Wns~cnRnrsewee 7H NL-2100 RA Leiden

PERRYMAN, M. ESA/ESTEC Astronomy Division Domeinweg NL-Noordwijk

SA-NCISI, R. K~pteyn Astronomical Laboratory Hoogbouw W.S.N. Postbus 800 NL-Groningen - ix -

SELLWOOD, J. ESO

DI SEREGO, S. ESA/ESTEC Astronomy Division Domeinweg NL-Noordwijk

8HAW, R. ESO

TAMMA~N, G. ESO

TARENGHI, M. ESO

TENORIO-TAGLE, G. ESO

TERLEVICH, R. University of Cambridge Institute of Astronomy Madingley Road GB-Cambridge CB) ORA

ULRICH, M.H. ESO

VIALLEFOND, F. KapteYn Astronomical Laboratory Hoogbouw W.S.N. Postbus 800 NL-9700 AV Groningen

VAN WOERDEN, H. KapteYn Astronomical Laboratory Hoogbouw W.S.N. Postbus 800 NL-9700 AV Groningen - xi

CONTENTS

Introduction "...... iii

Editors r Note ...... •• v

List of Participants ...... ••...... •..•••• vii

Contents xi

SURVEY OF DWARF GALAXIES (Session I; Chairman: G. Tammann)

G.A. Tammann: A Survey of Dwarf Galaxies 3

J.V. Feitzinger: Galaxies of Magellanic Type: Preli- minary Report •••••••••••••.•••••••• 19

H. van Woerden: A Survey of Galaxies w~th Velocities Smaller than 3000 km s-1 •••••••••• 23

VARIABLK STARS IN DWARF GALAXIES (Session II; Chairman: S. Laustsen)

S. van Agt: Variable Stars in Spheroidal Galaxies.. 33

S. Laustsen: Variable Stars in Dwarf Im Galaxies ••• 39

Im GALAXIES (Session III; Chairman: J. Lequeux)

(a) Frequency Determination, Galaxy Luminosity Function

G.A. Tammann: The Luminosity Function of Dwarf Galaxies 45

S. Laustsen and P. Gammelgaard: A Cloud of Dwarf Irregular Galaxies? ••••••••••••••• 49

(b) Hydrogen Content

W.K. Huchtmeier: HI Content of Dwarf Galaxies •••••• 57 - xii -

W.K. Huchtmeier: HI Observations of Irregular Galaxies 6S

J. Materne: Envelopes of Late-type Dwarf Galaxies. 67

J. Lequeux: Introductory Report ...... 79 J.V. Feitzinger: Studies of the Large Magellanic Cloud 83

A. Ardeberg, P. Linde, H. Lindgren and G. LyngR: Colour-Magnitude Diagrams and Luminosity Functions in the Magel- lanic Clouds 87

INTERGALACTIC HII REGIONS (Session IV; Chairman: D. Kunth)

D. Kunth: Intergalactic HII Regions 95 S. D'Odorico, P. Patriarchi and M. Perinotto: Astrophysics of Galactic -and Extr~ galactic HII Regions with IUE: A Preliminary Discussion of the Car bon Abundance and the Absorption Line Spectrum •..••••.•.....•••... 103

C. Barbieri and D. Kunth: IUE Observations of Blue Compact Galaxies ••...•••....•.... 113

G. Statinska and D. Kunth: The Helium Abundance De termination in Emission-line Dwarf Galaxies 117

J. Lequeux and F. Viallefond: HI Observations and Star Formation in the Blue Compact Galaxy I Zw 18 •....••..••...•.•.• 121

G. Tenorio-Tagle: Isolated Extragalactic HII Re- gions. "The Triggering Mechanism." 123

J. Audouze: Theoretical Implications of a Low Helium Abundance in Dwarf Galaxies 129 - xiii -

ELLIPTICAL DWARF GALAXIES (Ses~ion V; Chairman: L. Woltjer)

(a) Chemical Abundances, Main Sequence Observations

J. Danziger: Abundance Problems in Dwarf Spheroidals and Dwarf Ellipticals of the Local Group...... 133

(b) Globular Clusters in Dwarf Galaxies

V. Castellani: Globular Clusters in Dwarf Galaxies. 141

M.J. Disney: Are Some Dwarf Galaxies Crouching Giant s ? ...... •.. 1 51

F. Bertola: Energy Distribution in Dwarf and Giant Elliptical Galaxies ••.••..•••.•••• 161

MORPHOLOGY AND EVOLUTION OF DWARF GALAXIES (Session VI; Chairman: B.J.T. Jones)

B.J.T. Jones: The Origin of Dwarf Galaxies •...... 167

C. Norman and B.J.T. Jones: Gas Rich Dwarfs •....•.• 173

T.W. Hartquist: Galaxy Formation by the Agglomera- tion of Primordial Star Clusters 175

J.V. Feitzinger: Stochastic Selfpropagating Star Formation in.Magellanic Type Galaxies...... 181

S. di Serego Alighieri: Summary of ST Capabilities on Dwarf Galaxies •..•...... •••• 183 - 1 -

SESSION I

(Chairman: G. Tammann)

SURVEY OF DWARF GALAXIES 3

A SURVEY OF DWARF GALAXIES

G. A. Tammann

ESO; Astronomical Institute, University of Basel.

Can the Space Telescope be used for a systematic search

for dwarf galaxies? To answer this simple question, we

should specify what we mean by dwarf galaxies, and consi-

der some of their typical properties (or what we believe to

be typical).

I. Definition

No unique distinction can be made between "normal" and

"dwarf" galaxies on physical grounds. That is not to say

that the possibility of a physical separation can presently be excluded. It could be, for instance, that dwarf and normal ellipticals obey separate luminosity functions, or

that the space distribution of dwarfs is clumpier (or

smoother??) than that of normal galaxies.

All we can do for the moment is to propose a working de:finition; "Te will there:fore ca]l p,-a.laxles brighter than M B = m -16 "normal", and those fainter, "dwarfs." This definition excludes SMC (-17~0), but includes NGC 6822 (-15~3) and

NGC 205 (-15~7) as dwarfs. The unfortunate, but unavoidable

feature of this definition is that it depends on the adopted distance scale. For some Local Group (LG) dwarfs, the individual distances are still poorly known. For more distant dwarfs the luminosity depends on the preferred value 4 of the Hubble constant, which is taken here to be H = o 50 km s-l MPC-1 •

m The one rational vindication of the limit of ~ = -16 may be that the luminosity function of dwarf irregulars seems to peak near this absolute magnitude (see Session Ill).

For dwarf ellipticals, however, no turnover of the luminosity function is observed, and their number -- at least in the Virgo cluster -- increases with decreasing luminosity.

It would then seem that dwarf galaxies are more numerous than normal ones. That would bring astronomy hazardously close to psychology, where, too, the standard of

"normal" seems frequently to be exemplified by only a minority. However, as Reaves (1967) has pointed out, the contribution of all dwarf galaxies to the total mass in the universe is dwarfish. Masses of galaxies being today notoriously uncertain, one would like to replace the parameter mass by luminosity. Dwarf galaxies would then simply be that subset of all fainter galaxies, whose total contribution to the luminosity density (within a reasonably large volume) is negligible. Indeed, even if one adopts a luminosity function of the form log N(M) oc 0.2 M (Zwicky, 1957), which is the steepest increase so far proposed, the m galaxies fainter than ~ = -16 contribute less than two percent to the total light. 5

11. Classification

Three obvious classes of dwarf galaxies exist:

(1) Dwarf ellipticals (dE), first recognized as a sepa rate class of dwarfs by Baade (1944) and Wilson (1955), are characterized by the absence of bright stars (M~ ~ -1~5) , near absence of neutral hydrogen (Huchtmeier, 1980), and central symmetry of the (elliptical) isophotes (Hodge, 1971).

It has been proposed that the dEs with the faintest

surface brightness, like the Draco system, form a separate

class, the dwarf spheroidals (van den Bergh, 1959). However,

the dEs in the Virgo cluster exhibit a continuous transition

from high to very low surface brightness, such as to make

one question the necessity or even usefulness of the dis

tinction between elliptical and spheroidal dwarfs. From an

observational point of view, one might wish to preserve the

spheroidal class for those dwarfs whose surface brightness

lies below the detection limit, and which can be detected

only because of the concentration of resolved stars. Yet it

is this detection limit which shall be changed by the Space

Telescope.

(2) Dwarf irregulars, dIm (m stands for Magellanic,

although the dwarfs are considerably fainter than the Clouds),

are characterized by bright, blue stars, considerable amounts

of neutral hydrogen, and in many cases, by H 11 regions. 6

A bifurcation of the hydrogen-rich dwarf class has been proposed (van den Bergh, 1959), viz. dwarf irregulars and dwarf spirals. It must be noted, however, that with our present definition of dwarfs, it is doubtful whether any dwarf exhibits a clear spiral structure (type Sd or earlierL

At least 95% of all hydrogen-rich dwarfs are of type Srn or

Im as classified by several authors. It seems therefore m that galaxies with M < -16 are generally denied a spiral B structure.

(3) The "intergalactic H 11 regions", some of which m have M > -16 , are subject of a later session and shall not B be considered here.

111. Census

A dwarf with M > -16m and a redshift of 500 km s-l has B an apparent magnitude of m > 14rn • Therefore the great B majority of known (non-cluster) dwarf galaxies are expected to be contained in a catalogue of galaxies with v o -1 500 km s Indeed, such a catalogue (Kraan-Korteweg and

Tamrnann, 1979) lists 13 dEs and 76 dIms. Since the comple- tion of this catalogue, very few nearby dwarf galaxies have become known. In the following we shall refer to the dwarfs in the catalogue as the "10 Mpc sample".

Additional dwarf galaxies are known in the Fornax cluster (Hodge 1959, 1960) and in the Virgo cluster (Reaves,

1956). Recent efforts have increased the number of Virgo 7 cluster dwarfs to well over 1000 (Reaves, 19S0; Sandage and collaborators, 19S0).

The absolute-magnitude distributions of the dEs and dIms of the 10 Mpc sample are shown in Fig. 1. It is strik ing that the 13 dE's, with _Srn> ME > -16m, are companions of only three galaxies (the Galaxy, M31, and MS1). Another well known example of the clustering of dEs about a giant galaxy is provided by M100 in the Virgo cluster (Sandage,

1975). The present observational evidence does not exclude the conclusion that dEs are formed only in the neighbourhood of large (spiral) galaxies. The more regular distribution of the majority of the Virgo cluster dEs does not necessarily contradict this conclusion, because they may have been smoothed out by the cluster dynamics.

In Fig. 1 the dIms are separated into those which occur in groups of galaxies (Local Group, MS1 group, and a few other groups; cf. Kraan-Korteweg and Tammann, 1979). dIms in groups can be as faint as M ~ -S~9, whereas the faintest B m field dIm (D 8833) has M ~'-12.3. However, this difference B is not significant, because groups (including ~he Local

Group!) have been more thoroughly searched for dwarfs.

Contrary to dEs, there is no indication at all that dIms would favour the neighbourhood of large galaxies; they do occur in the field ("field" meaning here the halo of the

Virgo complex), and the field dIms may well obey the same luminosity function as their counterparts in groups. 8

N dE

M31 M31 M31

Gal Gal M31 Gal M31 N in groups dIm m o field 10 B 6

4 '-" 2 0 -8 -9 -10 -11 -12 -13 -14 -15 -16 Mqi B Fig. 1. The absolute magnitude distribution of the dEs {upper panel} and dIms {lower panel} of the 10 Mpc sample. For dEs it is indicated whose companions they are. dIms are differen tiated into group members and field dwarfs. The sample is strongly affected by selection effects; statistical properties must be derived with prudence.

IV. Detection Limits

Magnitude-limited catalogues of galaxies contain only a minute fraction of dwarf galaxies. This is for two reasons:

{I} To a given magnitude limit the volume sampled for underluminous galaxies is very small, and

{2} The low surface brightness of dwarfs discriminates against their inclusion into catalogues. The Shapley-Ames 9 catalogue, attempting completeness to m ~ 13m, contains from a total of 1246 galaxies only 7 dwarfs (Sandage and Tammann,

1980), whereas 7 dwarfs brighter than 12m (!) are missing from the catalogue. Also the Zwicky ca~alogue is incomplete for low-surface brightness galaxies.

From this it is clear that dwarf galaxies are badly underrepresented in existing galaxy catalogues. Although they may dominate the universe numerically, the vast majority remains anonymous. Special search programs for dwarfs are therefore of great interest. What can be done with the Space Telescope?

The dwarfs of the 10 Mpc sample are plotted in Fig. 2, which is an apparent magnitude-log diameter diagram. The apparent magnitudes, ~ , are taken from a compilation of the T available data (Kraan-Korteweg and Tammann, 1979); the diameters are mainly taken from the RC2 (de Vaucouleurs et al., 1976), the original sources, or are rough estimates.

Note that in most cases the diameters are D values, i.e., 25 25 mag arcsec-2 diameters. For those galaxies whose mean -2 surface brightness is fainter than 25 mag arcsec , the diameter must refer clearly to a fainter isophote. The plotted diameter data are therefore not in a consistent system, and this must be kept in mind when the data in

Fig. 2 are discussed quantitatively. 10

r------"\ 20 I ' I " \ I \ I"\ 19 : Domain of \ • I\ Space Tetescope \ I\ , I\ \ 18 1\' I' \ I' , I\ \ I\ , • 17 I\ .\ I\ ' I\ \ I \ \ 16 ! \ •\ \ \ 15~5 \ xx, \ \ 15 \ \ .... ~!\ ID \\. . \ E \ . \ \ 14 \ \ \ \ . \ \ xx \ \ 13 •• • • • \ \JeoB \ \ , \ \ 12 \ \ oDraco \ .. \ \ l.Jy1" . t \ \0 , \ \ rJrmgl c· 11 \ \ \ \ \ 10 \ \ \ 22,5mag/c·' • o 00 9 0 6" 10'

Fig. 2. The dwarf galaxies of the 10 Mpc sample plotted in the apparent magnitude-log diameter plane. Open circles are dEs, crosses are dIms. Small crosses are for dIms whose apparent magnitudes (and diameters) must be in error.

Most dIms lie within a triangle defined by the limits:

(1) apparent magnitude JIB > 8~5, (2) diameter D > 55", and T -2 (3) mean surface brightness roB < 25 mag arcsec . Very few - 11 - dwarfs are fainter than ~ = 15~5; indeed they all have to T be brighter than ~ = 16~5, as long as D > 55" and mE < 25 -2 T arcsec A mean surface brightness of 25 mag arcsec-2 corresponds to 10% of the blue night sky brightness (mE ~ -2 22.4 mag arcsec ), and that is the realistic limit to which a surface brightness feature can be expected to be found on a photographic plate.

There are, however, a number of dIms whose surface -2 brightnesses are considerably fainter than 25 mag arcsec ; -2 in two cases they are even fainter than 27 mag arcsec At these surface brightnesses they would clearly be undetect- able. Therefore either their apparent magnitudes or their diameters, -or both, must be in error. In fact an error in diameter alone could hardly remove the discrepancy, and hence one is forced to conclude that the luminosity of some dIms is severely underestimated, in the most extreme cases m ~ossibly by as much as 3 .

The corresponding conclusion is not indicated for the dEs, although several of them have a surface brightness -2 fainter than 25 mag arcsec The reason is that they were not detected because of their surface brightness, but because of the concentration of their brightest resolved stars.

The Space Telescope will widen the detection limits in

Fig. 2 in three ways: (1) The faint apparent-magnitude - 12 - cutoff will be pushed toward fainter magnitudes by almost Srn, (2) the minimum diameter will be decreased by nearly a factor of ten, and (3) the detectable surface luminosity will be lowered to 27 mag arcsec-2 due to the faint extraterrestrial sky brightness. Note that it is unlikely that any dwarfs exist below the ~ = 22.S rnag arcsec-2 line in Fig. 2, because low surface brightness is apparently an inherent property of dwarf galaxies (excluding here again the inter galactic H 11 regions).

Consider first what the new detection limits mean for the dIms. The present 10 Mpc sample shall be pushed outward to ~SO Mpc, increasing the sample volume by a factor of 12S.

Assume then, in rough approximation, that constant space density leads to 76 x 125 = 9500 dIms or a surface density of 0.2 dIms per square degree, and you will realize that a search for these dwarfs with the Wide Field Camera (WFC) - which has a frame size of 2.10-3 square degrees - would be absolutely hopeless. This estimate, however, is too conservative. Dwarf galaxies are actually found from the ground at distances of

22 Mpc down to mB~19m and diameters of a few arcseconds

(viz. in the Virgo cluster). The Space Telescope should there fo~e detect dwarfs at distances at least five times larger, i.e. at-100 Mpc. This distance limit increases the surface density of dIms to 1.8 dIms per square degree. This number should be further increased by an unknown factor for those dIms with surface brightness between 25 < mB< 27 mag - 13 -

,; 24 E dwarfs,_ 01 -C \ ",,,/' .§ " )// ~\\, E25 -' ..-- Im dw

2000 3000 6000

Fig. 3. The wavelength dependence of the extra terrestrial night sky brightness and the colours of dIm (full line) and dE (dashed line) galaxies. The dwarfs are arbitrarily normalized at a blue surface brightness of roB = 25 mag arcsec-2 •

-2 arcsec Whatever this factor may reasonably be, a syste-

matic search for dIms with the Space Telescope remains

prohibitive.

This conclusion is not changed if one turns to groups

of galaxies. In spite of the expected increase in space

density, the area of well defined groups is too large for a

survey. The situation is even worse in the Virgo cluster,

because there seem to be no dIms there rainter than m ~16m B (Reaves, 1980; Sandage and collaborators, 1980).

As to the dEs it is clear from the remarks above that it

is impossible to decide with the Space Telescope in a systematic way whether they exist in the field or not. The 14 - Galaxy has six known dEs within 220 kpc. To search for the corresponding dEs around M8l (at a distance of 3.3 Mpc), one would have to search an area of 15 square degrees, which is reduced to a still impossible value of three square degrees in the case of MlOl (at 7.3-Mpc). The situation is more favourable in the centre of the Virgo cluster, where the m surface density of dEs with ~ ~ 19.5 is about 20 per square degree. If the luminosity function of faint dEs increases as log N cr 0.2 m, one expects 200 dEs per square degree to a m m limiting magnitude of ~ 24.5 (M -7.2), or 0.4 dEs per = B = WFC frame. This number is roughly to be doubled because of the low-surface brightness dEs with 25 < roB < 27 mag arc- -2 sec With a few dozens of frames one would therefore be in a position to make a meaningful statement about the faint end of the luminosity function of dEs.

At the distance of the Coma cluster, which is 6.0 times more distant than the Virgo cluster, one expects 8 dEs m brighter than ~ = 27 per WFC frame, whose number is to be

increased again by the low-surface brightness dwarfs and by

the considerably higher space density of Coma cluster galaxies. Therefore~ the luminosity function of dEs could be well defined down to M -8~6 with only a few WFC frames of B = the centre of the Coma cluster.

We have tacitly assume~ so fa~ that the search for dwarfs should be carried out in the blue passband. The 15 faintness in blue light of the extraterrestrial night sky, taken here from Barbieri (1980), and the relative blueness of dwarf galaxies, as indicated by the few available UBV measurements (de Vaucouleurs et al., 1976; Visvanathan and

Sandage, 1977), suggest that a search in the U passband might be more promising. As Fig. 3 shows, the gain is appreciable only for dIms, a systematic search for which is out of the question anyhow. The search for dEs would be significantly facilitated only if the faintest dEs were still considerably bluer than those for which colors are presently available.

Partial financial support by the Swiss National Science

Foundation is gratefully acknowledged. - 16 -

References

Baade, W. 1944, Ap. J. 100, 137.

Barbieri, C. 1980, Internal E.S.A. Faint Object Camera

Report.

Bergh, S. van den 1959, Pub1. David Dun1ap Obs. ~, No. 5. Hodge, P. W. 1979, P.A.S.P. 21, 28. 1960, P.A.S.P. 72, 188.

1971, Ann. Rev. Astron. Astrophys. ~, 35.

Huchtmeier, W. K. 1980, this workshop.

Kraan-Korteweg, R. C., and Tammann, G. A. 1979, A. N. 300,

181.

Reaves, G. 1956, A. J. ~, 69. 1967, in: Modern Astrophysics, M. Hack ed.,

Gauthier-Vi11ars (Paris), p. 337.

1977, The Evolution of Galaxies and Stellar

Populations, B. M. Tinsley and R. B. Larson eds., New

Haven, p. 39.

1980, in preparation.

Sandage, A. 1975, Galaxies and the Universe, A. and M.

Sandage and J Kristian eds., The University of Chicago

Press, p. 1.

Sandage, A. 1980, in preparation. - 17 -

Sandage, A., and Tammann, G. A. 1980, A Revised Shapley

Ames Catalog, Carnegie Institution of Washington

(Washington) .

Vaucouleurs, G. de, Vaucouleurs, A. de, and Corwin, H. G.

1976, Second Reference Catalogue of Bright Galaxies,

The University of Texas Press (Austin).

Visvanathan, N., and Sandage, A. 1977, Ap. J. 216, 214.

Wilson, A. G. 1955, P.A.S.P. 67, 27.

Zwicky, F. 1957, Morphological Astronomy, Springer (Berlin),

p. 224. - 19 -

Galaxies of Magellanic Type: Preliminary Report

J. V. Feitzinger Astronomisches lnstitut, Ruhr-Universitat Bochum, BRD

Using the ESO-B Atlas 580 new galaxies of Magellanic type - partly dwarf systems according to their luminosity class - were classified; some examples are shown in fig.1. The sample shows dominant bar structures at the classification stages d , dm and m. A striking feature is the asymmetric position of. bar and disk. Tab.1 summarizes the relative frequencies of the sample galaxies.

Tab.1 Relative frequency of Magellanic type galaxies (N = 580); values in parentheses are uncertain.

cd d dm m SA 4% 5% (1%) (1%) SAB 5% 4% 3% (2%) SB 12% 14% 13% 14% lA. (1%) lAB (2%) IB 15% I 3% Pec (1%)

Further morphological studies are in progress. Surface photometry on large scale plates (3.6 m ESO) and spectral investigation of these galaxies is urgently needed. EUV studies from space will be fundamental for the under standing of the relationship between star formation and evolution of these galaxies. - 20 -

. . •

91 Fig.1 Examples of Magellanic Type Galaxies

Number Object Object§ Type Scale (") 1 048-G17 2204.5-7336 SB(s)m IV 27 2 048-G23 2224.4-7610 SAB(s)d 11 - III 24 3 079-G02 0029.9-6439 SB: (s)m sp 35 4 079-G05 0037.8-6343 SB(rs)cd Ill-IV 35 5 079-G07 :SB(s)dm IV 34 6 081-IG10 :SB(s)d Ill-IV 16 7 082-IG06 :SB(s)m IV 17 8 083-IG14 :SB(s)m IV 19 9 085-G14 0454.2-6606 SB(s)m IV 32

§after Corwin et al. 1977,1978; A.J. 82,557; A.J. 83, 1566 - 21 -

A more detailed report (Galaxies of Magellanic Type) is in press in Space Science Reviews (1980).

-22 .--,.----.---...---~ Se SbSaE£

-19

-16 -13 (

o 0.25 0.50 0.75 a-v 1.00 ~ refers to the S systems.

Fig.2 Modified and supplemented colour - magnitude diagram after van den Bergh (1977) and de Vaucouleurs (1977). Active dwarfs can be found below B-V= 0.3.

References de Vaucouleurs G., 1977, Topics in Extragalactic Astronomy, Occasional Reports No.2, Roy. Obs. Edinburgh van den Bergh S., 1977, in Evolution of Galaxies and Stellar Populations, B.M. Tinsley, R.B. Larson Eds., p. 19

Discussion

R. Shaw: Do you have any idea of the absolute magnitude range covered by your catalogue? - 22 -

J.V. Feitzinger: The aim of this catalogue is to present a sample of Magellanic type galaxies. Type classification (determination of the morphology of the stellar system) needs an image scale greater than 1 mm2 on the ESO -B plates (>'1000 pixels/mm2 ). Galaxies greater than this limit are included. Therefore no apparent magnitude or absolute magnitude range is taken into account. A rough guess of the absolute magnitude range can be obtained by looking at the luminosity classes of these galaxies (fig.2). The covered range of the absolute magnitude is approximately -18.5 <:' M<-15. - 2) -

A Survey of Galaxies with Velocities Smaller than 3000 km S-1

Hugo van Woerden, Kapteyn Astronomical Institute, Groningen

The standard galaxy cataiogues are known to be quite in complete at any magnitude or diameter limit. In particular, they favour the "earlier" morphological types, which have higher surface brightness. Hence, distribution functions of galaxy properties based on such catalogues are severely biased.

In the early 1970's Dr. R.B. Tully undertook an attempt to remedy this situation. He searched the ~alomar Observatory Sky Survey for galaxies which, on the basis of morphology and angular diameter, he estimated to probably have velocities V smaller than 3000 km s-1. The 1400 galaxies so identified were observed in the 21-cm line by Dr. J.R. Fisher at NRAO; about 70% were detected, and most of those indeed turned out to have V ~3000 km s-1. Many of these objects (though not all) are late spirals, irregulars, or dwarf galaxies. As a result of this survey, the population of the local volume of space is now much better known than before. A preliminary discussion of results, including the galaxy luminosity function, has appeared in IAU SYmposium No. 79.

The Tully-Fisher survey was, of course, limited to decli nations above _45°. A southern supplement was planned by Tully and myself in 1974. Tully searched the ESO-B Survey plates for probably-nearby galaxies, following the same criteria as in the North. W.M. Goss, U. Mebold, Mrs. B. Siegman and myself made two series of 21-cm line observations with the 64-m telescope at Parkes in 1975 and 1976. Subsequent series of observations were carried out, as the ESO-B Sur- vey search progressed, by Goss and Mrs. Siegman in 1977 and by P. Chamaraux, J.D. Murray and Mrs. Siegman in 1978.' Together with Fisher and Mrs. Siegman and several Mt. - 24 -

Stromlo students, I carried out the (almost) final series in February 1980.

The Southern list compiled by Tully contained about 430 objects, some of which already appeared in the Reference Catalogue. Including observations made in a survey of Reference Catalogue galaxies (van Woerden, Reif, Mebold, Goss, and Siegman, in preparation; see also Froc. Astron. Soc. Australia 1, 68, 1976) we have detected about 250 -1 galaxies, that is N60%. Most of these again have V~3000 km s although a considerable number are at higher velocities. Analysis of the material is well underway.

I wish to thank the Division of Radiophysics, CSIRO, for a . fellowship and for hospitality in 1975/76 and 1980.

Discussion

G. A. Tammann: With the exception of some clusters, the only way to increase the sample of known dwarf galaxies is,

as we have seen, -- by means of ground based optical and

21 cm observations. It is good to hear that new searches are on their way. One should, however, not overestimate the yield concerning dwarf galaxies. We have already seen that the magnitude-limited Shapley-Ames CataJog contains only 0.6 percent of dwarf galaxies. Even van den Bergh's 000 Cata

logue of dwarf galaxies, which is based on morphological - 25 - criteria for dwarfism (low surface brightness and flat m luminosity profiles), contains only ~30 dIms (~ > -16 ) out of 179 galaxies detected in HI (Fisher and Tully, 1975,

A. A. !i, 151); for the Nilson dwarfs the corresponding numbers are 26 and 145 (Thuan and Seitzer, 1979, Ap. J. 231,

680). One would have to apply extremely strict criteria to obtain a purer sample of optically selected dwarfs. This is actually not so surprising because the surface brightness of many dwarfs lies considerably below the night sky brightness.

H. van Woerden: The 76 dIms you were speaking of, do they -1 include all known dwarfs up to, say, 3000 km s ?

G. A. Tammann: No. The 76 dIms constitute only the "10 Mpc sample" (i.e., V ~ 500 km s-l). But at V > 500 km s-l a o o m dwarf is already fainter than ~ = l4 , which means that not many distant dwarfs are known yet. I have honestly not made a literature search to determine their total number.

H. van Woerden: Then, if one scales galaxy numbers by velocity, one should have ~4000 dwarfs in the South Polar Cap

(b < -30°) with v < 3000 km s-1 We are indeed very-far o from knowing even a fraction of these dwarfs. -- As to our survey, it should be pointed out that the last observations were made in February, 1980, and this explains why nothing is published yet. -- A difficulty for us is that photometry of our southern galaxies is virtually non-existent. This sorry state is now being rectified, but progress is difficult and - 26 - slow.

J. V. Feitzinger: We start a programme of southern galaxy photometry with the Bochum telescope on La Silla.

M. Dennefeld: Is the limit of 25 mag arcsec-2 a detection limit, or simply the limit where the classification of dwarf galaxies is possible?

G. A. Tammann: It is the actual detect~on limit, unless one presses things extremely hard.

B. J. T. Jones: To extend the present sample of dwarfs we have to find better ways in preselecting dwarfs from photo graphic plates. Present dwarf candidate lists are contaminated with brighter galaxies, because you cannot go below a certain surface brightness level in the visual detection.

It would be great to have somebody stand up and propose how we can actually go below the present detection limit.

M. J. Disney: For example, if one uses one of these techniques like David Malin, at AAO, and brings out the lowest surface brightness on IIIaJ plates. I don't know exactly at what level it stands, but I think you should see features -2 down to perhaps 28 mag arcsec [?]; I think in that case we should find a lot more dwarfs. But the main problem is that on such deep plates all kinds of funny things show up.

B. J. T. Jones: Are you suggesting we "Malinise" the whole sky? 27

M. J. Disney: Yes.

H. C. Arp: Madore and I are just completi~q a 4 to 5 year- long search of particular galaxies on the UK Schmidt survey.

I think we have about 20,000 objects, 1% of which seem to be dwarf galaxies of extremely low surface brightness. Some of these ~200 dwarfs show resolution, others don't. -- I think searches for dwarfs should not be made with large telescopes, but with Schmidt telescopes owing to their wide field and their ability to go to a very faint surface brightness. This is where we are going to succeed.

C. Barbieri: The WFC of the Space Telescope will observe approximately 1700 hours per year in the serendipity mode, thus one will find a number of dwarfs on these frames.

G. A. Tammann: This is an essential point. It raises, at very faint levels, however, a new question: can we distinguish between low surface brightness dwarfs and very distant galaxies whose surface brightness is faint because of the redsqift? The surface brightness within a metric diameter decreases like (1 + z)4, the isophotal surface brightness somewhat slower, -- the exact behaviour depending on the luminosity profiLe of the galaxy. Unless evolutionary effects come into the game, I could imagine that the distinction becomes quite difficult.

A. Ardeberg: I had a repbrt from H. E. Schuster here on dwarf galaxies, but most of his suggestions have been 28 covered or bypassed by now. Let me mention only one point, which he stresses sharply. Many low-surface brightness features on ESO sky survey plates have been mistaken for plate flaws. If the same fields are repeated, however, these features frequently show again, indicating that at least some of them are dwarf galaxies. Plates with suspi cious features should be repeated.

P. o. Lindblad: It was mentioned today that some local dEs are known only because they are resolved. At the distance of the Virgo cluster dEs remain clearly unresolved, and we can detect only those Virgo dwarfs which have a sufficiently high surface brightness. I understand that some dEs have nuclei. Could the nuclei help to detect new dwarfs, possibly even those of low surface brightness?

G. A. Tarnrnann: I cannot answer your question for two reasons: (1) I do not know if and how such nuclei can be distinguished from stars; and (2) we don't know yet whether the very-Iow-surface brightness dwarfs can also have nuclei.

B. J. T. J0nes: I would like to introduce a fourth class of dwarfs: those galaxies found only in HI with no optical counterparts. Sargent and Lo found some of these and later looked harder at the plate and saw something. There may be a vast number of objects which have no optical counterpart.

G. A. Tarnrnann: Dr. Huchtmeier will have more to say about these dwarfs in a later Session. 29

M. J. Disney: In some clusters there seem to be void spaces where there are no big galaxies at all. They would be good

places to look for very-Iow-surface brightness galaxies.

G. A. Tammann: This is true if the empty spaces are not

the effect of subclustering; they could, of course, be real

voids. -- But I fully agree that clusters are the best place

to search for dwarfs, particularly if the Space Telescope is

used. In conclusion, I am disappointed that we do not

really know of new ways to search for dwarfs. This is even

more true if we are not content to simply increase the num

ber of known dwarfs, but if we hunt for dwarfs with particu

larly significant properties, like lowest luminosities and

faintest surface brightnesses. To answer, then, our

opening question, -- how can we search for dwarf galaxies with the Space Telescope? -- we expect to pick up some

dwarf galaxies in the serendipity mode, we should devote a number of frames to search for very faint dwarfs in the centre of the Virgo cluster, and we should determine the number of dwarfs in the centre of the Coma cluster and

possibly their radial density distribution within that cluster.

H. van Woerden: I am not as pessimistic as this. I think we can do with improved techniques from the ground and with the Space Telescope much better than we have done so far. - 31 -

SESSION 11

(Chairman: S. Laustsen)

VARIABLE STARS IN DWARF GALAXIES - 33 -

VARIABLE STARS IN SPHEROIDAL GALAXIES

S. van Agt

The dwarf spheroidal galaxies provide us with a possi bility of checking observationally the theoretical pre dictions of conditions in space away from our Galaxy.

The dwarf spheroidal galaxies appear on the sky as very loose agglomerates of stars, no gas or dust is traced

thus far. There are seven of these objects known within a distance of 280 kpc and none are nearer than approxi mately 60 kpc. If they were, they would have been dis rupted by Galactic tidal forces. It is therefore expected that the dwarf spheroidal galaxies evolved isolated from, and uncontaminated by, our Galaxy. They are a possible means to trace back the conditions in the intergalactic medium at the time the members of the Local Group formed.

In' a general way one can say that the stellar population of the dwarf spheroidal galaxies is comparable with the old population II stars as found in the Halo globular clusters of our Galaxy. The colour-magnitude diagram for a number of the Sculptor type systemsis not unlike that of globular clusters. The dwarf spheroidal galaxies con tain large numbers of variable stars of the RR Lyrae type, which are found in lesser quantities in the globular clusters.

The discrete distribution of the characteristics of the

RR Lyrae in galactic globular clusters is not followed - 34 - by the short period variables in the dwarf spheroidal galaxies. Also in the colour-magnitude diagram of the dwarf spheroidal galaxies, the detailed position of the stars can only - in a first approach - be fitted with the results obtained from observations of globular clus ters. Survey papers on the variable stars and population of stars in dwarf spheroidal galaxies are given by van den Bergh (Comm. David Dunlap Obs. no. 195, 1968), van

Agt (Variable stars in globular clusters and related objects; ed. J. Fernie, 1973, p. 35), van den Bergh (Ann.

Review Astron. and Astrophys., 1975,11, 217) and Hodge

(Ann. Review Astron. and Astrophys., 1971, 9, p. 35).

The cepheids, which are well represented in a number of

Galactic globular clusters and which have periods longer than one day, have not been found in any of the dwarf spheroidal galaxies investigated thus far. But a type of variable has been found which seems typical for dwarf spheroidal galaxies. The period of these bright anomalous cepheids ranges from approximately 0.5 day to, not in ex cess of, 2.8 days, whereas the luminosity of these vari ables is about 0.5 to 0.8 mag brighter than the luminosity of population II cepheids with the same period. Although the bright anomalous cepheids seem typical for dwarf sphe roidal galaxies, mention has been made (by Zinn and Dahn,

A.J., 81, 527, 1976) of a variable in the globular clus ter NGC 5466, which might belong to this class. - 35 -

Their nature is commonly explained in terms of a larger mass than the RR Lyrae in the same dwarf galaxy. The ques tion - "how differences of the required order can exist or develop between the member stars in a dwarf spheroidal galaxy mass" should be answered.

A number of carbon stars have recently been found in the field of the Fornax and the Sculptor dwarf galaxy by Wes terlund (ESO Messenger No. 19" 1979) .

Work on the bright anomalous cepheids has been done by

Demarque and Hirschfeld (Ap. J., 202,' 346, 1975), Norris

(Ap. J., 202, 335, 1975), Zinn and Searle (Ap. J. 209,

734, 1976), Renzi~i et al. (Astron. Astrophys., ~, 369,

1977). Recent research into the nature of the dwarf sphe roidal galaxies is aimed at observationally improving the photometry and spectro-photometric work of the brighter stars at the tip of the giant branch, and of the giant branch stars (Kunkel and Demers, Ap. J., 214, 21, 1977;

Stetson, A. J., 84, 1149, 1979; Schommer et al., private comm. ).

A survey of variable stars in dwarf spheroidal galaxies is given in table I. - 36 -

Table I nr. of' nr. of' other vars. RR Lyrae anom8Jl. vars. Name 1 b total ab c ceph.

Fornax 237 - 66

Sculptor 286 - 83 603 9 3 1 LPV 1 Irr red var.

Draco 86 + 35 261 126 7 5

Ursa Minor 103 + 45 92 21 13 3 1 RR part of' eclipsing Algol syst.

Leo I 226 + 49 25 25*

Leo II 219 + 67 196 64 6 5 6 Irr red var. 6 Irr blue var.

Carina** 260 - 22

* Hodge and Wright, A.J., ~, 228, 1978 ** Discovered by Cannon et al., M.N.R.A.S., 180, 81, 1977

Discussion

A. Ardeberg: In most of' these surveys image crowding seems to be a serious problem. I wonder why we shouldn't take som~ ef'fort to do it electronographically.

S. van Agt: Image crowding is not such a problem in the Draco, the Sculptor and the Leo I and Leo II systems. I'm working in Sculptor and there it is definitely not neces- - 37 - sary to do it electronographically. Though in the other systems, in Carina for example, it might help us.

G. Tammann: You have in your table the total number of RR Lyrae stars. To what extent can they be used to cor relate with the total luminosity of a system? If these are complete samples and if we know the total magnitudes of the systems, then the question is, whether or not the number of RR Lyr stars is proportional to the total luminosity or the total mass of a system.

S. van Agt: Well, it should be said that only the search for Sculptor and Draco has been completed.

G. Tammann: One would like to see a colour-magnitude dia gram of these galaxies down to as faint a limit as possi ble, possibly down to the main sequence. Could we do that with the wide-field camera with 2:7 window? What are the diameters of these galaxies? How much could we cover with one frame?

S. van Agt: The diameters are quite large, up to about two degrees. 39 VARIABLE STARS IN DWARF Im GALAXIES s. Laustsen

The dwarf spheroidal galaxies are, as we just heard, well suited for stellar photometry, whereas the dwarf Im galaxies are a lot more diffi cult. In the Im galaxies the stars concentrate towards the central par~ and even in the outer parts the stars often appear in clumps; the galaxies normally have huge emission regions and they contain large amounts of dust. For these reasons the variable star research is, apart from the Magellanic Clouds, not too well advanced for Im galaxies. The Space Telescope with its high resolution capability can, especially for moderate angular size dwarfs, contribute much to such research.

In the Magellanic Clouds, variables of several types have been studied, as listed for instance by C. Payne-Gaposchkin (1971) at an early ESO Symposium. The Cepheids dominate in number, and in scientific importance as well, it seems. At the time of a study of van Genderen (1969) 1240 Cepheids were known in the SMC and he estimated a total of 2000 Cepheids in that galaxy, and Becker et al. (1977) have estimated the same number in the four times more luminous LMC.

Table 1 presents a list of Magellanic type dwarf galaxies considered as members or probable members of the Local Group. The list is made up from the catalogue of Kraan-Korteweg and Tammann (1979) and includes luminosity class IV-V and V objects only. The luminosity L is copied from the same catalogue. Cepheids have been studied in only a few of these galaxies. rc 1613 was studied by Baade and Sandage (Sandage 1971) who found 37 definite and 4 probable Cepheids. Kayser (1967) found 13 Cepheids in NGC 6822 while Laustsen and Tammann (1977) detected a few probable Cepheids. in IC 5152.

In order to plan for future observing, it would be good to have some estimates of expected Cepheid numbers. The best we can do at present is to assume proportionality of Cepheid number with galaxy luminosit~ based on an average of the LMC and the SMC. The factor of four between the two illustrates the uncertainty of the procedure. The estimates are listed in Table 1 together with angular sizes expressed in the number of ST-WFC - 40 - frames needed to cover the galaxy. The SMC will, obviously, have to be studied from the ground and the same is true for the outer, extended parts of several other dwarfs. But the smaller size dwarfs as well as the central, dense parts of others will need the higher resolution capability of the Space Telescope.

TABLE 1

Magellanic Type Dwarf Galaxies which are Members or Probable Members of the Local Group.

Name 000 No. L Cepheids Size

Observed Est.total WFC frames

SMC 964 1240 2000 6000

Sextans B 70 240 200 5 WLM 221 203 170 15 NGC 6822 209 198 13 170 50 sextans A 75 191 170 10 rc 1613 8 130 41 110 70 rc 5152 93 3 80 2 Pegasus 216 50 40 5 187 37 30 2 Leo A 69 27 20 5 GR 8 155 16 15 1 210 5 5 2 Sag DIG 2 2 2

Although rather few Cepheids are to be expected in the smallest dwarfs, these ought to be carefully studied for brightest stars and variables. Small dwarfs are believed to have a less complex structure and development than galaxies several orders of magnitude bigge~and could serve to give us basic information on the formation of stars and galaxies. - 41

References:

Becker, S.A., Iben, I., Jr., Tuggle, R.S., 1977, Astrophys. J., 218,633. Kayser, S., 1967, Astronom. J. ~, 134 Kraan-Korteweg, R.C., Tmmnann, G.A., 1979, Astron. Nachr. 300, 181 Laustsen, S., Tammann, G.A., 1977, The Messenger, No. 11 Payne-Gaposchkin, C., 1971, in the Magellanic Clouds, ed. A.B. Muller (Dordrecht-Holland: Reidel) p. 34 Sandage, A., 1971, Astrophys. J. 166, 13 van Genderen, A.M., 1969, Bull.Astr.lnst.Netherlands Suppl. l, 221

Discussion

F. Macchetto: When you say you need a certain number of frames I SUppose you need a frame of a certain colour and at a given time. How many frames do you need altogether? How many colours and how many exposures for each colour?

S. Laustsen: The best person to answer these questions is Gustav Tammann.

G. Tammann: I think it would be terribly difficult to get the period with less than 20 exposures. I think that is an absolute minimum, and since there is no period-luminosity relation but only a period-luminosity -colour relation, one should have it in two colours, so it would be 40 exposures per field.

F. Machetto: And how long would each of these exposures have to be?

G. Tammann: It's almost instantaneous, say 20 seconds. 42

M. Disney: Would it help to get ultraviolet colours? And is the age the same for all dwarfs? s. Laustsen: I wonder whether ultraviolet colours would help a lot as long as most variable star photometry elsewhere is done in Band V. As to the age problem, I do believe that the smallest dwarfs are more simple than the giants, also as to their ages. But we do not know whether in these very small galaxies most of the stars were created in a single burst of star formation, in a few bursts or in a more or less continuous formation of stars. So far we do not know much about ages. SESSION III

(Chairman: J. Lequeux)

Im GALAXIES

(a) Frequency Determination, Galaxy Luminosity Function 45

THE LUMINOSITY FUNCTION OF DWARF GALAXIES

G. A. Tarnrnann

ESO; Astronomical Institute, University of Basel

The term luminosity function is uSBd with two different meanings. For some it is the number of objects within a given absolute-magnitude interval per unit volume; for others it is just the probability that an object randomly chosen from a sample volume falls into a certain absolute- magnitude interval. If we choose the former interpretation, we can stop right here. We have already talked, in Session It about how little we know about the space density of dwarf galaxies in the field, in groups, and in clusters, respectively.

All we can hope at present, is to come to grips with the latter luminosity function of dwarf galaxies.

Turning first to the dwarf ellipticals (dEs); it is clear that locally we know far too few of these objects to put any useful limit on this luminosity function. From work in progress by G. Reaves and A. Sandage and collaborators concerning the dEs in the Virgo cluster, it seems that their number per magnitude interval increases down to the detec- m m tion limit of m ~ 19.5 (M ~ -12 ), and that their lumi- B B nosity function is the faint extension of the luminosity function of normal e1liptica1s. The preliminary data certainly allow an exponential luminosity function of the form log N(M) ~ aM, where Zwicky's high value of a = 0.2 is 46

still possible. The exciting question is, of course, at which magnitude N(M) reaches its maximum. Zwicky played with the idea that the exponential increase holds down to

individual stars. If a fraction of the dwarf galaxies are

the debris of interaction between larger galaxies, that

seems not impossible. We have seen this morning that the

only hope to decide if and where N(M) has a maximum rests on

Space Telescope observations in the Virgo and Coma cluster.

One interesting point should be noted in passing: the m faintest known dEs (~~ -8 ) are fainter than the brightest

globular clusters (~ < -lam). (There is some' evidence that

the luminosity function of globular clusters is Gaussian;

Harris and Racine, 1979).

It is my personal belief that the luminosity function

of dIms differs significantly from that of dEs, i.e., that m the former has a maximum which lies brighter than ~ = _14 . The evidence for this is twofold:

(1) The absolute-magnitude distribution of the Sc and

late galaxies in the Local Group and in the M81 group is

shown in the figure (cf. Tarnrnann and Krann, 1978); the m sample is certainly complete down to M ~ _14 (and b ~ 20°). B The distribution can reasonably well be fitted by a Gaussian m m with the maximum at ~ ~ -16 and a(~) = 3.3. The Gaussian fit is not correct for the brightest galaxies, because their

luminosity function for mainly el1ipticals in clusters - 47 -

SlIm I complete for N in LG and M 81 Group ~ M ~ - 14':"0

(Sandage, 1976; Schechter, 1976) and for pure spirals in the field (Tamman et al., 1979) is definitely steeper. But some bell-shaped, possibly skewed function is indicated. Note that additional, very faint dIms, which may be discovered in the future, tend to shift the maximum toward fainter magnitudes and to increase cr(~). To obtain an exponential increase as for the dEs (a = 0.2), one would have to find m still ~210 relatively bright dIms with -16 < ~ < -10 in the two galaxy groups considered, -- which seems nearly impossible.

(2) It seems that in the Virgo cluster dIms are very m m rare with ~ ~ 16 (M ~ -16 ). Although the classification B becomes difficult at faint levels, there is no doubt that 48 - the bulk of the Virgo dwarfs is of type dE (Sandage and

Shaw, 1980).

References

Harris, W. E., and Racine, R. 1979, Ann. Rev. Astron. Astrophys. 12, 241. Sandage, A. 1976, Ap. J. 205, 6.

Sandage, A., and Shaw, R. 1980, unpublished.

Schechter, P. 1976, Ap. J. 203, 297.

Tammann, G. A., and Kraan, R. 1978, I.A.D. Symp. 79, 71.

Tammann, G. A., Yahil, A., and Sandage, A. 1979, Ap. J. 234, 775. A CLOUD OF DWARF IRREGULAR GALAXIES?

S. Laustsen, P. Gammelgaard Astronomisk Institut, Aarhus Universitet DK-8000 Aarhus C, Denmark

H. Pedersen ESO La Silla, Chile

In a discussion on dwarf irregular galaxies it should be of interest to call attention to the cloud of objects seen close to the southern galaxy NGC 5291 and its companion, the Seashell galaxy in the IC 4329 cluster. These objects are probable candidates for the not very long lis~ of known Im dwarf galaxies. The cloud is rather distant, about 80 Mpc, and was only recently found by Pedersen et al. (1978) and independently found and studied by Longmore et al. (1979).

According to Longmore et al. (1979) a band north and south of NGC 5291 contains about two dozen non-stellar knots of which the largest in size are comparable to the r~gellanic Clouds. Spectra have been taken for a number of knots and show a continuum as well as emission lines of HII regions. The knots are enclosed in a huge cloud of neutral hydrogen, lOll solar masses, the center of which is located to the west of the band of knots. Knots and neutral hydrogen show the same sort of motion and are probably rotating.

Deep plates taken by us at the ESO 3.6 m te~escope, prime focus, have revealed hundreds of non-stellar knots. Most of them are concentrated in the band north and south of NGC 5291 but also outside this band, and in particular to the west of it, we find quite a number of knots, thus giving a better agreement with the distribution of neutral hydrogen. Comparison between red and blue plates clearly shows the knots to be more blue than normal galaxies in the field. One of our plates is shown in Fig. 1.

Longmore et al. (1979) have from spectrophotometry derived an average B magnitude for the brightest knots of B = 17~4, corresponding to 50

~ = -17. But no photometry of fainter knots is available and no lumin osity function can yet be derived. Pending such work to be done we have measured angular sizes of 105 knots in the dense band north and south of NGC 5291 and presented the result in a histogram (Fig. 2). The linear scale in kpc, corresponding to a distance of 80 Mpc, is given and, for comparison, the linear sizes of four well known dwarf galaxies of the Magellanic type are indicated. The counts are, naturally, incomplete for the smaller size knots, and we cannot tell whether or not the observed maximum at about 1.5 kpc is real.

Considering the appearance of the knots on our plates, their colour, their spectra, and their linear dimensions, it is most likely that we are here faced with a cloud of a few hundred galaxies of the Im type. The galaxies are probably young or we may, indeed, be witnessing the process of galaxy formation on a large scale.

A lot more observing is needed, however. Photometry and spectroscopy of many more knots are within reach with ground based large telescopes, whereas the better resolution and the fainter limit of the Space Tele scope will be needed for detailed photometry. With a distance modulus of 34.5 the brightest Cepheids should be accessible. The SMC has 3 Cepheids brighter than M = -7. An advantage here is that one single frame of the Wide Field Camera of the Space Telescope in the dense areas north or south of NGC 5291 will hold some 40 knots, which means that some 40 dwarf irregular galaxies covering a good range of magnitudes can be observed simultaneously.

References:

Longmore, A.J., Hawarden, T.G., Cannon, R.D., Allen, D.A., Mebold, V., Gass, W.M., Reif, K., 1979, Mon.Nat. R. astr. Soc. 188,285

Pedersen, H., Gammelgaard, P., Laustsen, S., 1978, The Messenger, No. 13 - 51

Figure 1. ESO 3.6 m prime focus photograph showing knots around NGC 5291. The plate is an unfiltered 60 min. exposure on IIIa-J. North is up and east to the left. The irregu lar band of knots north and south of NGC 5291 has a total extent of 9' which at a distance of 80 Mpc corresponds to 200 kpc. - 52 -

SagDIG IC 5152 NGC 6822 SMC II I I 30 1 2 3 4 5

Figure 2. Histogram of 105 knot sizes. See text. - 53 -

Discussion

H. van Woerden: The Space Te1escope cannot reso1ve individ ua1 stars in this group. The sensitivity 1imit is M N -7 for a distance of 80 Mpc and the reso1ution wou1d corre spond to 0~1 =40 pc 1inear reso1ution. Hence, we cou1d see associations, not stars.

G. Tammann: When I gave a c1assification of dwarfs this morning, I certain1y didn't mean this to be fina1 or com p1ete, and Dr. Laustsen's ta1k has brought us at 1east in the neighbourhood of a new c1ass, and this is debris. We know that ga1axies encounter, that they tear materia1 out of each other and that this materia1 exists afterwards. We know therefore there are interga1actic sing1e stars and there must a1so be interga1actic sing1e dwarfs after such encounters, and I do not want to pretend here that your case is the one or the other. It raises the funda menta1 question: does nature make dwarf ga1axies as indi vidua1s or/and as debris, or is there on1y one cosmo1ogy of dwarfs? - 55 -

(b) Hydrogen Content 57

HI Content of Dwarf Galaxies by \~. K. Huchtmeier, Max-Planck-Institut fUr Radioastronomie, Bonn. F.R.G.

Galaxies with absolute magnitudes fainter than -16 will be consid ered to be dwarfs. They are situated at the faint end of van der Bergh's (1960) luminosity classification. The luminosity function of the Kraan Tammann (1978, hereafter KT) sample of late-type galaxies within 10 Mpc peaks at M = -16. The dwarfs are galaxies with a very low surface bright ness, especially the dwarf ellipticals (dE) and spheroidals. Therefore they are difficult to be detected at great distances. All dE galaxies ex cept one in the KT sample are Local Group (LG) members. From the 73 late type objects in the KT-catalogue only 48 have been observed in the 21-cm line of neutral hydrogen (HI). Most of the others have not been observed because of their southern declination. The majority of single-dish HI data of dwarf galaxies is included in two investigations 11 : DDO galaxies have been observed by Fisher and TUlly (1975), Nilson dwarfs have been observed by Thuan and Seitzer (1979).

So far only one dwarf elliptical (NGC 205) has been detected in HI (7xl0-5 of the total mass). Upper limits to the HI in dE galaxies are of the order of 10-3 to 10-4 of their total mass (Table 1). Upper limits to any radio continuum emission are quite low, too: <0.05 Jy in the fre quency range '408 MHz to 2.7 GHz (Pooley 1969, Heeschen 1970, van der Kruit 1972).

Blue compact galaxies detected in HI are generally too bright to be called dwarf galaxies. However, I Zw114 (Chamaraux et al. 1970) is a dwarf (M = -14.7) with lOB Mo in HI. Properties of the blue compact galaxies seem to be similar to those of the dwarf galaxies.

Surveys for irregular dwarfs may be performed in the 21-cm line in principle. In fact several groups of observers have searched for iso lated intergalactic HI clouds. As the expected HI properties of these clouds are similar to those observed in dwarfs, these surveys are for dwarfs as well. Observations of HI in interacting galaxies (von der Hulst 1977, Weliachew et al. 1978), the Magellamic Stream (Mathewson et al. 1974), the interpretation of high velocity clouds (HVC) (Verschuur 1969) as intergalactic clouds lead to a model cloud with a scale length of 30-50 kpc, some 10 8 Mo of neutral gas, and a line width of ~35 km S-1. Upper limits to intergalactic HI clouds with such properties from dif ferent authors are given in Table 2. The lowest upper limits outside the LG are ~2x106 Mo Mpc-2 which corresponds to a few times 10 8 Mo in near by groups. Extended areas around some LG dwarfs have been searched to even lower limits. No new clouds .have been found. Clouds seen in pro jection onto the HI distribution of galaxies have not been counted as intergalactic. The four objects discovered by La and Sargent (1979) were identified with dwar~ galaxies. These HI lines belong to irregular

1) As this introduction does not claim any completeness in references we will not consider the occasional dwarf included in HI articles of late-type galaxies. 58 dwarfs, some of them the faintest ones observed so far. With the present ly available sensitivity in HI observations only a few new objects were discovered. Therefore the search for dwarf galaxies with optical techniques seems more adequate.

Unfortunately there is no complete sample of observations of dwarf galaxies available because of their low surface brightness and total lumi nosity. Within the DDO and Nilson sample there are more than three hundred dwarfs observed in HI. This is a diameter limited sample (blue diameter >1 arc min). Obviously there are selection effects in a way that this sample is biased towards brighter galaxies and high MHI/L values as upper limits (hydrogen poor objects) are not included in the discussion. Nearly all dwarfs from the DDO list were discovered by Fisher and Tully (1975) whereas Thuan and Seitzer (1979) detected only 145 out of 461 Nilson dwarfs. The objects from the second list are further away on the average and there are more non-dwarf objects. However their sensitivity limit does not allo~o detect galaxies with MHI/Lpg <0.1. Thuan and Seitzer (1979) believe their sample to be fairly complete within the following limits: D ~3 Mpc, o~-2~5, Ibl~200, and Mpg ~-10.5 ~' T? mini mize inclination corrections they excluded objects with incl~nat~ons <450 and retained 134 galaxies (DDO and Nilson sample together) whose global properties are discussed and compared with objects from Shostak's (1978) sample of late-type galaxies (see Table 3). There the low surface brightness (LSE) galaxies are divided into dwarfs (MPg ~-16) and non dwarfs (Mpg <-16). Generally the LSE galaxies seem to be a continuation in the morphological sequence beyond the normal irregulars (see global parameters in Table 3). LSE galaxies have larger MHI/Lpg and MHI/MT values than other late-type galaxies. To the first order MHI/MT and MHI/Lpg values are independent of absolute magnitude. Within the nat~ral scatter and observational errors the hydrogen mass-to-luminosity ratio of LSE galaxies is the same as for irregular galaxies from Roberts' (1969) or Shostak's (1978) sample. An exception are two of the Lo and Sargent (1979) dwarfs with a MHI/Lpg ratio of 15 and 18, respectively. For many of the dwarfs included in Table 3 MHI/Lpg values are greater than I, consequently there is a certain number ot considerably smaller values indicating the natural scatter in this quantity. Global para meters of a few irregular dwarfs observed in greater detail are given in Table 4. The HI-distribution of the brighter dwarfs shows generally an extent (column 5) of about the Holmberg diameter. A few show more extended HI envelopes. It is evident from their radial velocity fields that rotation is the dominant motion in these objects (column 11). There seems to be a continuous decrease in total mass and line width with absolute magnitude from brighter objects to fainter galaxies. Only for dwarfs fainter than ~ -13 random motion seems to be the dominant form of motion. There is still a small velocity gradient which may be interpreted as rotation (see Leo A in Table 4). Among the parameters in Table 4 the distance independent HI-mass to luminosity ratio (MHI/Lpg) is of special interest. From Roberts' (1969) investigation of global properties of spiral and irregular galaxies we expect a value of MHI/Lpg ~.8 for irregular galaxies; it becomes several times as large for some' dwarfs. Two dwarfs discovered by Lo and Sargant (1979) are included in Table 4. They continue the sequence to fainter luminosities. The MHI/Lpg value for M81 dA is extremely high even taking into account possibl~ errors in the total magnitude (a change of half a magnitude nearly changes the MHI/L ratio by a factor of two). 59

There might be a possible subdivision of dwarf galaxies at about M= -13. In fainter dwarf irregulars random motion becomes dominant. MHI/L is of order of 1 and greater, their linear size is about 2 kpc. In the pg case of elliptical galaxies there could be a subdivision between dwarf ellipticals (M< -13) and dwarf spheroidals (~ -13), the latter having an extremely low surface brightness.

References Allsopp, N. J. 1978, Mon. Not. Roy. Astr. Soc. 184,397 Cesarsky, D., Falgarone, E., Lequeux, J. 1977, Astron. Astrophys. 59, 15 Chamaraux, P., Heidmann, J., Lauque 1970, Astron. Astrophys. 8~424 Emerson, D. T. 1974, Mon. Not. Roy. Astr. Soc. 169,607 - Fisher, J. R., Tully, R. B. 1975, Astron. Astrophys. 44, 151 Hart, L., Davies, R. D., Johnson, S. C. 1980, Mon. Not: Roy. Astr. Soc. 191,269 Haynes, M.~, Roberts, M. S. 1979, Astrophys. J. 227, 767 Heeschen, D. 1970, Astron. J. 75,523 Knapp, G. R., Kerr, F. J., Bowers, P. F. 1978, Astrophys. J. 83,360 Kraan-Korteweg, R. C., Tammann, G. A. 1978, Astron. Nachr. 30~ 181 Lo, K. Y., Sargent, W. L. W. 1979, Astrophys. J. 227, 756 -- Materne, J., Huchtmeier, W. K., Hulsbosch, A. N. M-.--1979, Mon. Not. Roy. Astr. Soc. 186, 563 Mathe\-1son, D. S., Cleary, ~N. V., HurrflY. J. D. 1974, Astrophys.J. 190,29 1975, Astrophys. J. 195, L97 Mirabel, I. F., Cohen, R. J. 1979, Mon. Not. Roy. Astr. Soc. 188, 219 Pooley, G. G. 1969, Mon. Not. Roy. Astr. Soc. 144, 101 Roberts, M. S. 1969, Astron. J. 74, 859 --- Shostak, G. S. 1977, Astron. Astrophys. 54, 916 1978, Astron. Astrophys. 68, 341 Thuan, T. X., Seitzer, P. O. 1979a, Astrophys. J. 231, 327 1979b, Astrophys. J. 231,680 Tully, R. E., Bottinclli, L., Fisher, J. R., Couguenheim, L., Sancisi, R., van Hocrden, H. 1978, Astron. Astrophys. 63, 37 Unwin, S. C. 1980, ~lon. Not. Roy. Astr. So('. 190, ')')\ van den Ikrgh, S. 1960, PubI. David Dunlap Ohs-.-2, No. 6 van der Hulst, .1. H. 1977, Ph. D. thesis, UnIversity or Groningen van der Kruit, P. C. 1972, Astrophys. Lelt. 11, 173 Verschuur, r.. L. 19h9, Astrophys. J. 156, 771·- Weliachew, L., Sancisi, R., Cuelin, M. -1978, Aslron. Astrophys. 165,37 - 60 -

Table I: Upper limits to the HI-Mass of dE galaxies galaxy distance absolute ~I ~I reference magnitude M MT (kpc) (Me) a)Draco 67 -8.0 <68 <5.7.10-4 a UMi 67 -8.2 <280 <2.8.10-3 3 Sculptor 84 -10.2 <5.4.10 <1.8.10-3 Fomax 188 -11.90 <8.8.103 <4.4 '10-4 Leo I 220 -9.62 <7.2.103 <1.8.10-3 Leo II 220 -8.54 <1.1 .103 <1. 1. 10-2

5 b)NGC205 "'700 -15.72 3xlO 7.2xlO-5 b NGC22 I "'700 -15.33 <1.5xI06 <4.lxlO-4 c references: a) Knapp et al. 1978 b) Unwin 1980 c) Emerson 1974 - 61 -

Table 2: Search for intergalactic HI clouds

regions of sky areas ~I limits references searched for HI searched (l08 MO) (square degrees)

several hundred <2.6 106Me. / Mpc2 a Sculptor group <0.8 b, d M81 group 132 <0.8 «0.4) b, d CVnI group 90 <0.8 cl, c CVnII group <2.6 e NI023 group 92 <4.6, <4 e, e, d, g MIO) group <0.8 to N2841 group 8.108 M /H 2 pc cl Coma I group o depending on Cetus I group distance

(HI limits in paranthesis refer to smaller regions) around LG galaxies IC 1613 120 ~3.5.106 Me Sextaus 200 assuming D=1 Mpc e N6822/DDO 210 300 1=10 kpc, V=35 /1V=35 km s-1 ICIO 90

References: a Mathewson et al. 1975 b Shostack 1977 c Lo and Sargent 1979 d Haynes and Roberts 1974 e Materne et al. 1978 f Mirabel + Cohen 1979 g Hart et al. 1980 - 62 -

Table 3: Low surface brightness galaxies; DDO and Nilson sample of galaxies detected in HI

I I parameter Mpg from Shostak (1978) ~-16 I <-16 Sbe-Se . Sed-Sd Sdm-Sm Im (dwarf) I

number 40 94 45 44 3 2 diameter (kpe) 5.7±1.9 20.2±0.4

Mpg -15.O±O. 71 1-17. 9± I.I (km s-l) 495±243 1547±761 Vo , !W20 (km s-I) 100il6 I77±58 477 301 302 177 I MtII (109 Mo) O. 12±0. 13 I .44±1 .14 7.1 4.0 2.81 5.35 9 MT (10 Mo) 0.82iO.90 8.84±9.60 108 30. I 21.6 13.0 Lpg (109 L ) 0.17±0.09 3.22±3.6 29.2 10.4 14.7 14.3 , O I MHI/~ 0.21 iO. 19 0.26±0.37 0.08 0.14 0.13 0.41 Mti/Lpg i 0.66±0.43 0.55iO.36 0.26 0.44 0.25 0.61 MT/Lpg 4.56±3.4 3.37i2.28 3.86 4.03 2.31 I. 11 I ! , I I ! Table 4: Parameters of individual dwarf galaxies

, Mpg Holmberg HI- galaxy distance to ~I ~ ~I ~I Mr V velocity Re£. (Mpc) diameter Holmberg max dispersior ( 108Mf» L (km (kpc) extent (108Me) pg MT Lpg s-I) (km s-]) (1) (2) (3) (4) (5) (6) (7) (8) (9) (10) (1 ]) (]2) (]3)

SMC 0.065 -16.7 9.4 0.9 5.8 14 0.84 0.41 2.0 41 10.5 a N3109 2. -16.4 10.0 3.2 ]4. 60 2.7 0.23 12 56 12 N6822 0.70 -16.1 4.0 3.4 3.0 40 0.75 0.08 10 44 9 DD0125 4.5 -15.9 7.6 I. 13 1.3 6 0.41 0.22 1.9 19 9.4 Ho I 3.25 -14.4 5.0 1.0 1.3 3 1.6 4 20 II 0"1 0.4 \.I) Leo A 1.1 -12.7 2.2 1.5 O. 19 0.26 1.5 0.7 1.1 4 6 b

SDIG 2.9 -10.8 1.1 0.11 0.28 4 0.4 10: c M81 dB 3.25 -12.3 2.6: 0.045 0.35 0.37 d M81 dA 3.25 - 8.7 1.6 : 0.072 0.11 18.2 ,

References: a Tully et al. 1978 c Cesarsky et al. 1977 b Allsopp 1978 d La and Sargent 1979 - 65 -

HI Observations of Irregular Galaxies by W. K. Huchtmeier, Max-Planck-Institut fUr Radioastronomie, Bonn

A sample of twenty-one irregular galaxies of large angular extent has been mapped in the 21 cm HI line with the 100-m radiotelescope (half power beam width 9 arc min). Global parameters of the HI distribution have been derived directly from the maps for the extended objects and by a model-fit procedure for the smaller galaxies (extent of the HI distri bution at a level of ~19 cm-2 are given in Table 1, column 3). The abso lute magnitude (column 2) for most objects is situaLed at the bright limit of dwarf galaxies (galaxies with absolute magnitude fainter than -16are considered to be dwarfs).

The mean value of the HI extent at half power (column 4) is about the Holmberg diameter (].08±O.23). At a level of ]019 cm-2 seven galaxies of this sample show an extended HI envelope. With the exception of these envelop~s the mean extent at this lower level is ].7 times the Holmberg dia meter (column 5). The distance-independent hydrogen mass to luminosity ratio (MHI!Lpg in column 6) is high, as expected for irregular galaxies (~.8 in Roberts' 1969 sample).

The maximum rotational velocity (Vmax in column 7) is derived from ,model calculations for galaxies not indicated by a cross (*). This value is smallest for the absolutely faintest object~ as expected. Motion in brighter irregular galaxies is dominated by rotation. For the faintest objects random motion is of the same order as ordered motion. For example the profiles for the Pegasus irregular (DDO 216) still show a small velo city gradient not greater than the velocity dispersion along the line of sight (column 8).

J. H. Seiradakis (La Jolla) and J. Materne (Berlin).participated in this work. - 66 -

Table 1

.,1 a) Galaxy M HI-diametera) HI- to ~I V Velocity DDO pg max at Holmberg L dis- 10 19 cm-2 diameter pg persian at

(arc min) HPW 10 19cm-2 (Mo/Le) (km g-1) (km S-1)

(I) (2) (3) (4) (5) (6) (7) (8)

223 -18.8 10 1.2 2.1 0.09 111 13 IelO -16. I 72* 1.1 7.2 0.57 >30* 8 -14.7 40* 0.9 1.7 0.40 >12b)* 13 -17.3 4 0.8 1.2 0.90 30 9 28 -18.6 5.7 1.1 1.5 0.10 217 9

29 -18.9 5.1 0.85 1.4 0.96 110 12 39 -18.0 6 0.92 1.8 0.67 158 8 42 -16.8 30* 0.9 3 1.0 57* 47 -16.1 5.8 1.4 1.8 0.94 49 14 50 -16.9 27* 0.7 2.4 1.16 58* 10

63 -14.5 5.4 1.3 2. 1.46 18.5 10 70 -14.5 6.5 1.2 1.7 0.74 42 12 236 -16.3 66* 1.3 3.9 1.65 56* 12 75 -14.8 54* 1.5 5.6 1.14 25* 10 125 -16.5 6. I 1.5 2.1 0.56 18.6 11

133 -16.2 4.4 0.7 1.05 0.60 71 9 161 -17.1 2.30 185 -17.0 5.5 0.86 1.7 0.46 75 18 216 -13.0 0.29 9.4 7 217 -17.7 1.09 221 -14.4 45* 1.3 3.6 0.74 23* 9 a) Model derived parameter with the exception of * b) Uncertain inclination - 67 -

ENVELOPES OF LATE-TYPE DWARF GALAXIES

JUrgen Materne Institut fUr Astronomie und Astrophysik der TU Berlin+) and European Southern Observatory

I. Introduction

A search was undertaken for neutral hydrogen (HI) in and around dwarf galaxies of late type. Among the objects were two pairs of dwarf galaxies in the NGC 4631 group, the NGC 4490/85 pair and ~he NGC 4618/25 pair, and a field gala xy, namely NGC 3109. A detailed discussion of the observa tions is given in W.K. Huchtmeier, J.H. Seiradakis, and J. Materne, Astron. & Astrophys., in press, where also the relevant literature is given. In the present paper a special aspect is emphazised which concerns the so-called missing mass in systems of galaxies and an indication is given where to look for it, eventually also with Space Telescope.

11. The HI Shell Around NGC 3109

The HI shell around NGC 3109 (DDO 236) was found to have an extend of 64 arcmin at a density of 1019 atoms/cm~ which is 4.4 times larger than the optical standard diame ter (D 25 = 18 arcmin). It is still 2.5 times larger than the diameter at 27 magi "which we measured to be 27.5 arcmin. This clearly indicates that we have to expect large gas shells around late-type dwarf galaxies. In .addition to the observed HI these shells could also contain additional large amounts of mass.

+) present address 68

Ill. The Shell Around NGC 4490/85

The two galaxies NGC 4485 and NGC 4490 have a separa tion of around 4 arcmin. There is a common shell of HI around the pai~ which has an angular size of 57 arcmin x 30 18 arcmin at a density of 5.10 atoms/cm2 . This corresponds roughly to a metric size of 200 kpc x 100 kpc if the dis tance is assumed to be 12 Mpc. In figure 1} the HI distribu tion is shown overlaid on a photograph of the pair.

In the integrated velocity distribution the two gala xies do not show up. However, the peak velocity of the gas cloud is between the velocities of the two galaxies and the range of the velocities spans the region of the velocities of the two galaxies as well. We interprete the cloud there fore as a common HI shell of the two galaxies. The whole cloud has a hydrogen mass of 1.8.1010 ~ . It is not clear G how much of the HI must be attributed to the individual galaxies.

The whole shell rotates around a position which is roughly centered on the galaxy NGC 4490. The peak velocities of the rotation curves are about 150 km/so The rotation curve comes down on both sides again (it is not a flat ro tation curve) .

IV. The Shell Around NGC 4618/25

The two galaxies NGC 4618 and NGC 4625 have a separa tion of roughly 8 arcmin. Their common shell of HI has a 19 2 diameter of 30 arcmin (again at 10 atoms/cm ). This cor responds to an extent of about 100 kpc (again at the assumed distance of 12 Mpc). In figure 2 an overlay of the HI dis tribution on a photograph of this pair is given. - 69 -

In this case we do see the two galaxies in the integra ted velocity distribution. But the hydrogen mass of the two 9 8 galaxies, which is 4.7.10 and 9.10 ~ respectivel~ is not e sufficient to explain the whole cloud. 3.108 ~ of HI are ~ presumably in the extended envelope. We could not detect any sensible rotation of the common shell of the NGC 4618/25 pair but we cannot exclude it either. v. Masses

We estimated the masses of the individual galaxies using a Brandt ~odel which gives reliable masses as fas as the observations reach. In the following the individual mas ses are tabulated:

NGC: 4485 4490 4618 4625 Mass: 0.6 2.4 3.3 0.7

An average over all galaxies gives a mean mass to light ra tio of 1.8 ~ /L , blue magnitudes of de Vaucouleurs are used. e e

The common HI shells indicate, however, that the pairs are bound. The above quoted masses are not sufficient to accomplish this. The mass of the NGC 4490/85 pair has to be increased by a factor of 5.1, the pair must have a mass of 10 15.10 ~ . The mass of the NGC 4618/25 pair ~as to be in- e creased by a factor of 3.3, therefore this pair will have a . 10 total mass of 13·10 ~. e

It should be mentioned as a side-note that these in creased masses are still not sufficient to bind the NGC 4631 group of which the two pairs are members. If one compares the total mass of the individual galaxies derived on the ba sis of Brandt models with the virial mas~ the discrepancy amounts to a factor of 16. - 70 -

,.,/ ...... 5* . Peak 1:.57 K kms-1/' - ",. / .\ . / \ 20 / • /. \ /- \" - " / \ I \ I 10' \ . / I \ I I I I o· I J /" f I / I· f ,- / -10' I / / I. / I / I- / -20' J / / T \ / \ ./ I • / \ / HPBW "'-. /" ...... - .-- ...... - 10' 0' -10' a*

J~ig~._ 1: 'l'hC' JIJ di slribution superimposed on a photograph of l!w NGe 4490/8') pai r of late-type dwarf galaxies. The lowest contour levels are at 5, 2, and 1 % of the peak intensity (457 K ~m/s) - 71 -

20' . Peak 132 Kkm 5-1 • -:- 0* - ':' '/ -- /. " • / " 10' / " \ ( \ \ . \ \ \ . \. \ 0' ".. , "- I .'\ • • I .\ j. "- .,,- / , / -10 ./ .. " ..,-/ " """"'------

20' 10' 0' -10'

fig. 2: The HI distribution superimposed on a photograph of the NGC 4618/25 pair of late-type dwarf galaxies. The lowest contour level is at 5 % of the peak in tensity (132 K km/s). 72

The rotation curve of the HI shell of the NGC 4490/18 pair indicates a somewhat higher mass for the extended system, but the increase is only marginal.

VI. Conclusion

We have two pairs of galaxies, NGC 4490/85 and NGC 4618/ 25, where we know that they have factors of 5 and 3 respec tively of hidden matter compared with the masses derived using Brandt models. We know furthermore that this mass is attached to the pairs themselves, making these objects very attractive candidates for the search of the so called "missing mass". The area to be surveyed is relatively small, only a few arcminutes squared. That makes the surroundings of these late-type dwarf galaxies an interesting target for Space Telescope~

Discussion

H. van Woerden: Materne's suggestion of a "missing mass" in the system shown in his first slide is unwarranted. The Effelsberg telescope resolves only two picture elements in the hydrogen distribution; hence one is hardly able to assign velocities to particular locations. And, especially, the mass in the system can only be derived from a velocity difference if one assumes circular orbits or an equilibrium configuration. The available observations would also allow for a galaxy pair in a non-closed orbit. 73

G. A. 'ra.mmann: There is a tendency to assume that low- .

luminosity dwarfs of the dIms class are particularly rich

in H I. This would be important for the search of what

Zwicky might have called the pygmy galctxies. Dr. Huchtmeier

shall tell us about some very faint dwarfs, which seem

to be very hydrogen-rich. However, it might be that we

know of these objects only just because of that property,

and it is not clear to me whether we can say with any certainty that the !YEical MHI/L ratio increases with. decreasing luminosity L. To illustrate my point I would

like to show you the following diagram, which contains data

that D. Jenni and I have compiled from the literature for a

(nearly) complete sample of galaxies of type Sc and later

I j I,Sr- • I I + 0 Se/SBe .. Sed i : • Sd/SBd r 0 Sdm L + Sm/SBm i ~ , lC + le Im/IBm ~ ) c: ~ I :::J lC ... \0' - Im Cl.OI) ..j .Q r l lit ~ 0 ~ I lC - Sm(087) , .5 • I .0 CD r 1 -'.... ~ Cl ---. Sdm (0.71) .j z + ! :E 00 0 00 ~ ~ .0• • - 5d (O.S4) I 05~ 0 • •! 0 - Sed (046) l • 0 . 0 0 0 - Se

(i.e., the corresponding galaxies from the Revised Shapley ~ l2~5 Ames Catalog with < and, for the case of -16 < MB < -14ID , the members of the Local Group and the M8l group).

The diagram shows the MHI/L ratios versus the absolute B magnitude M~'o, corrected for galactic and intrinsic absorp T tion. The overall impression of the diagram is that there is

lumin~sity. strong increase of MH/LB with decreasing But this increase is actually a consequence of mixing the types Sc, Scd, .•. Im. For each individual type, no such increase is manifested. The reality of the increase therefore depends entirely on the question, whether or not it is legitimate to consider Sc - Im galaxies as one continuous class. Dr. Huchtmeier informs me that the top witness for the hydrogen-rich "pygmies" is M8l dA with ~/LB ~ 18.

It is this galaxy, which I referred to in Session I, whose published apparent magnitude implied an impossibly low surface brightness. I believe the galaxy's magnitude must be

~ underestimated by at least 2 mag, and in that case, MH/LB 2.8. This value still leaves a wide gap between hydrogen- rich dwarf galaxies (dIms) and intergalactic HI clouds.

M. Dennefeld: There seems to be a trend in your diagram for the scatter in MH/L to increase for the later types. Would the larger scatter explain why one finds some exceptionally

H I-rich dIms?

G. A. Tammann: That is well possible. 75

G. Tenorio-Tagle: I wonder what fraction of the scatter could be due to interaction between galaxies.

G. A. Tarnmann: That is an interesting question, which could be tackled by plotting the relative hydrogen content in function of the distance to the nearest neighbour galaxy.

But not too much weight should be attributed to the scatter in the diagram, because our compilation of published HI masses was done rather simple-mindedly. A careful scrutiny of the available data and new observations of high quality

(especially for extended sources) could possibly bring a significant reduction of the scatter.

M. Dennefeld: The M~/LB versus ~ diagram brings another idea to mind. If one substituted the luminosities by the

~ ~, total mass and plotted MH/MT versus and if one still would find a slope in such a diagram, one could determine the minimum galaxy mass which is required to form stars.

This is because MH/M 1 is a natural limit. The experi- T = ment is probably not possible at the moment, because one would need reliable, true masses for a wide range of galaxies. - 77 -

Introductory report by J. LEQUEUX, Meudon Observatory I- Introduction

I would like to emphasize the interest of studies of the Magella nic Clouds (MC) on three points

1) They are the closest galaxies of all, thus they offer unique opportunities for studying intrinsically faint stars outside our own Galaxy.

2) They can be considered as prototypes for more remote irregular galaxies. The Space Telescope will reach very faint stars in the latter, the equivalent of which can be observed from the ground in the Mcs.

3) They are less evolved than our Galaxy, thus are closer to primordial matter in effe~t, the abundance of heavy elements is smaller by a factor I.B in the LMC and by a factor 5.B in the SMC with respect to the Orion Nebula. Thus many properties are expected to be different.

11 - Studies of Interstellar }futter

1. HI! regions

HII regions in the MC have been extensively studied with modern means during the last years. All authors essentially agree on the values of the abundance of He, N, 0 and a few other elements, all of which are deficient with respect to the solar neighborhood (see e~. Lequeux et a1., 19BO).

2. Interstellar lines

The study of interstellar lines, particularly in the far-UV, will allow to obtain the abundance of other elements and to study the physics of interstellar matter in the MCs. Some IUE observations are available for stars in the LMC and in the SMC (see e.g. de Boer and Savage, 19BO ; Prevot et aI, 19BO) in the high resolution mode, the MC lines are well separated from the galactic ones. These observations have not yet been used for obtaining abundances ; this will be diffi cult because the column density of hydrogen is not available, since no observation of interstellar Lyman a exists. We probably will have - 80 - to wait for observations with the High-Resolution Spectrograph attached to the Space Telescope to obtain really s~ificant abundances. Some important atoms ions and molecules (D, H , CO, OVI, and to some extent 2 C, Nand 0) can only be studied at wavelengths shorter than ~ 1200 A, beyond the range of the ST. Several French laboratories are joining with Princeton and Lockheed to propose to NASA a specialized Shuttle high resolution far-DV spectrograph, the MISIG project, which would fill this gap. For the interpretation of these data, it is crucial to have very high resolution profiles of optical interstellar lines. The observations by Blades (1980) are a major step in this direction and I hope that ESO will contribute to such programs.

3. Dust extinction law

We now know that the DV extinction law differs in the LMC and in the Galaxy (Nandy et al., 1980, Koornneef, 1980). We are presently studying the DV extinction law in the SMC, which is not yet known, also using lUE observations. Presumably the dust/gas ratio is different in the MCs and in the Galaxy. For such studies, complementary ground-based observations should give better E(B-V) for the studied stars, the column density of HI in their directions (high-resolution 21-cm obser vations) and better determinations of the galactic extinction towards the MCs.

III - Studies on Stellar Populations

1. Faint stars The main goal for the studies of faint stars in the MC is to see how this population is affected by metallicity. For example one may wish to study the field luminosity function (preliminary results have been published by Butcher, 1977), the lower main-sequence in clusters, horizontal branch stars etc. Such studies will be pursued with the Space Telescope down to magnitudes 27-2~ but there is already a large amount of work to do with ground-based telescopes down to m = 22-24. The main problem is the lack of photometric standards at such faint magnitudes, the only published ones being stars in NGC 1866 down to V = 22.4 (Walker, 1974). This needs conventional photoelectric photometry for the brighter stars, 81 extended by electronographic photometry at fainter magnitudes. Such a programme is presently conducted with the }lc}fullan camera at the I.Sm Danish telescope by A. Ardeberg.

2. Brighter stars

These h~vcbeen studied much more in a recent past. It now appears that not only are there photometric and spectroscopic differences between the supergiants in the MCs and the Galaxy (see e.g. Dubois, 1979), but that also in spite of gross similarities (Humphreys and Davidson, 1979 Vangioni-Flam et al., 1980)} some differences in the population of the upper HR diagram exist. The most striking effect is the huge variation in the ratio of the number of red supergiants to Wolf-Rayet stars between the inner Galaxy, the outer Galaxy, the LMC and the SMC : this is interpreted as being due to the effect of mass loss on post-main - sequence evolution of massive stars, the mass loss rate decreasing with decreasing metallicity (Maeder et al. 1980). We may also expect differences in the ratio red supergiants/blue supergiants and perhaps also in the upper-luminosity function of O-B associations, due to the effect of metallicity. Much more ground-based work is needed, in particular a complete survey of red supergiants in the SMC. The importance of such studies is reinforced by the fact that they are preliminary to similar ones with the Space Telescope in more distant irregular galaxies with different metal contents.

IV - Global or large-scale parameters

It is very important to obtain good large-scale parameters fnr the MC such as photometry in various bands, in order to understand ~heir evolution in comparison with more distant galaxies (see e.g. Vangioni Flam et al., 1980). Surprisingly enough, we have only the old B,V,R isophotometry of de Vaucouleurs (1960), and photometry in 3 far-UV bands (Maucherat-Joubert et al. 1980). Photometry in other bands, par ticularly in U, in the near and in the far IR are badly needed, as well as higher resolution isophotometry in various bands (this is in progress for the far UV, by Carruthers). It should also be noted that no CO survey exists for the MCs ; we hope that this gap will be filled by the 4-m Australian millimeter radiotelescope or by another survey instrument in the Southern Hemisphere. 82

REFERENCES -

Blades,J.C., 1980, Monthly Not. Royal Astr. Soc., 190, 33 Butcher, H., 1977, Astrophys. J, 216, 372 de Boer, K.S., Savage, B.D., 1980, Astrophys. J., 238, 86 Dubois, P., 1979, Astron. Astrophys., 79, 143, and Ph D thesis, University of Strasbourg, 1980 Humphreys, R.M., Davidson, K., 1979, Astrophys. J., 232, 409 Koornneef, J., 1980, preprint Lequeux, J., Peimbert, M., Rayo, J.F., Serrano, A., Torres-Peimbert, S., 1979, Astron. Astrophys., 80, 155 Maeder, A., Lequeux, J., Azzopardi , M., 1980, Astron. Astrophys., in press ~~ucherat-Joubert, M., Lequeux, J., Rocca-Volmerange, B., 1980, Astron. Astrophys., 86, 299 Nandy, K., Morgan, D.H., willis, A.J., Wilson, R., Gondhalekar, P.M., Houziaux, L., 1980, Nature, 283, 725 Prevot, L., Laurent, C., Paul, J., Vidal-Madjar, A., Audouze, J., Ferlet,R., Lequeux, J., Maucherat-Joubert, M.L., Prevot-Burnichon, M.L., Rocca-Volmerange, M.L. 1980, Astron. Astrophys., in press. Vangioni-Flam, E., Lequeux, J., ~~ucherat-Joubert, M., Rocca-Volmerange,B., 1980, Astron. Astrophys., in press Vaucouleurs, G. de, 1960, Astrophys. J., ~, 574 Walker, M.F., 1974, Monthly Not. Royal Astr. Soc., 169, 199 STUDIES OF THE LARGE ~lAGELLAKIC CLOUD

J. V. Feitzinger, Astronomisches Institut, Ruhr-universitat Bochum

The kinematics and dynamics of the Large Magellanic Cloud (L!'lC) are inYestigated. The main structural features of the LMC are disk, bar, rudimentary and well-developed stel lar arms as well as spiral filaments (not necessarily con nected with density waves); the r-structure is a broken-up ring structure. Embedded into these features are young, asymmetrically located spiral arm filaments. As an explana tion for these structures stochastic star format~on in an ordered chain reaction is proposed. The pattern of the spiral arm filaments is determined by the rotation curve. The mor phological peculiarities of the LMC can also be detected in other galaxies of that type.

The mean absolute displacement of the centers of bar and disk, determined from 18 galaxies, is.A = 800 pc. The displacement between the bar center and the symmetry center of the rota tion curve is of the same order.

The presently known radial velocities of planetary nebulae, star clusters, HI and HII regions.and stars belonging to the Lt-1C have been collected in a catalogue as the basis of a discussion about the kinematics and dynamics of the LMC. Con trary to earlier work, we have used, for the first time, the radial velocities of objects of all subgroups together with a proper weighting scheme. Thus the basic kinematics and dynamics of the LMC has been deduced.

The radial velocity field shows no central symmetry; it is characterized by large-scale (2-J kpc) disturbances. By comparison with the velocity field of other galaxie~ three main disturbances are identified: an oval distortion of the velocity field in the bar region, a radial velocity 84 field around JO Doradus, and disturbances connected with a warp or material above the disk in the southern quadrants. The results of a detailed numerical analysis of these three facts can be summed up as follows:

(1) The rotation curve is determined over a 100 diameter; it shows differential rotation, an asymmetric behaviour in the south and a double structure in its HI component. The rotation center is displaced by 0?7 from the bar center. The orientation of the kinematic line of nodes and the systemic velocity vary as functions of the distance from the centui". Therefore, it is possible to show that large scale disturbances (warping, z-structure and streaming~ motions) definitely do exist.

(2) By variation of the kinematical parameters (systemic velocity, inclination, orientation of the line of nodes, rotation center) the dispersion of the measured radial velocities was minimized and the basic rotation curve determined. The rotation curves for the north and south side of the LMC are significantly different. The south side is either warped or there is material above the main plane. There seems to be a connection between this structure, the Panmagellanic Gas and the Magellanic Stream. The north side appears to be free of distortion.

(J) The residual velocity field (observed minus model1 de duced from a basic rotation curve, shows that the displace ment between the rotation center and the bar center is not caused by local streaming motions. The rotation center must be the mass center. The bar shows a radial velocity field; in the JO Doradus region,inward and outward motions are found.

The mean velocity dispersion of population I objects is 10.5 km/s, of population 11 objects 16.0 km/so Red and blue globular clusters show different kinematical behaviour. 85

In a comparison of eight mass models, taking into conside ration thickness effect~ and controlled by surface photo metric data, the mass of the LMC is found to be (0.5 ± 0.1) 0 1010~Q (assuming an inclination of 33 , a systemic velo city of 46.9 km/s, and a distance of 56 kpc). Dynamically, the LMC can be described by a dominating disk potential with an additional bar potential as a disturbance. The mass of the bar is 0.6 109~@. The stable neutral point of such a configuration can be found in the residual velo city field. The absorption feature crossing the bar co incides with the maximum velocity gradient of the computed radial velocity field in the plane of the system. - 87 -

COLOUR-MAGNITUDE DIAGRAMS AND LUMINOSITY FUNCTIONS IN THE MAGELLANIC CLOUDS,

A. Ardeberg, P. Linde, H. Lindgren and G. Lyngg, European Southern Observatory and Lund Observatory

Summary

Colour-magnitude diagrams and luminosity functions are studied for fields in the Magellanic Clouds. Electronography is used. For ground-based observations the main sources of error are sky background and image crowding. Both types of errors can be vastly decreased, if observations from space are made.

Key words: colour-magnitude diagrams - luminosity function Magellanic Clouds - electronography - image crowding.

Introduction

The colour-magnitude diagrams and luminosity functions are studied in selected fields in the Magellanic Clouds. The fields now include the old halo population of both galaxies and the young population in the Wing of the Small Magellanic Cloud. A Spectracon electronographic camera (l1cGee, 1976; Walker, 1976) has been used on the ESO 3.6-m and 1.S-m tele scopes. Ilford L4 nuclear track emulsion is used. Exposures of up to 2 hours are made in V, B and, to a limited extent, U. 88 flcasurements are made with a PDS machine at the Lund Observ- ator)' (Elvius et al., 1978; Lindgren et al., 1979). The pixel size corresponds to approximately 0.14 x 0.14 arcsec 2 • Photoelectric calibration is made with the ESO 3.6-m telescope, using several stars within each ~rogramme field. The brightness range of these calibration stars is roughly V= 13 to V= 20. Further, electronographic exposures have been made of some galactic clusters with existinp, photoelectric photo metry. This gives us an additional check of the linearity of the electronographic measurement data. Our new observations are made with a ~tdlullan electronographic camera (Hartley and McNullan, 1976; Hcf\1ullan et aI., 1976; Hdtullan and Pm-Jell, 1979). The larger target area of this camera is advantageous

esneciallv.. ~ for the statistics of brighterl,..' stars.

Preliminary Results

Preliminary results exist for an U'C-halo field situated 1200 pc from the Bar. There are quite a number of red stars, es~ecially among the brightest objects. There appears to be a gap in the distribution of red stars between Mv = -1 and Hv O. This has to be confirmed before any interpretation is attempted. In any case, our results are quite different from those of Butcher (1977), who investigated a field in the neighbourhood of the cluster NGC 1866. - 89

Sources of Error

In principle, there are three main sources of error af fecting the information contained in the resulting clectrono graphic images. These errors are detector noise, sky back ground and image crowding. Test observations with closed camera shutter have proved detector noise to be virtually negligible, even for fainter objects, if target and emulsion are properly selected. It is noted that ~reat care has to be exercised in order to minimise influence of emulsion flaws and extraneous particles.

The sky background is an error source which has to be treated with considerable attention. If this is not the case, results for the faintest stars can be completely spurious.

As an example, on a 1-hour B exposure taken at the 3.6-~ telescope with a seeing slightly worse than 1 arcsec, the outermost part of faint stellar images covers a diameter of about 160 ~m. Inside such an image, dark sky contributes around 3 x 10 6 photons. Therefore, sky background defines a theoretical limiting magnitude slightly beyond B = 25, if a standard deviation of 5% is accepted.

Still, the worst threat to faint-star photometry in the

~1agellanic Clouds comes from image crowding. Certainly, these effects can be diminished if a proper instrument is selected and exposures are made during p,ood seeing conditions.

The ESO 3.6-m telescope provides good images. However, also for exposures taken during good seeing, around 1 arcsec, 90 - image crowding defines the limiting magnitude obtainable with acceptable errors. On the one hand, a linear detector seems mandatory, if statistically useful information is the aim also for fainter stars. On the other hand, also with linear devices, it is in practice impossible to avoid considerably increased instrument errors when complexes of several images have to be resolved. This is, of course, especially true for the fainter members of such comlllexes.

We have paid considerable attention to the disentangling of overlapping images (Lindgren et al., 1980). Still, even for best-quality exposures with the 3.6-m telescope in the direction of well-defined halo areas of the Large Hagellanic Cloud, there seems to be a definite limit beyond which crowd ing of the outer profiles of neighbouring brighter objects rapidly increases measurement errors. Also if we set the "acceptable standard-deviation limit" as liberal as 0.2 mag nitudes, the corresponding limiting magnitude can hardly be pushed fainter than V = 24. As a result, comparisons with solar-neighbourhood data will be considerably weaker than de sirable. Further, V = 24 is just beyond the magnitude for which we expect to find the knee of the luminosity function of the Large Nagellanic Cloud.

A further source of error may well be due to the measurement procedure. Possible non-linearity effects in the response of the measuring machine have to be carefully checked. 91

Need for Space Observations

It seems to us that the only feasible way to deeper

photometry is through space observations. From such observ

ations we expect image diameters to decrease dramatically,

more than a factor 10 for diffraction limitation.

Space observations will improve our possibilities to

reach faint objects in two ways. First, the sky background

gets generally lower and its influence is even more notably

decreased by the decrease of image diameters. Second, and

still more important, the problem of image crowdinp, will be

vastly diminished. This being the case, we would at once

come far beyond the solar-type stars in the Large Hagellanic

Cloud. This would allow some exciting comparisons between

the evolution of this galaxy and that defined by the stars in

the solar neighbourhood.

We hope soon to be able to present data down to V = 23ffl m 24 for a number of fields in the t1agellanic Clouds. We

also hope that space observations can be made for the most

interesting of these fields and for some of the local dwarf

galaxies.

References

Butcher, H. 1977: Astrophys. J. 216, 372.

Elvius, T., Lindgren, H., Lyngfi, G. and Wihlborg, N. 1978:

Reports from the Observatory of Lund N~ 14. 92 -

Hartley, K. F. and HcMullan, D. 1976: Adv. E. E. P. 40A, 493. Lindgren, H., Ardeberg, A., Linde, P. and Lyneft, G. 1980: Proc. ESO Workshop on Two-Dimensional Photometry, Noordwijkerhout, 21-23 November 1979, eds. P. Crane and K. Kj1ir, p. '155. Lindgren, H., Lyngft, G., Linde, P., and Hagerbo, H. o. 1979: Proc. International Workshop on Image Processing in Astronomy, Trieste, 4-8 June 1979, eds. G. Sedmak, M. Capaccioli and R. J. AlIen, p. 65.

McGee, J. D. 1976: IAU Colloquium N~ 40, Applications Astro nomiques des Recepteurs d'Images a Reponse Lineaire, eds. M. Duchesne and G. Lelievre, p. 4-1. McMullan, D., Powell, J. R. and Curtis, N. A. 1976: Adv. E. E.

P. 40B, 627.

McMullan, D. and Powell, J. R. 1979: Adv. E. E. P. ~, 315.

Walker, M. F. 1976: IAU Colloquium N~ 40, Applications Astro nomiques des Recepteurs d'Images a Reponse Lineaire, eds. M. Duchesne and G. Lelievre, p. 37-1. - 93 -

SESSION IV

(Chairman: D. Kunth)

INTERGALACTIC HII REGIONS - 95 -

"INTERGALACTIC HII REGIONS" D. Kunth European Southern Observatory

Introduction

It is not 100% clear to me what the title of this section means, for the objects I am supposed to talk about were originally called "Extragalactic HII Regions" (Searle and Sargent, 1972) and are also discussed in the literature under the pseudonYm "blue compact galaxies".

In fact, galaxies with 'HII region like' spectra are fairly common, they constitute a large fraction of those in the Markarian, Tololo and the Zwicky lists. Among them, the low luminosity galaxies are of particular interest, because they are very unevolved objects with a low abundance of heavy elements.

They are ideal for studying star formation in a primordial gas and for the determination of the primordial helium. They also exhibit a simple structure and their low mass implies a small gravitational potential; many are isolated. They are therefore reference objects on which galactic evolution models can be tested. In table 1, I give the essential properties of these objects. - 96 -

Table 1.

PROFILE

ARCHETYPES II Zw 40, 11 Zw 70, I Zw 18

Distances > 10 Mpc

Size of emitting region 0.2 to 1 kpc

Holmberg diameter 8 kpc

Colours . U-B -0.4 ., -0.75

B-V 0.4 0.0 8 Mgas 1 to 4 x 10 M (i) 9 M total .8 to 1.5 x 10 M (i) Mgas/Mtotal .2 to .4

Mtotal/L 1 B Z 0.0033 to 0.0004

Y 0.250 to 0.233 -1 Velocity dispersion 30 km sec (optical, HI)

Environment interacting, isolated - 97 -

Up to now, two possible models have been put forward to account for their HI content, colour and metal content. Either star formation only started recently (lIyoung , ll galaxies ) ,or we are just observing galaxies whose present rate of star formation exceeds their average rate in the past. I would like to spend some time on the following questions relating to the nature of these objects, which should outline their astrophysical importance for galactic evolution:

1. Do the blue compact galaxies differ from the Irregulars? Recent work by Lequeux et al (1979) and Balkowsky et al (1978) support the view that they are similar and define similar relations, as far as global properties are concerned, such as Mtotal, Mgas/ Mtotal and Z. They hav~however, a bluer colour thQn average and a much higher surface brightness. They are also mainly found at larger distances than the known, well studied Irregulars, which may suggest a selection effect: one has to explore larger volumes to find objects in which the star formation rate is currently higher and which started 7 only 10 years ago. Clearly these galaxies should be studied in connection with closer systems like the SMC, ICIO etc, to establish their relationship.

2. Which mechanism triggers star formation in these galaxies? To give a complete• description of the star formation activity in these galaxies, one should know the relevant ingredients which enter the recipe : low metallicity, a core embedded in large neutral hydrogen clouds and velocity dispersion of the order of those found in giant HII regions (Terlevich and Melnick) are some of them. - 98 -

It is not clear whether some are relevant : existence of high luminosity galaxies? Tidal interaction effects? It is clear that some are still unknown, as suggested by the following points:

3. How can one describe the Stellar birthrate? a - are more massive stars formed in underabundant systems? Relation between metallicity and IMF (Larson and Tfnsley 1978) or metallicity and density (Burki 1977) b - are massive stars the real and only ionizing source? c - if any, how many bursts do such galaxies undergo? What are the quiescent counter parts? Studies of individual HII regions should provide ways of measuring their velocity differences, their chemical composition differences, and to estimate parameters of interest for galactic evolution models, such as the mixing time scale of newly synthesized elements in the gas.

The following list is not intended to be exhaustive but to suggest possible projects which could be pursued from the ground or carried out with the ST. The chief advantage of the ST for this topic is its:

0.1 11 resolution

UV) spectral ranges IR) 1. Are there objects with more extreme properties? Surveys from the ground involving searches for HI clouds and "extragalactic HII regions" of even lower luminosity should be encouraged. 99

The identification of astrophysical objects undergoing a quiescent or milder phase should be carried out from the ground. A statistical approach may help to identify the right kind of objects (Thuan and Seitzer 1979) as well as extensive and detailed studies of low surface brightness dwarf galaxies.

2. Study of the present burst. There is a chance that the ST can help to establish the IMF by direct observations of the ionizing stars. The FOS will be useful to obtain ultraviolet spectra so that properties of the photonionizing stars can be derived. The effect of the dust on the ionizing radiation could be directly measured by observing the bump at 2200 ~.

The wide field camera would also allow measure ments of the extension of the emissive nebulosity around knots and isolate locations of star form ation outside the nucleus. Youth hypothesis: an interesting suggestion has been made by Pagel at the last ESO Conference on the Space Telescope (1979) and involves 'accurate photometry down to M > 5 below the main sequence turnoff. This is v obviously out of reach for the blue compact galaxies, but results gathered on the SMC would be quite informative. A smooth underlying old population could be detected using the WFC, but this could also be advantageously carried out from the ground. Spectroscopy at longer wave lengths - up to 11000 R- could reveal absorption lines typical of an old main sequence population.. - 100 -

More challenging would be the attempt to detect globular clusters. Since globular clusters are barely resolved at the distance of M31 one could hope to resolve them at distances as large as 7 to 10 Mpc. Otherwise UBVRI photometry could be used to distinguish them from field stars. Since one expects their luminosity to be of the order of -8, low resolution spectroscopy could still be obtained out of distance moduli of typically 30.

3. From an empirical point of view, the study of the metal content in these galaxies and particularly the helium content is of specific interest (Searle and Sargent, 1972). Such a project, in my opinion, is within the reach of ground based telescopes. There are, however, two areas where the ST could lead to a breakthrough: a - It will be possible to determine abundances in the nebulae outside the nucleus which will cons equently be free from underlying stellar components. This possibility arises because the resolution will enabl~ us to isolate stellar complexes from the emissive nebulosities. b - A major key to the chemical evolution is the C production by stellar nucleosyntheses. The ST offers a unique possibility to determine the carbon abundance in unevolved objects. That such a project is feasible is attested to the preliminary IUE observations of carbon lines such as CIII 1909 and CIV 1549 in planetary nebulae of respectively low and high excitation such as IC 418 and NGC 7662. From the ground, only the recombination line CrI 4267 has been reported in Orion by Peimbert (1979) and is 400 t~es weaker than HB. - 101 -

References

Ba1kowski, C., Chamaraux, P., We1iachew, L. : 1978, Astronomy and Astrophysics ~, 43 Burki, G. 1977, Astron. Astrophys., ~, 135. Larson, R.B., Tinsley, B.M.: 1978, Astrophys. J. 219,46 Lequeux, J., Peimbert, M., Rayo, J.F., Serrano, A. and Torres-Peimbert S., 1979, Astron. Astrophys. 80, 155 . - Page1, B.E.J. : 1979, in F. Machetto, F. Pacini and M. Tarenghi (eds.), "Astronomical Uses of the Space Telescope", pp. 229-239. Peimbert, M. : 1979, in W.B. Burton (ed.), "The Large Scale characteristics of the Galaxy" pp 307-316 Sear1e, L. and Sargent, W.L.W. : 1972, Astrophys. J. 173, pp. 25-33 Ter1evich, R. and Me1nick, J. 1980,ESO Preprint no. 95 Thuan, T.X. and Seitzer, P.O. 1979, Astrophys.J. 231,327 - 103 -

• Astrophysics of galactic and extragalactic HII regions with IUE: a preliminary discussion of the carbon abundance and the absorption line spectrum.

S. D'Odorico, European Southern Observatory

P. Patriarchi and M. Perinotto, Arcetri Observatory, Florence, Italy

SUMMARY

The carbon abundances in HII regions of the SMC, M33 and MlOl have been

compared with the solar and Orion values. The clo ratio is constant in systems

of quite different metal content with the possible exception of NGC 604 in M33.

The low dispersion lUE spectra of HII regions in spiral and irregular galaxies

are characterized by absorption lines of interstellar, cold gas with equivalent

widthslarger than 1.5 R• The P Cygni profiles in the spectrum of NGC 604 suggest

that WR type stars dominate the UV spectrum of this H II region.

I. INTRODUCTION

Several galactic and extragalactic HII regions have been observed with the lUE

low dispersion camera. The large entrance aperture of the spectrograph (elliptical

in shape and about 10x23 arc sec) is well suited to sample selected areas of ga

lactic and Magellanic Cloud HII regions and to include essentially all emission

from objects beyond the Local Group. The spectra are in general of very good

quality in the short wavelength range and contain a large amount of astrophysical

information on the stars which ionize the HII region, on the absorption and

scattering properties of the associated dust, on the chemical composition and

physical conditions of the gas. unfortunately, this information is often

blended together in a way that one cannot disentangle the physical parameters on

the basis of UV data alone. A typical example: the shape of the continuum in the

far UV is determined both by the spectral energy of the associated stars and by - 104 - the influence of the dust. For both effects, astronomers often rely on the extra~ polation of the properties of galactic object~neglectingthe possible effects of the different chemical compositions.

An ad hoc IlPdel for each case and data, at all wavelengths, will be needed to reach some significant conclusions. In this paper, we use the UV spectral data for 5

HII regions in different galaxies to discuss preliminarily the carbon abundance and the absorption line spectrum.

The subject is relevant to this workshop because the giant HII regions which are observed in external galaxies resemble the dwarf HI rich systems (intergalactic

HI! regions) which have been discussed here: in both the bulk of the energy is radiated in the UV by the hot, massive stars which ionize them. Experience with lUE data on extragalactic HII regions will therefore be essential to plan observations of dwarf galaxies with ST.

II. The Carbon abundance from the UV emission lines

IUE spectra of galactic and extragalactic HII regions are usually dominated by a strong continuum which raises toward the camera limit (llOO~), with direct or scattered light from hot stars and atomic continuum of H and He as the main contributions in this spectral region. The emission lines from the ionized gas are not as prominent as in planetary nebulae or nuclei of Seyfert galaxie~because of the lower degree of ionization and temperature. It is difficult to measure La

in emissio~because it is blended to the geocoronal La in systems of low redshift

(as is the case for the regions we are discussing here) and partially scattered by intervening neutral gas. Among elements other than H, emi~sion lines from relevant

ions of carbon are represented in the 1000-3000 ~ spectral range and provide

the best means to determine the carbon abundance, because of the problem with the

interpretation of the faint optical lines of c+ and c++ (e.g. Seaton 1980). In

Table 1, we give the carbon abundances for HII regions in four galaxies and compare them with nitrogen and oxygen abundance. Sources of the data and assumptions for - 105 -

the computation of the abundance are given in the notes.

Table 1

Object Galaxy C N/O C/O 0 Notes 4 xl04 xl0

Sun The Galaxy 4.7 0.12 0.56 8.3 1

Orion The Galaxy 2.1 .10 0.42 5.2 2

NGC346, IC1644 SMC 0.6 0.04 0.51 1.1 3

NGC 5471 MlOl .5-1 0.05 0.36-0.7 1.4 4

NGC 604 M33 <0.5 0.05 <0.1 5.3 5

Notes

1. Lambert 1978

2. In Orion, Ai909 of CIII} and A2325 of CII] are observed with lUE. Carbon

abundance in Table 1 is an average of values given by Perinotto and Patriarchi

(1980) and Torres-Peimbert, Peimbert and Deltabuit (1980). Nitrogen and Oxygen

abundances are from Peimbert and Torres-Peimbert (1977).

3. Data are from the work by Dufour, Talbot and Shields (1980). For carbon, He

abundance is based on the strength of the A1909 line. Physical conditions

and abundances of Nand 0 are based 'on groundbased observations by the same

authors.

4. The HII region, the brightest in MIOl, was observed by Rosa (1980) who measures 12 -2-1 a value of 1.4 10- erg cm s for the A1909 ~mission. Lines from CIV or

CII are not observed. To correct for reddening, we used ~he UV extinction curve

of Nandy et al (1980) for 30 Dor (E(A-V)/E(B-V) = 6.5 at A1908A) and CH a = 0.18 from Shield and Searle (1978). The reddening corrected value of A1909 -12 -2-1 CIIII is then 4.1 10 erg cm s • Since NGC 5471 enters completely into 2+ the lUE slit, the C abundance has been derived by comparison with the inte-

gral HB flux,as given by Israel and Kennicutt (1980) (1.8 x 10-12 erg cm-2s -1 - 106 -

2+ 0 after reddening correction). With the T of the 0 region, taken as 14000 K from e Smith (1975) and the equations and atomic data given by Perinotto and Patriarchi -5 2+ (1980), an abundance of 5 x 10 for C is derived. At this temperature the + amount OD C is negligible (as confirmed by the absence of A232~which in Orion 3+ is as strong as A1909) while C could represent a significant fraction of the

0 total carbon abundance (see the results for 0.2 and T > 13000 K in the z/zo = e -4 grid of models by Stasinska 1980). We estimate lxlO as an upper limit for the total carbon abundance. The oxygen and nitrogen values are from Smith (1975).

Note that Shield and Searle (1978) measure a fainter A4363 of (OII~ and hence derive T 11.900 OK. The oxygen abundance has the same value as in Smith (1975) e = 2 but we would obtain a C +abundance three times larger. On the other hand, the 3 amount of C + at this lower temperature would be reduced so that the change in the

total carbon abundance would be less than a factor of 2.

5. A tracing of the SWP 7349 of NGC 604 is presented in Fig. 2. By assuming A v 0.3 for NGC 604 and the same reddening curve as for NGC 5471, we estimate

an upper limit for the reddening corrected flux of the A1909 CIII line of -13 -2 -1 2.3 xlO erg cm s For NGC 604 the IUE entrance slit was centered on the o brightest section of the nebula at a position angle of 327 (Rosa 1980). It seems appropriate to compare the UV measurement with the H8 flux given by Searle (1971', which was measured at the same position and through a similar aperture. This -12 2 gives FH8 = 3.7 x 10 erg/sec cm after a correction of a factor of 1.7 for the difference in aperture areas and for reddening. Adopting a T 9100 OK from e = 2 -5 Smith (1975), we derive an upper limit for the ratio of c +/H = 3.5 10 • At the

ionization stage of NGC 604, we can neglect C,3+ and C+ has to be less than 30%

(the Orion value). The final upper limit for the total C content is then fixed 5 at 5 x 10- •

Values for Nand 0 abundances in Table 1 come from Smith (1975). - 107 -

There are several sources of errors on the carbon abundances given in Table l.

Apart from the uncertainties in the lUE observations and calibration, the results are strongly dependent on the reddening correction and the electron temperature.

An estimate of the uncertainty is 0.3 in the log; smaller errors are quoted for

the Orion and SMC data by the respective authors.

It is now possible to compare the carbon abundances versus the metal content

(as represented by 0) in four galaxies.

The carbon abundance is correlated to the oxygen abundance in all objects except in

NGC 604, where carbon is strongly underabundant. The uniformity of the c/o ratio over a range of Z, supports the prediction, by the theoretical models,that C nucleosynthesis follows closely that of 0, Ne, Sand Ar (Arnett, 1978).

Several hypotliesis could be made for the carbon deficiency in NGC 604, but we feel that this result has to be confirmed by observations at other positions in the nebula before attempting any interpretation.

lIt. The UV absorption line spectrum of extragalactic HII regions

We discuss briefly in this chapter the low dispersion absorption line spectra of the extragalactic HII regions, NGC 604 in M33, N63 in the LMC and NGC 5471 in

MlOl to provide a preliminary reference for the interpretation of UV observations of other extragalactic systems.

Fig. la shows the AAllOO-2000 ~ spectrum of the HII region N63 in the LMC from the IUE SWP 3490. The low ionization interstellar lines of SII 1257+ Sill 1263,

SiII+OI A1305 and CII A1335 are well detected with equivalent widths of about

1.5 ~. At our resolution, the absorption by galactic and LMC interstellar gas are blended together and such intensities are not unexpected (Savage and de Boer 1979, de Boer and Savage 1980). Of the high ionization stages, SiIV A1402 is found to be as strong as the low ionization lines, but CIV appears unexpectedly weak.

The HII regions also show emission at I.. 1909 of CIII] and two unidentified lines which app~l, at 1746 and 1765 i. Tracing "b" from the same lUE spectrum refers - 108 -

to the closeby supernova remnant N63A, discussed by Benvenuti et al (1980).

Resonance UV lines become prominent due to the high temperature reached in the

shock ionized gas. The absorption by high ionization stages, if present, would

?e filled by emission and can't be detected. Low ionization absorption lines are

poorly seen because of the low continuum, but their EW are similar to those ob

served in the HII region spectrum, as expected from their interstellar origin. l In Fig. lc, the spectrum of the 06p star e Oric is shown for comparison.

The short wavelength spectrum of the Orion nebula is essentially the same

(Perinotto and Patriarchi, 1980), but for the presence of the A1909 i line of

CIII] in emission. This spectrum is saturated between Al240-l400 R, but the inter-

stellar lines are seen in shorter exposures in this spectral range. Fig. 2 shows the integral spectrum of NGC 604 obtained by Rosa (1980), who mainly

discussed the spectral energy distribution, and two tracings of different

regions along the slit. Variations in the absorption line strengths are observed

in co:incidence with discrete objects in the HII region. unfortunately, the position

of the slit is not known with sufficient accuracy to identify the individual

stars or star associations.

Tracing a in Fig. 2 shows all the interstellar low ionization lines observed in

NGC 604, with equivalent widths, of about 2 ~. Again, a strong contribution from

interstellar gas in M33 is blended to absorption by galactic gas. At

position a, the slit intersects a group of bright stars (Rosa, 1980 and Benvenuti

et al. 1979). We observe strong absorption of N V, 1240 i, Si IV 1398 ~, C IV 1549 ~

with marked P Cygni profiles, which suggest the presence of a WR or WN star.

In this spectrum, strong absorption lines of Fell 1608 ~, NIV 1719, Al III 1857 i

and an emission feature at 1640 R (Hell) are also probably of stellar origin.

At position b, the low ionization lines have the same equivalent width as in a,

but the Fell 1608, Si IV and C IV are weak if present at all.

Finally we briefly comment on the spectrum of NGC 5471 in MlOl reproduced by - 109 -

Rosa (1980). The interstellar lines of atoms in low ionization stages are present with an intensity comparable to that observed in M33 and the LMC, even if the galactic latitude of M101 is higher (60° versus 30°). Lines of Si IV,

CIV and NIV have equivalent widths similar to those of the low ionization stages and do not show, at the resolution of the low dispersion camera of

IUE, the PCygni effect. As a conclusion, we remark that strong absorption lines (EW>1.5 2) by Sill and CII are to be found in the UV spectra of objects associated to spiral and Mage11anic Irregular galaxies. For redshifts smaller than 2000 km/sec, the contributions of galactic and local gas will be unresolved with the lUE low dispersion camera. The presence of hot gas in a galactic and LMC corona has been suggested by Savage and de Boer (1979) and de Boer and Savage (1980). In our objects the absorption lines from high ionization stages, like NV, SiIV and CIV are observed with EW smaller or equal to the lines of the low ionization stages and in some cases are contaminated by absorption due to mass loss from eruptive stars. This is clearly the case in a section of NGC 604 as revealed by the strong PCygni profiles in the high ionization lines. In such a case the peculiar star (or stars) dominates the spectral energy distribution so that a discussion of the global stellar population in that HII region based on the UV continuum is not possible.

We are grateful to M. Rosa and R. Dufour for communicating their results in advance of publication. - 110 -

REFERENCES

Arnett, W.D. 1978, Ap.J. 219, 1008

Benvenuti, P., D'Odorico, G. and- Dumoute1, M., 1979, Astrophysics and Space Science, 66,39

Benvenuti, P., Dopita, M.A., and D'Odorico, S., 1980, Ap.J. 238

de Boer, K.S. and Savage, B.D., 1980, Ap.J. 238, 86

Dufour, R.J., Ta1bot, R.J., Shields, G.A. 1980, Proceedings of the IUE Conference at the Goddard Space Flight Center, May 1980.

Israel, F.P. and Kennicutt, R.C. 1979, preprint

Lambert, D.L. 1978, MNRAS 182,249

Nandy, K., Morgan, D.H., Willis, A.J., Wi1son, R., Gondha1ekar, P.M and Houziaux, L. 1980, Nature 283, 725.

Peimbert, M. and Torres-Peimbert, S., 1977, MNRAS, 179,217

Perinotto, M. and Patriarchi, P., 1980, Ap.J. 235, L13

Perinotto, M, and Patriarchi, P., 1980, Ap.J. in press

Rosa, M. 1980, A & A, in press

Sear1e, L. 1971, Ap.J. 168, 327

Seaton, M.J. 1980, Proceedings of the IAU XX Assembly, Montreal 1979

Shield, G.A. and Sear1e, L. 1978, Ap.J. 222,821

Smith, H.E. 1975, Ap.J. 199,591

Stasinska, G. 1980, A & A 84,320

Torres-Peimbert, S., Peimbert, M. and De1tabuit, E. 1980, Ap.J. 238,133, - 1 11 -

l/2o.0p IfOO.OP If80.0p ,I 60.0,o 1440.00 2O 0 IrOO.op 8O 0 1760·00 1?40.0p Ir20 .0,o ! r II Ir . ,o Ir . ,o II ..N 0

~ 0

Cl> ~ 0 C

0 ~

0 .. = (J =Gl 0 cii I :E: = ?!; I ~ cii =(J iii "Cl)= I I 0 0 I I I 0'"

0 N 0

a: 0 '- ~ u w UJ '-0 N~ ~o U '- Cl o::::il w·0

X :::::>0 u...o-l ...

0 N 0

0 0 0 1120.00 1200.00 1280.00 1360.00 1440.00 1520.00 1600,00 1680.00 1760.00 1840.00 1920.00 WAVELENGTH

Fig. 1 a HII region N63 in the LMC extracted from IUE SWP 3490. (raws 24-27 of the line by line spectrum). Exposur~ time 23700 seconds. The spectrum was obta1ned by P. Benvenuti, S. D'Odorico and M.A. Dopita.

b Supernova remnant N63A in the LMC extracted from IUE SWP 3490 (raws 30-33 of the line by line spectrum) 1 c star 0 Ori C, spectral type 06p. -13 The multiplying factor of the intensity scales is 10 for a and b, 10-8 for c. - 112

1760·00 II

... c.l o

o c.l o

~ ci = ~ i = ~ In in ci =(J ~ I I in ~ ~ I (J ==CII '"I , If :z: '5 0", - , I ...... 0 4( (J III ci , I

0:. ,,,? wo b w ,(J) N 2:: ,W e;, o:::~ W· 0 x =>", u..o-l0

.0 p

0 0 0 1120.00 1200.00 1260.00 1360.00 1440.00 1520.00 1600.00 1680.00 1760.00 11840.00 1920.00 WAVELENGTH

Fig. 2. The full slit spectrum, SWP 7349 of NGe 604 in M33 (c) and two sections sampled perpendicular to the dispersion: a (raws 26-28), b (raws 30-32). The spectrum was originally taken by M. Rosa. Exposure time 6360 seconds. - 113 -

IUE Observations of blue compact galaxies

C. Barbieri 1 and D. Kunth 2

We would like to show preliminary results obtained 2 months ago with the IUE on three blue compact galaxies. The three galaxies were selected upon criteria based on their absolute luminosity (Mp: -16, -17) and their optical appearance on photographic plates taken at Asiago Obser vatory (Barbieri et al, 1979). The galaxies chosen were Mkn 170, Mkn 907 and 11 zw 70 one shift or less was given per object with the low resolution primary short wavelength camera. Figure 1 shows the spectra obtained. The best signal was achieved on 11 zw 70, the slit was centered on the bright nucleus of this galaxy and the exposure lasted seven hours. The processed spectra are displayed in Figure 2. No cons picuous emission lin~are noticeable, the spikes are due to reseau marks or bad spots unremoved. Interstellar lines are certainly present in 11 zw 70. Interstellar lines such as SII 1257, Sill 1265, 1304, CII 1335, and Al III 1857 are clearly present in IIZw70 and possibly the SII and Sill absorptions in the two remaining objects. Redshift absorption lines of stellar origin or due to a hot coronal gas are observed in IIZw70 such as SiIV 1394, 1403, and CIV 1548, 1551.

To get an idea of the average temperature of the ionizing stars in those objects, we attempted to fit the continuum with a black body curve. The curves are plotted in figure 2 and the correspondin~ black body temperature

1 Istituto di Astronomia, Padova, Italy 2 European Southern Observatory, Geneva, Switzerland - 114 - indicated. We believe that the fit is only meaningful for 11 Zw 70 which shows the well exposed continuum and that leads + to a rather high temperature of the order of 40000- 4000. From the total flux one estimates a number of 0 stars 4 of the order of 10 in the emitting region. One may also attempt to correct the observed spectrum allowing for reddening. If we take the reddening correc tion factor from the one obtained by Lequeux et al (1979) using the Balmer decrement, it introduces a significant steepening of the continuum corresponding to a black body of more than 80 000 OK. This result shows either that - the continuum is not reddened by the same amount as the emission lines. This can be achieved if the filling factor inside the Galaxy is very small and the continuum escapes practically un reddened - or that massive stars contribute significantly to the continuum in the U~ which in turn would then very much decrease the number of 0 stars derived above. Another shift is planned on the long wavelength camera and may help to clarify this point by estimating the reddening from the bump at 2200 R.

References Barbieri, C., Bonoli, C. and Rafanelli, P. 1979, Astron. Astrophys. Supple 1I, 541 Lequeux, J., Peimbert, M., Rayo, J.F., Serrano, A. and

Torres-Peimbert, S. 1979, Astron. Astrophys. ~, 155. - 115 -

1 ...3 IIZ.,78

• • •• ~ N • ...._ ...... __------...,;

8957.488 1488.288 2821.654 I 8815 "KN178

'...... ~ ',' .~ . • •• • '. 'f tI· .: • ..' ~ ••:, _1"·~"'~· ,~- ...... '" ·..,. . . . ;.. . . . • , . . -". .'. .. . e. I• ". . . J 8954.288 1485.411 2821.912 I 8828 "KN987

•• 4t:. '. '. , . , ~ -,;• .;.'<' •• 0'. '''~~'' '-, ' . ~z ...... ------/' " . '. . ~ . ..

8957.481 1488.211 2821.654

INTERGALATIC HII REGIONS IUE SPECTRA

Fig. 1 - 116 -

flZ'l/18 ILE Sf'E[TRUt I ei BB 4000Q·K ~ z « m Vi. &=I ~ Kt 1Il = ?: N- - E I u ?: ~ u tu I C'o - IS :l0s;2 &i

I ei 11:ill iDl 137 917 1 ~ WAVELENGTH 751 2flI'5l.161

~ 118 ILE Sf'E[TRI.)t

BBJOOOO"K

I ei '-11--:ill"-iB/J:::::------,r------"'---..L..Ll5B4.-.-1-76------JLL..-.,------::2B51l.-:--"--'3S3 WAVELENGTH1

BB 20000"K

~ 751 2flI'5l.161 137 917 WAVELENGTH1 Fig. 2 - 117 -

The helium abundance determination in emission-line dwarf galaxies

G. Statinska 1 and D. Kunth 2

Introductory remarks

The gas rich galaxies with definite metal deficienc~ such as II zw 40 and I zw 18, are ideal for determining the abundance of the primordial helium. Such a study has been made by Lequeux et al (1979) who find a value of He/H = 0.074 ± 0.006. However the quoted uncertainty takes into account only the scatter in the abundanc~found among different objects and not the possible errors in the deter mination of these abundances. Clearly the problem requires further studie~ involving good quality observation and a larger sample of galaxies in order to minimize the influence of random errors in the interpretation of the data which

as su~marized below are not negligible. i Apart from the uncertainty due to the reddening correction, the observational problems include a possible contamination of HeI 5876 by the NaD line at 5890 Rdepending on the redshift, and event~ally the presence of an underlying absorption by hot stars in the Balmer lines of hydrogen and/or in the HeI lines. One possibility to overcome this difficulty is to work at high enough resolution (~3R) and build up a good SiN in the continuum. In the case where the Hell 4686 emission line is observed, the hypothesis of its stellar origin has to be considered (Bergeron, 1977). In this connection, the Space Telescope might be useful to isolate nebulosities free from stellar background in which spectroscopy could be performed.

Observatoire de Paris, Meudon, France 2 European Southern Observatory, Geneva, Switzerland - 118 - Interpretation + + The determination of the He /H ratio shouln be quite accurate if radiative recombination alone were at work to produce the observed HeI lines. Indeed, the emissivities involved are well-known (Brocklehurst 1971, 1972) and the result is little dependent on the electronic temperature. However, Cox and Daltabuit (1971) pointed out a possible contribution due to collisional excitation and self-absorp tion from the metastable level 2 3 S. Brocklehurt (1972) argues that the collisional rates given by Cox und Dalta buit were overestimated by a factor 3 at least. 2 -3 Thus the effect of collisions will be negligible at ne

Self-absorption does not affect too strongly the two most intense linesof HeI, namely 5876 and 4471, but still could contribute by ~lO% at the most (Cox and Daltabuit, 1971). The total He/H abundance is obtained once the amount of neutral helium eventually present in the galaxy has been estimated. Unfortunately, all the "recipes" used so far are in fact illusive and it turns out that the total ++ helium abundance can be reliable only if 0 /0 ~.8 (Statinska, 1980) which fortunately is often the case in metal-poor galaxies. - 119 -

References

1) Barker, T. 1978, Ap.J. 220, 193

2) Bergeron, J. 1976, Ap.J. 210. 287

3) Br 0 C k 1 e h u r s t, l-1. 1 9 7 1, H. N . R . A . S . 1 5 3, 4 7 1

4) Brocklehurst, M. 1972, M.N.R.A.S. 157, 211 5) Cox, D.P. and Daltabuit, E. 1971, Ap.J., 167, 257. 6) Lequeux, J., Peimbert, M., Rayo, J.F., Serrano, A., and

Torres-Peimbert, S. 1979, Astron. Astrophys. ~, 155 - 121 -

HI observations and star formation in the blue compact galaxy I Zw 18

J. Lequeux, F. Viallefond

This is a short summary of a paper to appear in Astronomy and Astrophysics

21-cm line observations of I Zw 18 with the Westerbork Synthesis Radio Telescope have resolved it into a complex substructure. A possible model consists in about 6 hydrogen clouds with hydrogen masses between 3 and 6 30 10 MO (at an assumed distance of 10 Mpc). Acontinuum emission of 2.5 mJy has been detected at the 3a level and presumably corresponds to the thermal radiation of the HII regions seen optically. Some HI compo nents and the HII regions have a particularly high mean gas density. From our 21-cm line observations we obtain a virial total mass for two HI clouds and for the whole galaxy, which in all three cases is signifi cantly larger than the HI mass. The excess mass can be that of molecu- lar hydrogen, or of old stars, or of a combination of both. We suggest that I Zw 18 is a galaxy presently in the process of formation by a merging of primordial clouds which were previously supported against gravitational collapseby quiet formation of stars of intermediate masses. 6 The present burst of star formation is probably not older than 5 10 years. It might be the first in the history of the galaxy. I Zw 18 appears to be the best - and perhaps even the only known - candidate for a "young" galaxy. - 123 -

Isolated Extragalactic HII Regions. "The Triggering Mechanism".

G. Tenorio-Tagle* European Southern Observatory

Abstract. - The collision of intermediate and high velocity clouds with a galactic disk is proposed as the mechanism res ponsible for the formation of isolated extragalactic HII re gions. The observational facts and the properties of the model are stressed here.

I. Introduction

Estimates of the number of (and mass in) stars in extra galactic HII regions have been obtained (Sargent and Searle, 1970) by using Blaauw's (1964) data for 0 associations (M/L ~ 2 5xl0- ). Knowing the luminosity (L) of the objects this has led 6 to masses. M N 10 M0 in the form of young stars. From Zantra's theory and the observed H fluxes one can infer a consistent result that implies some 103 to 105 0 stars born during the last 7 10 yr. All these stars are concentrated in a small volume of" 2 typical size D (D~ 10 pc) where the gas density is high (n -~ 10 2- 10-3 , O'Connel et al., 1978). The core of the HII region is surrounded by an extended halo of low density ionized gas (typical size R~ 300 pc), and the whole nebula is imbedded 8 in a massive (M ~ 10 Mo) HI cloud. From the colors of the HI objects Sargent and Searle (1970) concluded that either star formation just started or that the galaxies have only created o stars during the past 1010 yr. In a further investigation, (Searle and Sargent, 1972) the chemical composition of a few objects was obtained. Oxygen and Neon were found to be under abundant with respect to cosmic values while the He abundance was found to be normal. From this information and the color~

*Present address: Max-Planck-Institut fUr Astrophysik, Karl-Schwarzschild-StraBe 1, D-8046 Garching. either the galaxies are young or the star formation in them occurs in strong bursts which are separated by long quiet periods. 8earle et al. (1973) obtained an estimate of the strength (8) of the bursts (10 ~ 8 ~ 20, where 8 = R /R , and R is the 8 10 8 star formation rate averaged over 108yr and R is the rate 10 10 averaged over 10 yr) and the number of bursts (5 to 10) required to explain the isolated extragalactic HII regions. Also through statistical arguments the idea that the objects are young was dropped. However, no attempt was made to explain the origin of the bursts.

11. The Trigger

The problem is to find a mechanism able to compress a large amount of material and to trigger a gravitational instability that could lead to large numbers of stars concentrated in a small volume. Further, the process should repeat itself after 8 a long (>10 yr) quiet phase, as many times (during 1010yr) as necessary to produce the observed metal abundances. Therefore, from the mechanisms already available in the literature, one can immediately exclude processes such as fragmentation of the protogalactic cloud and the Parker instability (Woodward, 1978). In a similar manne~ other mechanisms able to initiate the col lapse of a cloud by a sudden increase in the external pressure can be excluded. This is either due to lack of observational evidence (such as the implosion of gas clouds by shocks pro duced by spiral arms), or because the mechanisms operate for a short time but perhaps in a continuous manner (as the sequent ial formation of stars due to the expansion of supernova rem nants, of HII regions, etc.) which disagrees with the time E~ paration between bursts. Other major events such as the col lision of galaxies will fail to explain the compactness of the isolated giant HII regions. In this way one is only left with effects such as thermal instabilities, the agglomeration of stable clouds (due to cloud-cloud collisions) and the collisions of intermediate and high velocity clouds (I-HVCs) with a galac tic disk, to compose a scenario which may account for the va rious observational facts. 125 -

Let us assume that I and HVCs exist in blue dwarf (and in all) galaxies. Suppose that they were formed either during the collapse and fragmentation of a rotating protogalactic cloud and were trapped as satellites of a newly formed galactic disk (composed of many more fragments), and/or that a galaxy can continuously re-generate them as it moves through an inter galactic medium (Oort, 1970). The collisions of clouds (or fragments) in the disk will eventually lead to agglomeration (i.e. larger stable clouds) and to a diffus~more extended me dium (see Stone, 1970; Tenorio-Tagle 1980). These processes will be favoured by thermal instabilities in the gas and by collisions (and consequently losses) of I-HVC with the galactic 2 3 disk of densities (n ) in the range 1 ng~ 10 cm- (see g < Tenorio-Tagle 1980). other satellites colliding with the lower density medium (say more than 90 % of the disk cross-section) will traverse the disk and continue almost in a similar orbit. Only a few of them will eventually collide with a dense conden sation (generated in the disk). This latter kind of collisions will lead to the formation and gravitational instability of a 6 massive (M > 10 M)o layer of shocked gas,out of which several hundreds of early type stars may result, even if star formation is a very inefficient process (say 1%). The stars will be con centrated in a small volume of radius determined by the size 2 of the impinging cloud (say, 10 pc) and the ionization of the gas left over from the collapse will generate a giant HII re gion. During the following 107 yr the disruption of the dense condensation will take place (through super novae explosions 8 and the expansion of HII regions etc.). More than 10 yr-will be necessary to build a new target (a dense cloud) in the same area (Woodward 1978). During this time a blue dwarf galaxy will return to its equilibrium color (B-V~0.36, Searle et al. 1973), 8 and only after several times 10 yr the probability of an I-HVC colliding with it and generating a new giant HII region will approach 1.

Ill. Conclusions

Our model requires that a small fraction of the total HI mass of blue dwarf galaxies exists in the form of I and HVCs. - 126 -

The probability of a successful event (i.e. the generation of a giant HII region) depends on the cross-section of the target and the number of satellite clouds (which may be depleted with time). Several density condensations may form in the disk over 8 long periods of time (?10 yr), but not all of them will ne cessarily become a successful target and turn into stars at the same time. In this way one can explain the occasional dis placement between the optical and 21 cm emission centers (as in 11 Zw 40, Gottesman and Weliachew, 1972).

The present, necessarily speculative, model seems to ex plain the large number of stars, the compactness of the objects, and offers a logical explanation for the time separation bet ween bursts of star formation. Other processes such as tidal interactions between galaxies may be a way of re-generating I and HVCs in a galaxy and/or may be responsible for orbital de flections leading to a larger number of successful collisions. However, a tidal process is certainly not necessary and/or sufficient to produce an isolated extragalactic HII region.

Our ideas are also consistent and complementary to the conclusions of Searle et al. (1973) based on the color obser vations of late-type galaxies.

Detailed calculations of the model described her&are now underway.

I'm grateful to Drs. D. Kunth, D. Alloin and P. Bedijn for many fruitful discussions on the subject.

References Blaauw, A.: 1964, Ann. Rev. Astr. and Ap. ~, 213. Gottesman, S.T. and Weliachew, L.: 1972, Astrophys. Lett. ~, 63. O'Connell, R.W., Thuan, T.X. and Goldstein, S.J.: 1978, Ap. J. Lett. 226, L11. Oort, J.H.: 1970, Astron. and Astrophys. I, 381. Sargent, W.L. and Searle, L.: 1970, Ap. J. Lett. 162, L155. Searle, L. and Sargent, W.L.: 1972, Ap. J. 173, 25. Searle, L., Sargent, W.L. and bagnuolo, W.G.: 1972, Ap. J. 179, 427. - 127 -

Stone, M.E.: 1970, Ap. J. ~, 227. Tenorio-Tagle, G.: 1980, E.S.O. reprint No. 102, Astron. and Astrophys. (in press) . Woodward, P.R.: 1978, Ann. Rev. Astr. and Ap. ~, 555. - 129 -

THEORETICAL IMPLICATIONS OF A LOW HELIUM ABUNDANCE IN DWARF GALAXIES

Jean Audouze, Institut d'Astrophysique du CNRS, Paris

Dwarf galaxies are known to be objects where the metallicity Z and also the Helium abundance are significantly lower than in our own Galaxy (see e.g. Peimbert, 1975, Lequeux et al., 1979): Z can be lower there factors up to 10 and Y might be as low as 0.22 with still large uncertainties on this last value.

Today, it is generally accepted that 4He (such as D,3He and also 7Li ) is sYnthesized during the Big Bang. While the D and 7Li abundances depend strongly on the present density of the Universe (and are good indicators of an open dontinuous ly expanding Universe), the 4He abundance depends on (i) the neutron lifetime which is not yet precisely determined by the various experiments, (ii) the rate of the Universe expan sion. If the expansion is rapid, the resulting He abundance is high while if it is slow, the primordial He abundance should be low. The characteristic time scale of expansion depends on many factors:

(a) The homogeneity and the isotropy of the Universe.

(b) The choice of the time space metric. The "canoni cal" Big Bang models use the metrics derived from the General relativity.

(c) 'The total number of neutrino species. The primor dial He abundance increases with the number of different neutrino species (see fig. 1 coming from Yang et al., 1979): there are at least three dif ferent types of neutrinos: the electronic neutrinos, the muonic neutrinos and the tau-neutrinos. - 130 -

Other types of neutrinos are searched and may be discovered in the future.

Therefore the search of He in dwarf galaxies is very impor tant for two different reasons: (i) It has strong implications on the elementary particle physics, especially the physics of the leptons. (ii) If the He abundance is found to be very low, Y ~ 0.20, the "canonical" models of the Big Ban~ such as those developed by Wagone~ would run into very severe difficulties.

(from Yang et al.)

Cl> U c: o :2 "0 c: ~_--:-__-'o- :::> .0 ~ -8 Cl> 10 -32 31 30 ..:I: 10 10 10- PHI T:27°1( )(g/cm3 )

FIG. I.-Primordial abundance of 'He for !:l.NL = 0 I 2 FIG. 2.-Primordial abundance of 2H for !:l.N = 0, 1,2,3, 3,4, and S. ',, 4, aDd S. L

REFERENCES

Lequeux, J., Peimbert, M., Rayo, J. F., Serrano, A., Torres Peimbert, S., 1979, Astron. Astrophys., 80, 155. Peimbert, M., 1975, Ann. Rev. Astron. Astrophys., 12, 113. Wagoner, R. V., 1973, Ap.J., 179, 343. Yang, J., Schramm, D. N., Steigman, G., Rood, R. T., 1979. Ap. J ., ~, 697. - 1)1 -

SESSION V

(Chairman: L. Woltjer)

ELLIPTICAL DWARF GALAXIES

(a) Chemical Abundances, Main Sequence Observations - 133 -

Abundance Problems in Dwarf Spheroidals and Dwarf Ellipticals of the Local Group.

J. Danziger European Southern Observatory clo CERN - CH-12l1 Geneva 23

1 - Introduction A comparison of the properties of dwarf spheroidals and dwarf ellipticals in the Local Group shows that, in respect of absolute luminosity and mass, the dwarf spheroidals define a continuous sequence between dwarf ellipticals and globular clusters. None of the dwarf spheroidals however have the central surface brightness (or semi-stellar nucleus) observ~in members of both other groups. The fragility of dwarf spheroidals and their radii, interpreted to be tidally limited by our Galaxy, mean that even at their perigalactic distances they have not been close to the galactic center. On the other hand with M32 we have one example of a dwarf elliptical which has plausibly been tidally stripped by M3l to the extent of diminishing its luminosity by 2-3 magnitudes (Faber 1973).

2 - Information on Abundances of the Elements There are, various pieces of information which provide limited insight into abundances of heavy elements in these systems. - 134 -

(a) Morphology of HR diagrams. In the dwarf spheroida1s the ear1ie~ work produced HR diagrams which, although different from one another in detail, showed similarities with those of galactic globular clusters. Recent photometry by Demers, Kunke1 and Hardy (1979) revealed a broad giant branch among the field stars of Fornax. This was interpreted as evidence for a significant range in metal abundances, as had been done earlier for the globular cluster W Cen.

(b) Globular Clusters. Spectroscopy (van den Bergh 1969) and intermediate band photometry (Danziger 1973) of the integrated light of the globular clusters in Fornax showed a range of meta11icity, which covered the range from the median to the extreme metal poverty observed in galactic globular clusters.

(c) Planetary Nebula. The planetary nebula in Fornax (Danziger et al 1978) has heavy element abundances only a factor 2-3 less than solar values. This result applies not only to Nand 0 but also to elements heavier than carbon, nitrogen and oxygen,these latter being the only ones likely to be produced in solar mass stars. The planetary nebula in NGC 185 observed by Ford (1978) has element abundances similar to those of galactic halo planetary nebulae.

(d) Individual Stars (i) Individual stars observed in Scuiptor (Norris and Bessell 1978) and Draco (Zinn 1978) show a range in element abundances in both. Neither of these objects have broad giant branches. - 135 -

(ii) Recently Aaronson and Mould (1980) have observed carbon stars in Fornax , and in deriving M = -5.6 for them,conclude that they cannot bol 9 have ages much greater than 2 x 10 years. From the standpoint of the evolution of the galaxies concerned,t?ese various observations have been interpreted individually in different ways. They are: 1) Primordial scatter in element abundances within one galaxy, 2) Heavy elements produced 1n one generation of evolving stars, 3) Several generations of star formation giving rise to a range in metal abundances as a result of multiple processing of material.

Although the third interpretation might seem empiricallY to be the most plausible scenario, because it might account for most of the observations, it should be remembered that Fornax (as well as the other dwarf spheroidals) has a very low escape velocity (10 km/sec). Models of galaxies 8 9 starting life with masses of 10 - 10 Ma and producing several generations of stars before the gaseous partly processed material is swept out by stellar or supernova driven winds remain to be calculated in detail. Ford (1978) and his colleagues have argued on the basis of detection of planetary nebulae,that even the more tightly bound dwarf ellipticals in the Local Group have lost a sizeable fraction of the gas expelled from stars as planetary nebulae. - 136 -

3 - Future Observations with the Space Telescope and from the Ground. (1) Spectroscopic observations of planetary nebulae in dwarf spheroida1s and dwarf e11iptica1s, paying attention to both the CNO group and heavy elements, would help our understanding. The Space Telescope would be particularly useful in observing in the DV 10ns of He, C, N, 0, Ne, Ar, Mg and possibly Si and S. It may even be necessary for optical observations of the faint p1anetaries in the dwarf e11iptica1s.

(2) Main sequence photometry (narrow or intermediate band) or low resolution spectroscopy of the same stars should be feasible with the ST for the dwarf spheroida1s. These results provide not only information on ages but also on metal 1icity and particularly the possible spread in this parameter. (3) Call absorption line photometry of RRLyr stars in the globular cluster W Cen has shown a range in the meta11icity. This technique might be used in the dwarf spheroida1s from the ground and in the dwarf e11iptica1s (if RRLyr stars are found) from the ST. (4) Intermediate band photometry on giant stars in dwarf e11iptica1s is a logical extension of the technique used in nearer systems. Whether demonstrating that a range of meta11icity exists or not in these systems will help to elucidate the problems we are confronted with no~ is not clear. - 137 -

(5) It might be ~nteresting to observe spectrosco pica11y the dwarf e11iptica1s in the Virgo cluster to ascertain what fraction, if an~ have had an evolutionary history similar to that proposed for M32 proposed by Faber (1973), and mentioned above.

References

Aaronson, M. and Mould, J. 1980, preprint Danziger. I.J. 1973, Ap.J. 181, 641. Danziger, I.J., Dopita, M.A. Hawarden, T.G., Webster, B.L. 1978, Ap. J. 220, 458 Demers, S., Kunke1, W.E., Hardy, E. 1979, Ap.J. 232, 84 Faber, S.M. 1973, Ap.J. 179, 423 Ford, H. 1978 IAU Symposium ~' p. 19 (Reide1 Dordrecht Holland) Norris, J., Besse11, M.S. 1978, Ap.J. (letter) 235,L49 van den Bergh, S. 1969, Ap.J. Supp1. ~' 145 Zinn, R. 1978, Ap.J. 225, 790 - 139 -

(b) Globular Clusters in Dwarr Galaxies - 141 -

GLOBULAR CLUSTERS IN DWARF GALAXIES

Vittorio Castellani Istituto di Astrofisica Spaziale,Frascati

Globular Clusters are highlY relevant test objects for the understanding of the nuclear evolution of matter, since they date from the very early phases of the universe and -more in general- they can supply information on the relation between age and metal enrichment in cosmic matter. As an example, one finds that GC's are really picturing the very initial phases of evolution in our Galaxy. As a matter of fact, when looking at the density of "metal poor" galactic Globular Clusters projected on the XZ plane (that is: on the plane passing through the sun and the galactic center and per pendicular to the galactic disk) we can get evidence of a spheroidal distribution, as disclosed in fig.1, which is very likely connected to a spheroidal phase of the protogalactic cloud. This is not the case for "metal rich" Globular Clusters, their distribution having been originated in a more advanced stage during the galactic collapse, shoving how GC's ca"o be used as probes for the study of the protogalactic matter and early galactic evolution. This is only an example of how GC's can provide us the "1ocal" relation between time and chemical enrichment. Hope fUlly, a sufficiently large amount of such relations will tell us the history of matter in our universe. - 142 -

Sp.T: FO -. F7

o 5 10 kpc

Sp.T. )F7

Fig.1: The projected density distribution on XZ plane of "metal poor" galactic globular clusters, whose spectral types range from FO to F7 (upper part) compared with the same dio tribution but for "metal rich" clusters with spectral types later than F7. Isodense lines have been obtained by applying Walsh-Hadamard transforms to the clustE:~r coordinates (Caste2 lani and Melchiorri 1980) - 143 -

How far from our Galaxy can we use such a probe? If the brightness distribution of the galactic Globular Clusters is fitted with a gaussian profile, an extrapolation to M 31 shows , (Battistini et al. 1980) that "4G'"" diameters are expected between 1".4 and 24". Assuming the Virgo Cluster at 24 Mpc

(i.e. Ho= 50) it turns out that the Space Telescope will be able to study the structure of the largest GC's at this dis tance. The role of GC's in Dwarf Galaxies is obviously dependent on the nature of the galaxy we are dealing with. As far as the problem of chemical evolution is concerned, we are generally interested in DG's and, in turn, in GC's in DG's, since they have evolved in isolation from the principal chemical enrich ment occurred in large galaxies. We are particularly interes ted in the lower limits of this chemical enrichment, because the helium content will give us strong constraints on the big bang structure, and the metal content will tell us something more about the still mysterious mechanism which spreaded eve rywhere (as fur as we know at present) an abundance by mass of heavy elements of the order of Z ~:.\, 10-4,i.e. an amount of "metals" much larger than expected from the nucleosynthesis occurred in the big-bang. So we are particularly interested in Dwarf Spheroidal Ga laxies (Sculptor type) as they are true Population 11 objects, which overlap GC's in luminosities, colors, masses (see table 1), without gas or dust, and which cannot survive close tidal ~counters with large galaxies. - 144 -

) R(kpc) M(M 0 D(kpc) w Cen 0.09 1.2+6 8.

Ursa Minor 1.2 1• +5 70 Draco 0.5 1.2+5 80 Sculptor 1.2 3. +6 80 Carina 170 Fornax 3.1 2. +7 190 Leo I 0.9 4. +6 220 Leo 11 0.6 1 • +6 220

NGC 147 2.2 6. +7 700

Table 1: Radii,masses and distances from the galactic center fbr DSG's and for GC \}J Cen and DG NGC 147

=====-======

DGs: Fornax N= 6 (B-V)= 0.51,0.61,0.65,0.66,0.74,? NGC 205 8 0.51,0.55,0.57,0.58,0.64, 0.72,0.76,? NGC 185 6 0.69,0.71,0.79,0.94: ,?? NGC 147 4 0.64,0.80,?? M 33 5: 0.74,0.77,0.80,1.16,1.26 (reddened GC's?)

GGC's: M92 (B-V) 0=0.61 Fe/H -2.12 M3 0.68 -1.57 NGC 6723 0.72 -0.85 47 Tuc 0.85 -0.44

Table 2: The number N of Globular clusters d~overd in DG's membering the local Group and the integrated (B-V) colors of these objects, when available. Integrated colors and metalli cities of typical galactic Globular Clusters are reported at the bottom of the table for the sake of comparison.

======- 145 -

What about the observational evidence in the local Group? Till now we have no direct information on ages, as we cannot acheive turn-off luminosities out of our Galaxy. We only ob serve a spread in integrated colors among GC's membering the various DG's (table 2), among the DSG's themselves, and among the Globular Clusters membering the DSG in Fornax. This spread in colors has been generally interpreted in terms of a spread in original metallicity. Were it true, com parisons with galactic GC's would indicate a spread of this parameter among the Fornax clusters by a factor of 10. It turns out that these objects may not be direct probes of the early universe, but rather of some other equally interesting pheno mena, responsible of the local enrichment in the various clu sters. What the ST will enable us to know is sketched in fig.2: no doubts we will be able to derive a more clear picture of matter evolution and a powerful test for the mechanism of me tal production. We know that something interesting will appear, because we see that something odd happens in DSG's, though we have only minor suggestions on the origin of such oddities. As a matter of fact, "peculiar cepheids" in DSG's are often quoted: I would like to stress that, like cepheids, also RR Lyrae va riables show clear peculiarities, as shown in table 3 where properties of DSG-RR Lyrae are reported in comparison with the behaviour of galactic pulsators. When remembering that also HR diagram locations show other peculiarities, one can only conclude that DSG's are completely peculiar objects, and we have to wait for ST to look more in datail in this problem. - 146 -

-2 • • • 0 • •• 0 0 • +M31 • • CIS... • • u 2 • -Q) • a. • en • +Fornax 4 •• • +M31 • ...en 6 • :::I • 0 0 • u • +Fornax 8 • •• B-V -.4 .4 1.2 1.6

Flg.2- What we wll1 be able to observe In GC's wlth ST at the dlstance of Fornax or M 31. Photometric and spectroscopic u1 (6/'..//"=10-1.) limits mve been assumed at 2g"'and 25 ,respecti vely. - 147 -

min P ab

Draco 0.617 0.537

Ursa IJIinor 0.636 0.490

Sculptor 0.566 0.479

Leo 11 0.593 0.487

Galaxy:

Oosterhoff I 0.551 0.454

Oosterhoff 11 0.637 0.543

Table 3:The mean period of Bayley's ab-type pulsators in Dwarf Spheroidal Galaxies membering the local Group and the minimum observed period for this kind of pulsators. The same quantities are reported at the bottom of the table for the. two typical group:in which galactic globular clusters can be subdivided. Note that, in particular, one expects Pa~in to be directly connected with evolutionary parameters like age and original chemical composition.

======.======:-0·:'::= - 148 -

What about ST job? I would like to give you just my per sonal list of priority,i.e. i) HR diagrams for GC's in Fornax ii) pulsational properties of RR Lyrae ( if any) in the above quoted clusters, iii)turn-off luminosities for DSG stars. After this first step, DG's and Globular Clusters in DG's within and outside the local group have to be searched to ex tend this analysis in a very general frame, in which ttre pro blem of the possible spread of evolutionary parameters within each cluster has also to be discussed. Along this way, no doubt we will face new evidence and new problems which in turn will require further and new investigations.

As a conclusion, I just want to stress that this proble matic still needs quite a lot of job from ground based tele scopes: we are still lacking information about turn-off lu minosities for more than 90% of the galactic globular clusters, which means that we are still waiting for the key-data to be inserted in each theory of galactic forma~ion.

REFERENCES

Battistini,P.,Bonoli,F.,Braccesi,A.,Fusi Pecci,F.,Malagnini, M.L.,Marano,B.:1980,Astron.Astrophys.Suppl. (submitted to) Castellani,V.,Melchiorri,M.:1980, in preparation Discussion

M.-H. Ulrich: 11 some of the stars in the halo of our Galaxy were -captured from dwarf spheroidals, why do we not find RR Lyrae in the halo with periods similar to the periods of the RR Lyrae stars in dwarf spheroidals? v. Castellani: It is not easy to decide whether or not the halo (field) RR Lyrae have characteristics connected with spheroidal galaxies. As a matter of fact, the proper ties we were talking about are essentially "collective" properties, as the "mean period" or the "minimum period" of' a sample. So, for a single RH Lyrae nothing can be said as - in general the period is compatible with both galactic or extragalactic origin. As far as collective properties of the halo are concerned, we know that the halo has to be at least "contaminated" by stars which have escaped from the clusters, so that collective properties are very easily "deformed". This is particularly true of the "minimum.period", whereas the mean period may pe moved only if the number of galactic and extragalactic pulsators are comparable. And we have no indication of the relative amount of both kinds of pulsators. - 151 -

ARE SOME DWARF GALAXIES CROUCHING GIANTS? M. J. Disney Department of Applied Mathematics and Astronomy, University College, Cardiff, Wales, U.K.

1. INTRODUCTION Two good motives for interest in dwarfs are: (a) It could be that dwarfs contribute significantly to the mass density in the field, in groups, in clusters and in the Universe as a whole. (b) It may be that the galaxy formation is a comparatively short lived process. If that is so, and if formation continues around us today, then to catch one of them lat it' observers will need to study the most numerous galaxies about, which happen to be dwarfs. I would here like to discuss proposition (a), and to point out that even if, on the face of the evidence, the proposition looks to be unsound, the evidence itself may be seriously biased by observational selection. Here I speak of selection in the sense that the observed 1uminosities, and hence the masses of galaxies, may depend rather crucially on their surface brightnesses relative to the background sky.

2. THE APPARENT MASS SPECTRUM OF GALAXIES Numerous authors have discussed the luminosity function of galaxies, and by imputation their mass spectrum. For instance there are excellent reviews by Ho1mberg (1975) and by Peeb1es (1971), and a particularly good recent discussion of the mass to light ratios by Faber and Ga11agher (1980). To summarize: (a) Spirals and SO's can be separated from the rest. Their number density as a function of absolute magnitude would appear to be roughly Gaussian with a mean Ho1mberg magnitude of -17·7 and a dispersion of ± 1.2 magnitudes. Why there should be no dwarf spirals we do not know but at least we can ignore them here. (b) For the remaining e11iptica1s and irregulars, Zwicky (1957) appears to have been right in that the numbers increase exponentia11y - 152 - to fainter and' fainter absolute magnitudes. Assuming there to be.no systematic relation between m/l ratio and absolute luminosity the observed luminosity function can be converted to N(m).dm, the number density of galaxies per unit volume by mass. All but the very brightest galaxies would appear to fit the relation:

N(m) = K m-a.dm with a fixed (1 ) and hence Emo' the total mass of all galaxies in the range mo to 10 mo will be given by: tm ~ m 2-a (2) o 0 from which it is clear that dwarf galaxies in toto will contribute significant mass only if a ~ 2.

let Elo denote the relative number of galaxies of intrinsic luminosity La to 10 Lo in an apparent-magnitude-limited sample, e.g. a catalogue, then from (1): EL ~ m 5/2-a (3) o 0 from which we deduce that most of the galaxies in the catalogue would be dwarfs if a > 2·5. Since this is patently not the observed case we know at once that a < 2·5. Holmberg (1969) studied 300 galaxies in a number of local groups. The survey, which goes down to MH = -10, but is only complete down to -13·5, yields a = 1'5. Abe11 (1975) studied cluster galaxies. His sample is larger but because of the much greater distances to clusters it goes down only to M= -16, but yields for the fainter galaxies essentially the same value for a as Holmberg and Zwicky found. This result allows us to draw up some crude but useful aide memoires. If we assume: (a) most non-spiral galaxies exist in the absolute magnitude range Mv = -10 to -20; (b) dwarfs and giants are separated at Mv = -15; (c) the exponent a = 1·5 prevails throughout then: (a) Dwarfs are ten times more numerous in space than giants. (b) Giants are on average 100 times more luminous (and hence more massive?) than dwarfs. - 153 -

(c) In any apparent magnitude limited catalogue there will be 100 times as many giants as dwarfs. (d) The dwarf population contributes only 10% of the total ,light (and mass?)

T~ese statements are not greatly affected by the inclusion of spirals.

3. SELECTION EFFECTS Although 2(d) would seem to rule dwarfs out as effective mass contributors to the Universe, one or two uneasinesses remain in my mind which I hope are not simply a sign of personal ignorance. Since I have always understood that the distinction between a workshop like this, and a colloquium, is that one can make a fool of oneself in a workshop, I will proceed to air those uncertainties. The first uneasiness, arid it is a slighter one, concerns 2(c) above. With a = 3/2 in magnitude limited catalogues, dwarfs will apparently be swamped by giants, so that it will not be easy to come to a true estimate of their relative frequency without fairly extensive redshift, or other distance-indicating information. Even if a = 2, when the dwarfs would be significant mass contributors, they would be outnumbered ten to one in the catalogues. Do we have all the redshift information we need, particularly for the lower surface brightness objects (see below)? I do not know. A more substantial worry concerns the quoted luminosities of dwarfs. We know how difficult it is to agree upon the total luminosity of even such a spectacular and well-observed giant as M87. How much more should we worry then over galaxies which, on the face value of the observations, are on average a hundred times less luminous than the giants. This magnitude worry is compounded by surface brightness effects. de Vaucouleurs (1959) has shown that the surface brightness profiles of most galaxies can be characterised by two parameters which in operational terms might be the central surface brightness cr(o) and the

scale-length r 2I • If the independently distributed physical_ parameters which describe galaxies are mass m and mean density p then on - 154 - dimensional grounds we expect: 1/3 _-1/3 r~ '" m p (4) 1/3 _-2/3 a(o) '" m p (5) which means that on average dwar~wi11 have surface brightnesses and radii each down by a factor of 5 on giants. And accurate values of both a(o) and r 1 are needed before the total 1uminosities LT can be inferred from LT '" ~(o) r~2. While this may be obvious in principle, I sometimes wonder if the delusions of isophotal magnitudes are not sometimes understated. To illustrate thi~ Figure 1 shows shorter and longer, but otherwise identical, exposures of the pair of giant el1ipticals NGC 3309 and NGC 3311. Which twin wears the highest total luminosity? The extra observational difficulties confronting the observer of dwarfs means that quoted dwarf magnitudes are more likely than giants to be isophotal rather than total, and simply because of relation (5) above, for dwarfs the last measurable isophote will contain a lower, possibly a much lower, proportion of all the light. And then again, how many dwarf galaxies are masquerading as stars? A wide angle space camera might be able to tell us. My most substantial worry, however, concerns selection effects imposed upon us by the surface brightness of our own sky. Not only do we reside inside a giant (and hence, according to (5), one of relatively high surface brightness) galaxy, but we suffer from substantial zodiacal and atmospheric contributions as well. We are like prisoners in a lighted cell trying to discern our whereabouts by peering out through a small casement into the darkness outside. We can see the street lamps easily enough, and the lighted windows, but can we see, or correctly infer, the houses and the trees? It should be obvious that objects perceived as extended must lie on a broad band in the log apparent luminosity/log apparent diameter plane, and indeed they do, as Arp (1965) long ago pointed out. What is not so widely recognized is that most well-observed galaxies for some reason lie in a highly restricted strip of,that allowed band. Freeman (1970) showed that the discs of 40 well-observed spirals and irregulars 155

• • • - 156 - have central surface brightnesses B(o) defined in the de Vaucou1eurs sense, of < B$(o) > = 21·65 ± 0·3 magnitudes per square arc second, while Fish1s (1964) observations of 30 e11iptica1s yield (Disney, 1976) < BE(o) > = 14·80 ± 0·9. The difference < BS(o) > - < BE(o) > = 6·85 is intriguing, but is mostly an artefact of the mathematical way in which the two types of de Vaucou1eurs' profiles are defined. What is far more surprising is the very narrow range of surface-brightnesses observed in each case. They are far narrower than observational constraints would allow. For instance the archetypal compact I Zw I is still discernable as a galaxy with B(o) close to 8 magnitudes per square arc second. On the low surface brightness end, I am not so sure where the observational limit lies, because it will depend upon the emulsions, the focal ratios, and the plate-scale in use, but it is said (Heidmann, et. al., 1971) that nebulosity as faint as 24 B magnitudes per square arc second is apparent even on the Pa10mar prints, while resolved Scu1ptor/Fornax dwarfs have B(O)I S as low as 25 magnitudes. I would ~ubmit that the narrow observed surface brightness rang~s are remarkable, and call out for an explanation. I know of no sound physical explanation for them, though Fish (1964) made a very ingenious, if not convincing, try for the case of e11ipticals. On the other hand, if we concede that the effects of observational selection are at work in a particularly insidious and dramatic way, then the Fish-Freeman results are straightforwardly explained (Disney, 1976) by what I call the liceberg effect1 which I will summarize only briefly here. Granted only that galaxy profiles are broadly described by the de Vaucouleurs laws then, for any given true total luminosity, it is possible to calculate the apparent or isophotal radii of galaxies on the sky as a function of their morphological types and central surface brightnesses' cr(o) (c.f. Disney (1976) for details). We assume all galaxies to be face on for simplicity. For galaxies of low cr(o) the outer isophotes are quickly lost against the sky and the outermost noticeable isophote will be of small apparent angular diameter. Conversely, for a given LT once again, very high surface brightness galaxies must necessarily be relatively compact with lower apparent - 157 -

angular diameters. It follows that for a galaxy of any morphological type, and any true total luminosity LT' a central surface brightness can be chosen which lends to that galaxy the largest apparent area on a photographic plate of given depth, that is to say the largest area within, an isophote where the nebulosity is apparent to the eye without sophisticated reduction techniques.

I will not belabour the point. Interested astronomers c~n best convince themselves by carrying out the trivial computation involved. But the result is rather startling. All the Fish and Freeman galaxies have surface brightnesses calculated to lend them the most spectacular appearances on contemporary photographic plates. And a glance at the Fish and Freeman samples shows them to contain most of the best photometered galaxies available at that date. Indeed Freeman's sample is largely a compendium of the more accurate published results. All this represents a series of numerical coincidences we simply have to face. Either some interesting physics is at work or we are being seriously misled by observational selection. If the latter is true then our true knowledge of the galaxy luminosity functlon is fairly primitive, and we are being misled always in the same sense, i.e. to underestimate the total luminosity and mass of the galaxY population. If, as one suspects, dwarfs are systematically lower in surface brightness, then we shall systematically underestimate the dwarf contributions to the cosmic density. In particular I suspect that some of the dwarfs we shall be discussing here are not really dwarfs at all, but giants crouching for the most part below the brightness of our sky; Perhaps I should mention that the original paper on this subject (Disney, 1976) has been heavily criticised by Allen and Shu (1979). Their criticism is based on a misinterpretation of my paper for which I suppose I must be partially to blame. The Fish-Freeman samples do not contain all the largest angular diameter galaxies in the sky (nor did I mean to imply that) because, particularly in Fish's case, where he was interested in a wide range of absolute magnitudes, some of the galaxies were chosen to be true dwarfs (such as M32) of small angular size•. But for any given total luminosity LT' the sample galaxies do have central surface brightnesses which give maximal apparent isophotal - 158 - diameters. Careful reading of both papers will show, I thin~, a more or less complete agreement as to the significant facts.

4. FUTURE WORK Clearly what we need are reliable distribution functions for the surface brightnesses and half-light radii of a much larger sample of galaxies than we presently possess, so that the selection effects can be clearly understood. Fortunately some of the technical means to tackle this ambitious programme are now coming to hand for the first time. I am thinking of digital area detectors like CCO's, of powerful image processing machinery of various kinds, and of photographic techniques like high contrast printing which will allow us to readily see the lower surface brightness features on plates. Much can be expected of these tools within the next decade. In collaboration with colleagues at the Royal Observatory, Edinburgh, and the Anglo-Australian Telescope, we are embarking on a modest programme of galaxy photometry at Cardiff, in particular using high contrast printing techniques to pick out galaxies, including dwarfs, with low surface brightnesses on the Science Research Council J survey. Space Telescope is not the' instrument to find many low surface brightness objects, but it will be able to distinguish compact dwarfs from stars. Since the fields of view are far too small for a useful survey, we can only hope for something useful to come from serendipity observations. Since the optical sky background (both from space and on the ground) is such a hindrance to low surface brightness observations, then perhaps we should move to other frequencies in our attempt to map out the local mass density? In principle we should not be much better of~ because the real problem is that surface brightness, at any wavelength, is to a first approximation independent of distance, and so if our own galaxy has a typical energy spectrum we shall be fighting a relatively bright sky in whatever band we try to work. Only if we can, in some sense, resolve the background, can we hope for a fundamental advance. For instance at 21 cm one can use spectral resolution to discriminate between local and extragalactic hydrogen. So far we appear to have had - 159 - very little success (e.g. ~hostak, 1~77) in searching for the extra galactic variety unassociated with visible galactic images, but I shall be interested to hear from the radio astronomers present what improvements we can expect from technical developments in the years to come. Personally I find the Einstein X-ray images of clusters (Jones et. al., 1979) to be very exciting. Loose clusters such as A 1367, A 2634 and A 2147 show, in addition to the X-rays from individual bright galaxies, islands of X-ray emission where no visible galaxies are to be seen. Are these the optical 'icebergs' or crouching giants I referred to earlier? Galaxy astronomers must hope that the X-ray background (Giacconi et. al., 1979) will prove to originate very largely from resolved sources for then, at last, we may have found in X-ray surveys the tool we need to identify at least some of that low-surface brightness population not easily detectable by any other means.

ACKNOWLEDGEMENTS I would like to thank Dr. Laustsen for permission to reproduce the two plates of N 3309/3311 which he took with the European Southern Observatory 3·6 metre telescope.

REFERENCES Abe11, G.O., 1975, Stars and Stellar Systems, Vol. IX, (Univ. Chicago Press), 610. A11en, R.J. and Shu, F.H., 1979, Astrophys. J., 227, 67. Arp, H., 1965, Astrophys. J., 142,402. de Vaucou1eurs, G., 1959, Handbuch der Physik, ~' 311. Disney, M.J., 1976, Nature, 263, 573. Faber, S.M. and Ga11agher, J.S., 1979, Ann. Rev. Astron. Astrophys., 12, 135. Fish, R.A., 1964, Astrophys. J., 139, 284. Freeman, K.C., 1970, Astrophys. J., 160,811. Giacconi, et. al., 1979, Astrophys. J., 234, L1. Heidmann, J., Heidmann, N. and de Vaucou1eurs, G., 1971, Mem. R. astr. Soc., 75, 85. Ho1mberg, E., 1975, Stars and Stellar Systems, Vol. IX, (Univ. Chicago Press), 123. - 160 - rlolmberg, E., 1969, Medd. Lunds Astr. Obs., 5, 305. Jones, et. al., 1979, Astrophys. J., 234, L21. Peeb1es, P.J.E., 1971, Physical Cosmo1o~y, (Princeton Univ. Press), 56. Shostak, 5., 1977, Astron. and Astrophys., 54, 919. Zwicky, F., 1957, Morphological Astronomy, (Berlin, Springer Ver1ag), 176.

Discussion

H. van Woerden: According to Dr. Disney the blue colours observed 1n NGC 1510 suggest that this elliptical galaxy may have been formed only recently; accretion of gas onto an old dwarf-elliptical could not well produce the observed colours. However, in a recent paper (T.G. Hawarden, H. van Woerden, U. Mebold, W.M. Goss and B. Siegman 1979, Astron. Astrophys. 76, 230) we have shown the following facts: -- I. The SBa galaxy NGC 1512, at 5' (~ 20 kpc) from NGC 1510, has a strongly disturbed morphology, suggesting tidal interaction with a mass> 109 M in NGC 1510. e 2. The system contains 10 10 M of neutral hydrogen, centred on NGC 1512 and extending over ~ 100 kpc ~iameter; the rotation indicates 2 x 101t M in NGC 1512. e 3. The emission line spectrum and blue colours of NGC 1510 can be well fitted by a young population of age ~ 10 7 • 5 years and mass ~ 10 7 • 5 M e formed from gas accreted onto an old (red) dwarf elliptical of ~ 109 • 5 M• e . . 4. The locat1on and mot1on of NGC 1510 relative to NGC 1512 allows accretion onto the former of 107 - 10 8 M of gas in 10 7 - 108 years. We claim that this accretion model ~s the only one consistent with all the observations. - 161 -

ENERGY DISTRIBUTION IN DWARF AND GIANT ELLIPTICAL GALAXIES

Francesco Bertola Osservatorio Astronomico, Padova

Observations with IUE have recently made available accurate energy distribution curves extended into the ultra violet"Jlidl give important information on the' different stellar conten~of dwarf and giant elliptical galaxies.

While in the spectral range from the infrared up to about 5500A the energy distribution of- elliptical galaxies is very similar, differentiations occur shortwards of this wavelength and become more and. more pronounced towards the ultraviolet.

In Fig. 1., which is adapted after Bertola et al. (1980) and Oke et al. (1980), one represented the energy distributions of NGC 3379, the standard elliptical galaxy, NGC 4486 (M87) , the well known active galaxy, and NGC 221 (M32) , a dwarf elliptical. Although available data for the giant elliptical NGC 4472 and for the bulge of M3l were not plotted since they are very similar to those of NGC 3379.

From 5500A to 2700A M32 is brighter than NGC 3379, reaching differences of as much as half magnitude. This is in agreement with the prediction of the CM relation for elliptical galaxies and is interpreted as a metallicity effect.

Quite different is the behaviour of M87, which becomes brighter than NGC 3379 only shortward of 3800A. It cannot be attributed to blanketing effect since the blanketing is the same for both ,galaxies from 3800A to 4500A. The cause of this UV excess in M87 is probably to be found in the fact that it is an active galaxy.

In the last portion of the UV spectrum the energy curve has, for all the galaxies considered, a sudden rise - 162 corresponding to the addition of a source which can be approx imated by a black body at 30,000oK. The phenomenon can be due to the presence of upper main sequence or horizontal branch stars. The different levels of the rising branch can be interpreted as due to different ratios of hot to cool stars in these galaxies. The ratio is very high in M87, intermediate in NGC 3379 and very low in M32. It should be noted that the extrapolation of the black body component towards longer wave lengths does not account for the UV excess of M87.

The combined energy distribution of NGC 3379 and NGC 4472 not corrected for aperture effect, has been compared with that derived from the ground for distant galaxies (z = 0.30-0.45) and reduced to the rest wavelength. It is a straightforward way to investigate the presence of a cosmological evolutionary effect. The two distributions are shown in Fig. 2. The difference between the two curves is just what is expected if the difference in colour between the inner region,as observed with rUE, and the total galaxy, as observed in the case of distant galaxies, is taken into account. ~~erefore colour changes due to different evolution of local and relatively distant galaxies are excluded within the accuracy of these measurements.

References

Bertola, F., Capaccioli, M., Holm, A.V., and Oke,J.B. 1980, Ap. J. Letters ...• Oke, J.B., Bertola, F., and Capaccioli, M., 1980, preprint. - 163 -

••••• NGC 4486 I~ . NGC 3379 . M 32 o 6'".,~.:-~ o·o~-': Y, ..,,'·....1 ••...... •.t .... rv • • •

2000 4000 6000

Fig. I- Energy distribution curves obtained with IUE and with the Multichannel Spectrometer on the Palomar Hale Telescope for NGC 4486, NGC 3379 and"M32. The zero point has been shifted in order to puperimpose the flat parts of the curves. The full and the dashed lines connect mean points.

-13.4

A".r.lIe HOC 337914472 G••o ... Z.0.30 to 0.45

log fl

3000 4000 6000 7000

Fig. 2 - Average energy distribution of NGC 3379 and NGC 4472 from rUE and 200" data, compared with the composite energy distribution of first rank cluster galaxies - 165 -

SESSION VI

(Chairman: B. Jones)

MORPHOLOGY AND EVOLUTION OF DWARF GALAXIES - 167 -

The Origin of Dwarf Ga1axies*

Bernard J.T. Jones Institute ·of Astronomy Mading1ey Road Cambridge CB3 ORA

Introduction

How do dwarf galaxies fit into our ideas on how galaxies form? When talking about galaxy formation, we speak of 'characteristic galaxy masses' and we have in mind large galaxies of the kind that are found in most pages of the Hubb1e Atlas (Sandage, 1961). Indeed, most observations of galaxies have concentrated on these large, impressive, systems and we have a situa tion where what is possibly the majority e1ement-'(the 'dwarfs') are forced into the background by an impressive minority. Yet dwarf galaxies may hold an important clue as to how galaxies in general formed: they may well be fossils from a time long past which have evidently taken a different evolu tionary track from the big galaxies.

What are dwarf galaxies?

Traditionally, when picturing dwarf galaxies we think in terms of the lesser members of the Local Group and other nearby groups of galaxies. To provide a focus for ideas, one might chose to consider 'dwarfs' as being galaxies fainter than MB = -16, and then consider classes of such objects: the dwarf e11iptica1s (dE), the Mage11anic Irregulars (Im), extragalactic HII regions, and gas rich dwarfs of the kind discussed at this meeting by Sancisi and by Norman. These definitions may however prove misleading. Firstly, it is really the mass function that we are interested in, not the luminosity function: the luminosity of such systems may be very sensitive to star forming activity and we should beware of 'optical chauvinism'. Secondly, it is possible that these different types of dwarf galaxies may on~ be different evolutionary phases of small galaxies, and thus the divi sion into types may obscure a more fundamental aspect of small galaxies.

One of the main difficulties we have to face when studying dwarfs is the problem of obtaining a fair sample of small galaxies. At present,

* Talk presented at the ESO Workshop on "Dwarf Galaxies: the need for coordinated space and ground based observations", Geneva May 12/13, 1980. - ,68 - objects are drawn from the Local Group and nearby groups of galaxies, from neutral hydrogen surveys (Lo and Sargent, 1979, Davies, 1980), or from searches for low surface brightness galaxies (the DDO catalogue for example). The fact that low surface brightness galaxies often have con siderable mass (Fisher and Tully, 1975) shows already that having low surface brightness is not synonymous with having low mass. The converse, that low mass galaxies have low surface brightness is arguable if we omit the extragalactic HII regions on the grounds that they are quite unlike 'ordinary' galaxies, though as stated above such an omission is likely to prove highly misleading. In Figure 1 (taken from Saito, 1979) the mean density of well-observed objects is plotted against their mass. The dia gram includes 'normal' galaxies, globular clusters, and some nearby 'dwarfs'. Notwithstanding the uncertain y in the values of the densities, the objects having masses in the range 106Me < M < 1010MG fall well below the line fitting the normal galaxies (and as it happens the globular clusters). If it were indeed the case, as this diagram suggests, that small galaxies fall below this line, it would be difficult to maintain the hypothesis that glob ular cluster sized objects are the basic building blocks of galaxies.

So, where are the small galaxies of 'normal' density? Evidence that they do not exist in substantial numbers comes from Holmberg's (1969) thorough survey of companions to Shapley-Ames galaxies. He finds a strong correlation of surface brightness with absolute magnitude

p = - log A + const. It seems unlikely that Holmberg missed small companions of higher surface brightness than this relation suggest, though one may more plausibly argue that this relationship is merely an upper envelope (though in that case, where are the lower surface brightness bright galaxies?) Independent evi dence in support of this comes from studies of the luminosity functions of galaxy clusters (see for example Godwin, 1976), though the range of absolute magnitudes measured is not as great as Holmberg's, and the selection effects that come into choosing the sample of galaxies to be measured are poorly defined.

The Formation of Dwarf Galaxies

There is no lack of suggestions as to what role dwarf galaxies have played in the general scheme of galaxy formation and evolution: the major problem is a lack of any systematic body of data with which to confront these ideas. The central problem with regard to the dwarf ellipticals is to explain how they have acquired such a relaxed-looking appearance (as judged from their radial light profiles) given such low stellar densities 3 3 ($ 10- Me /pc in their central regions). The standard way of dealing with this problem is to argue that these dE's were denser in the past and have expanded due to mass-loss. If dwarf ellipticals were no denser in the past, we have no way of understanding how stars could form at such low densities. Moreover, the fact that some dwarf galaxies have their own globular clusters suggests that they looked more like normal galaxies in the past.

Saito (1979) argues that supernov~ driven winds are particularly impor tant for galaxies having masses M < 10 Me and lying on the 'bright e1lip tica1s' line of Figure 1. Such small galaxies experience significant mass loss over short periods of time and should now be found in a state of - 169 -

globular clusters

2 N 44868 10 o N5846A o C'"""""' '(,) C- :ea> - 1 normal galaxies ...> -Cl) N185 Z o LLI C 10. 2 o DD o o o dwarf spheroidals •

10 10 MASS (Ma)

Figure 1 - 170 - expansion. It is not clear, however, why globular clusters have not experienced the same mass-loss. Silk and Norman (1979) argue that the likely mass-loss mechanism is the ram-pressure stripping by the hot gas in the halo of the galaxy. They argue for a population of gas-rich dwarfs in groups of galaxies; occasionally one of these falls into a large galaxy, forming a dwarf elliptical and providing a source of HI gas an important point of their scheme being that this explains the existence of HI gas in early type galaxies. The Silk-Norman scheme is based on a number of assumptions about the nature of the precursors of dE's: the so-called gas-rich dwarfs. It does however have the merit of explaining why some galaxies are evidently still accreting hydrogen gas, and of equal importance, tying together the various kinds of dwarf galaxies into a single scheme.

The Silk-Norman idea is not necessarily inconsistent with the older scheme discussed by Hodge and Michie (1969) and Innanen and Keenan (1973) who argued that the dwarf galaxies were formed with the Galaxy, and indeed Silk and Norman comment that at present we may be witnessing the tail-end of an initially more extensive primordial distribution of gas-rich dwarf galaxies and interga1actic clouds. An important question is: can we tell how recently the dE's of the Local Group were deprived of their gas? There is no evidence in these for recent bursts of star formation that might have been associated with their having had large quantities of gas in recent times. We might hope to find such evidence in integrated intermediate band colours, or by studying the co1our-magnitude diagrams of those systems that can be resolved into stars (the existence of 'blue stragglers' and a broad ened giant branch are perhaps the symptoms to look for). On the other hand, if the dE's of the Local Group lost their gas 1010 yrs ago in a violent event of the kind suggested by Saito (1979), we would hardly expect to see the remaining stellar system now.

Relevant observations of dwarf galaxies

Given our present state of knowledge (or lack of it) about the dynamical, physical and chemical properties of dwarf galaxies, any data at all is desirable! From the point of view of understanding the origin of these systems it would be useful to establish some key facts in certain areas. The following is a brief list of some interesting questions.

1. Chemistry: what is the composition of the gaseous and stellar com ponents of dwarf galaxies, and how dies it relate to morphology and size? Already there has been a considerable amount of work on emission lines from the extragalactic HII regions which has been reviewed at this meeting by Terlevich. There is less information on the underlying stellar population. Jones and Jones (1980) in their spectroscopic study of the Fornax cluster of galaxies found evidence for considerable variations in meta11icity among the dwarf galaxies with absolute magnitudes around -16.5. In their list, G72 and G43 are clearly as metalrich as many galaxies which are two magnitudes brighter. G139 shows strong Balmer lines in absorption, with Ha and [OIIJ in emission, and G150 has strong [OIIJ, [OIIU, and weak Ha on an otherwise low-metals K-type stellar continuum. These data will be published in the near future since they represent an optically homogeneous and complete sample. 2. Spatial Distribution: how are dwarfs located in space and what is their relationship to 'normal' galaxies? Here again we are hampered by the lack - 171 - of surveys. The distribution of DDO galaxies (which are not always dwarfs!) has been studied by Sharp, Jones and Jones (1978) who found evidence that they were anti-correlated with clusters and groups of galaxies, but never theless positively correlated with bright galaxies and with each other. Since the DDO galaxies are mostly gas-containing systems, the avoidance of groups may not be surprising. Of more relevance may be the studies of Reaves, Shaw and J. Jones of dwarfs in the Virgo cluster. These are generally dE's and there appears to be more of these per bright galaxy in the Virgo cluster than ln the field. If this impression is confirmed it will have an important bearing on ideas for the evolution of galaxies in clusters as well as for the origin of dwarf ellipticals. 3. Internal Dynamics: what is the mass-to-light ratio of dwarf galaxies of various kinds, and are dE's expanding? A considerable amount of work has been done in recent years on dwarfs having HI and/or optical emission lines in their spectra. The variety of dynamical behaviour is large, one extreme case having been discussed by Sancisi at this meeting. A completely unanswered question is 'what is the dynamical state of dE's?'. If follow ing Norman and Silk they are stripped hydrogen-rich galaxies, we might expect the remnant stellar system to be expanding at ~ 10 km/sec. Such an expan sion of course places a constraint on the lifetime of dE galaxies. That may not be a problem for the dE's of the Local Group, they may well have been stripped no more than a few billion years ago (see earlier comment on 'young' stars in these systems). However it poses a problem for the origin of the dE's in the Virgo cluster - it seems that they cannot be stripped gassy systems otherwise they would have dissolved a long time ago (their numbers are such that one would hesitate to contemplate the possibility that they are just now falling in).

How can ST help?

It is clear that there is much to be done from ground-based obse~va tions in the way of surveys and spectroscopy of individual objects, and that the main impact of ST will be in studying dwarf galaxies that it can resolve into stars. The sample of resolved dwarf galaxies will increase many times in size, and hence (subject to the obvious limitations in avail able observing time) ST will provide a direct measure of the stellar popula tions that make up these galaxies. The dwarfs of the Local Group are a rather restricted sample of the various types of dwarf galaxy, and so such a venture beyond the Local Group is of some importance. The ability to spatially resolve the central regions of dE's and so provide good light profiles will shed light on the kind of dynamical models that are relevant. It is impossible to predict where in the study of dwarf galaxies ST is likely to make the greatest impact, though its use as a tool in investi gating these objects can hardly be doubted. What is of central importance, however, is that ground-based facilities be used more fully now to increase our all-too-poor understanding of dwarf galaxies, and then used in a co ordinated effort to work with ST in its observational programs.

Acknowledgements

I wish to thank Prof. P.O. Lindblad for his invitation to join. this stimulating workshop, and ESO for their hospitality during that time. - 172 -

References

Davies, R.D., 1980, Phi1.Trans.R.Soc.Lond., A296, 407. ~ Fisher, J.R. and Tu11y, R.B., 1975. Astr. &Astrophys., 44, 151. ~ Godwin, J., 1976, Ph.D. Thesis, Oxford University.

Hodge, P.W. and Michie, R.W., 1969, Astron. J., 74, 587. ~ Ho1mberg, E., 1969, Ark.f.Astr., 5, 305. vv Innanen, K.A. and Keenan, D., 1973, J.Roy.Astr.Soc. Canada, 67, 248. ~ Jones, J.E. and Jones, B.J.T., 1980, Mon.Not.Roy.Astr.Soc., 191, 685. ~ Lo, K.Y. and Sargent, W.L.W., 1979. Astrophys.J., 227, 756. ~ Saito, M., 1979, Pub1.Astron.Soc. Japan, 31, 181 and 193. ~ Sandage, A., 1961, The Hubble Atlas of Galaxies, Carnegie Inst. (Washington)

Sharp, N.A., Jones, B.J.T. and Jones, J.E., 1978, Mon.Not.Roy.Astr.Soc., 185, 457. ~ Silk, J. and Norman, C., 1979, Astrophys.J., 234, 86. ~ - 173 -

Gas Rich Dwarfs

Colin Norman* + Bernard J.T. Jones

~Huygens Laboratory, University of Leiden, Netherlands Institute of Astronomy, Madingley Road, Cambridge.

At this meeting we have heard about the low surface brightness dwarf galaxies (Lo and Sargent, 1979; Sancisi, this meeting). They have masses in the range 106 - 109 MG, MaI/L ~ 1, low surface brightness and negli gible rotation. The origin, stability and structure of these systems, their accretion onto massive galaxies and their role in the galaxy forma tion process are discussed in two papers: Silk and Norman (1979) and Norman and silk (1980).

The basic model of galaxy formation consists of collections of bound gas clouds that cluster hierarchically with dissipation occurring in the core regions. The nature of the original building blocks depends little on which theory of galaxy formation we adopt. Extrapolation of the pre sent clustering correlation function suggests density perturbations of unit amplitude on scales of 108 - 109 Me. If we adopt isothermal fluctua tions, the relevant mass is the Jeans mass at recombination (~ 106n-! Me). The 'pancakes' theory provide a late injection of objects with masses in the range 106 - 10 9 ~. The late injection in the Pancake theory is a good thing since one of the main problems is to explain the longevity of those basic building blocks that still survivg and that we believe are the gas rich dwarfs. We need lifetimes of 10 - 1010 years.

We can understand these long lifetimes as a consequence of a balance between the rate of star formation and the energy input to the gas from supernovae. The condition that the clou~ radius should be greater than the radius of a supernova remnant puts a lower limit on the cloud mass

T -1 M 4 c;> and the requirement that the supernovae be not too frequent provides an estimate of the lifetime of the cloud t ~ 109• 7 MT -1 El •28 years. cl 7 4 51 These clouds (or gas-rich dwarfs!) can indeed exist for as long as 1010 yrs. The rate of chemical enrichment of the cloud is very low. - 174 -

We see evidence for significant accretion of gas in a number of early type galaxies: the 'spindle', NGC 4278 and NGC 1023. Silk and Norman suggest that the accreted gas is generally gas-rich dwarfs. The suggestion is then that the number density of such objects may have been significantly greater in the past than now, and this encourages specula tion about the evolution of radio sources (the fuel runs out at z ~ 0.5), the origin and distribution of QSO absorption lines and other evolution ary effects like the Butcher-Oemler effect. The general swallowing of large numbers of such objects by a galaxy provides a source of energy for a galactic wind that can strip future infalling dwarfs of their gas. The left-over stellar system dissolves in the galactic halo and any sur viving structure might be identified with dwarf spheroidal systems such as are seen in the Local Group.

References

Lo, K.Y. and Sargent, W.L.W., 1979, Ap.J., 227, 756.

Norman, C. and Silk, J., 1980, to be published.

Silk, J. and Norman, C., 1979, Ap.J., 234, 86.

Discussion

H. van Woerden: I find this paper full of stimulating ideas, but I wonder about its relation to observables. In particu lar, there is no evidence that NGC 1023 owes its neutral hydrogen to copious consumption of gas-rich dwarfs. With one exception, the vast amounts of HI around NGC 1023 do not correlate with dwarf galaxies. The distribution of HI around NGC 1023, and especially its velocity field, rather appear suggestive of a ring with a tidal tail, although I grant that it is unclear what galaxy should have caused the tidal effect. Moreover, the Lo-Sargent search has shown so few gas-rich dwarfs that I doubt that there would be enough of those objects for your purpose. - 175 -

Galaxy Formation By The Agglomeration Of Primordial Star Clusters+

T.W. Hartquist

Department of Physics and Astronomy, University College London; Harvard-Smithsonian Center for Astrophysics; Max Planck Institute for Physics and Astrophysics

Abstract

Galaxies may have formed by the agglomeration of smaller primordial star clusters. I investigate the role of two body relaxation processes in the evolution of such a protogalaxy. - 176 -

I. Introduction

Peebles and Dicke (1968) suggested that the first stellar systems had masses equal to the Jean's mass at recombination (~ 106 M) and that larger stellar systems formed by the C!) agglomeration of these urclusters. Mutual tidal interactions (Spitzer 1958), shock heating as an urcluster passed through a disk component of the larger system (Ostriker, Spitzer, and Chevalier 1972), and evaporation resulting from two body relaxation processes (Ambartsumian, 1938; Spitzer, 1940) could have acted to dissolve the urclusters. In recent years it has become evident that angular momentum plays little role in determining the structure of elliptical galaxies (Binney 1976), and it is uncertain whether rotation caused a disk component to form in them, and unless the urclusters had masses of roughly a few thousand solar masses and less the-y could not have '- evaporated in the first billion years after recombination. Here we consider an alternative picture most applicable to elliptical galax~es, in which the urclusters were destroyed by their mutual tidal stripping.

II. Tidal Stripping Time Scale

If a protogalaxy consists of Nu urclusters of mass M 1/3 u dist~ibuted uniformly to a radius of 2 Rg the root mean square speed of an urcluster is ~)1/2 = 30 km s-1 ( (R )-1/2( M )1/2 ( 1 ) 500 1 ipc 10~ MC!) We will assume that rotation is insignificant and that the velocity dispersion is comparable to v as may be the case for g elliptical galaxies. The root mean square speed of a star in an urcluster will be

)1/2 = 12 km s-1 Ru )-1/2( M (2) (10 pc 106 M

(3)

(cf. Lightman and Shapiro 1978) which is the time in which any energy gained will become distributed among all of the stars of the urcluster.

Comparison of Vu and Vg indicates that Vg » Vu for large galaxies with M 1011 M, R 10 kpc and that during g = <:> g = collisions the urclusters will not alter each others paths significantly. Then for collisions of short duration it is possible to calculate the heating of each urcluster by using the impulse approximation given by Spitzer (1958) 2 G H 3R 2 4 u u 6Er Cp) = 3 4 2 (4) P Vg where p is the collision impact parameter and G is the gravi tational constant. Spitzer found that this approximation is reasonable for collisions lasting up to about 3 tdu and that for large Lmpact parameters expression (4) greatly overestimates the heating. Hence, we assume that the energy gained by each ur cluster in a collision is given by

6E = 0 2 R P < u 6E (p) 2 R < < t ( 5) = r u --P 1.5 Vg du 0 1.5 V t < = g d u P

The rate at which energy is gained by an urcluster is then 2 1. 5 v t 3N G M 3R 2 d dE u 4 u u -3 v f g u ~ dt = 3 2 g 41TP dp 81TR 3 Vg 2 R g u

34 N_U )1/2 Rg )-5/2 Mu )5/2 x 5x10 erg s-1 ( 500 ( 1 kpc ( 6 10 Mo Max 10, 1 - (1.52vR~ (6) g du r - 178 -

We can get the rate at which an urcluster radius grows from 2 dE 3 G!'1u dRu dt = "5 ;z- dt (7) u and we find that an urcluster of initial radius R grows to an u infinite radius in a time of ~)-1 ~ 9 ( /2 ( Mu )-1/2 Rg )5/2 (R )-1 TS 4x10 yrs 500 106 M (1 kpc 10u pc o (8)

Ill. Mass Segregation

Inspection of equation (8) shows that for NM = 5x1010 M R )5/2 u u 0 11 10k~C and Ru = 10 pc that t s = 10 yrs ( which for present galactic radii far exceeds the age of the universe. If galaxies formed as a consequence of urclusters tidally stripping one another the more massive protogalaxies must have been denser than galaxies during the current epoch. It is well known that a considerable fraction of a galaxy's gravitational energy can be contained in an inner core of size less than 1 kpc. Energy transfer from the urclusters which survive the initial phases of the agglomeration could lead both to the formation of a dense galactic nucleus and to the re expansion of the bulk of the galactic material.

Dynamical friction (Chandrasekhar 1943) can cause the remaining urclusters to sink relative to the newly created homo geneous stellar field on a time scale of

8 ~)1/2 Mr )-1 Rg )+3/2 L ~ 2x10 yrs ( ( df 500 ( 10~ M 1 kpc <:> - 179 - where M r is the mass of one of the remaining urclusters, A is u roughly the ratio of the urcluster size to the galactic size, and f is the fraction of the mass contained in the stellar field. However, while migration can take place for a large range of f's it is clear that re-expansion to about 10 kpc can occur in the age of the universe only when the masses in the remaining ur clusters and in the stellar component are comparable. This is equivalent to requiring that galaxies which have re-expanded had dynamical friction time scales comparable to their ur cluster destruction times at an early stage in their evolution. From equations (8) and (9) we can derive the initial radius of a galaxy which substantially re-expanded.

= .05 kpc (10)

If we set Mu= M, r substitute (10) into (8), and require 9 that L ~ 4x10 yrs we can derive the mass that an urcluster S must have had if re-expansion was important.

u u M M (NuM )4/5 (R )3/5 (lnA)-1 (11 ) u = (;) 5x108 10 pc 10 10 Taking a galactic mass of 5x10 M,M(;) u must have been about 6 2x10 M. The initial radius of such a galaxy would have been (;) about 2.5 kpc and one would expect about half of its mass to be contained inside this radius at the present epoch.

IV. Smaller Galaxies

We do not consider the evolution of larger galaxies which were probably greatly affected by capturing other nearby objects. We will assume that all smaller galaxies collapsed to approximately the same density as did the 5x1010 M galaxies. (;) The dynamical friction time scale then scales as N 1/ 2 and the u urcluster destruction time scales as N 1/3. Hence dwarf u galaxies evolve in a different fashion. - 180 -

Specifically the number of two body relaxation times that the individual urclusters survive will be approximately (M /5X1010 M )-1/6 so that an elliptical galaxy of mass, M = ~ 0 g 10 M will have somewhat more angular momentum through two o body processes (Marchant and Shapiro 1977). More importantly the system of urclusters would have had time to develop a more strongly packed core-halo structure. If smaller galaxies did not collapse to as high densities as larger ones the effects of two body processes would have been even more pronounced, and observations of the rotation and the core-halo structure of dwarf elliptical galaxies may provide insight into the dynamics of galaxy formation.

References

Ambartsumian, V.A. 1938, Ann. Leningrad State Univ., No 22. Binney, J. 1976, Mon. Not. R. Ast. Soc., 177, 19.

Chandrasekhar, S. 1943, Ap. J., ~, 255.

Lightman, A.P. and Shapiro, S.L. 1978, Rev. Mod. Phys., 50, 437.

Marchant, A.B. and Shapiro, S.L. 1977, Ap. J., 215, 1.

Ostriker, J.P., Spitzer, L. and Chevalier, R. 1972, Ap. J. (Letters), 176, 51.

Peebles, P.J.E. and Dicke, R.H. 1968, Ap. J., 154, 891. Spitzer, L. 1940, M.N.R.A.S., 100, 396.

Spitzer, L. 1958, Ap. J., 127, 17. - 181 -

Stochastic Selfpropagating Star Formation in Magellanic Type Galaxies

J. V. Feitzinger Astronomisches Institut, Ruhr-Universitat Bochum, BRD

Beside the spiralarm density-wave theory the stochastic selfpropagating star for~ation theory (SSPSF, Gerola and Seiden 1978, Seiden and Gerola 1979, Seiden, Schulman, Gerola 1979) is an addition and / or an alternative for the explanation of the spiral arm phenomena in galaxies. Stochastic selfpropagating star formation is able to build up spiral arm filaments in slowly rotating galaxies. These short, string-like patterns clear~y show only little similarity with the familiar arms of spiral galaxies and can be found especially in late type galaxies. One typical example is the Large Magellanic Cloud; the youngest constituents (age <:3.5 107 yr) are lined up in short filaments (Schmidt-Kaler 1977). Using the SSPSF theory and the parameters o~ the LMC, especially the new rotation curve (Feitzinger 1980), model sequences for this stellar system are calculated (Feitzinger et al. 1980). The observations are in agreement with the basic predictions of the theory: It is possible to simulate spiral filaments with a common asymmetrical located center and irregular distributed small regions of star formation. The nonuniformity of star formation over the disk is well established (compare fig.3 Feitzinger and Schmidt-Kaler 1980). In many cases of late type galaxies it was re vealed with the help of photographs in different wave length regions that a given stellar system may be simul taneously a well streamlined barred spiral, a disorganized - 182 - rudimentary barred gal~ with spiral arm filaments, or irregular, depending on wether we study the structural features by the blue (young), yellow or red (old) stars or by the distribution of HII regions, shown by EUV pictures (in the case of the LMC compare Page and Carruthers 1978). Therefore the capability of taking photographs in EUV (or IH) with the Space Telescope will lead toanew understanding of the morphology and evolution of late type galaxies (with the subgroup dwarf systems) and to further testing of the SSPSF theory. One basic supposition for successful space observations is the existence of good ground-based observations of the objects in question.

Feitzinger J.V., 1980, in Photometry, Kinematics and Dynamics of Galaxies, Austin Conference, D.S. Evans Ed., p.435 Feitzinger J.V., bchmidt-Kaler Th., 1980, in Two Dimensional Photometry, ESO workshop, P. Crane, K. Kjar Eds., p. 249 Feitzinger J. V., Gerola H., Seiden P.E., 1980, in preperation Gerola H., Seiden P.E., 1978, Ap.J. 223,129 Page Th., Carruthers G.R., 1978, Far Ultraviolet Atlas of the Large Magellanic Cloud, NRL Report 8206

Schmidt-Kaler Th., 197'7, Astron. Astrophys. ~, 771 Seiden P.E., Gerola H., 1979, Ap.J. 233, 56 Seiden P.E., Schulman L.S., Gerola H., 1979, Ap.J. 232, 702 - 183 -

SUMMARY OF ST CAPABILITIES ON DWARF GALAXIES

S. di Serego Alighieri

Astronomy Division, Space Science Department of ESA, ESTEC, NOORDlVIJK, HOLLAND.

The main characteristics of the Space Telescope and of its first genera tion Scientific Instruments are listed in Tables 1 and 2. The main ad vantages over ground-based telescopes are the improvement in resolution (from I arcsec to at least 0.1 arcsec) and the extension of the observable wavelength range into the ultraviolet (down to 1200 R) and into the infra red. The expected impacts on the observations of dwarf galaxies may be summarized as follows: a) Mainly because of its higher resolution. the ST will be able to detect very faint point sources; the limiting detectable flux for such sources is related to the telescope parameters 1n the following way: !1-- (S/N < 10 ) f u SiN ~V s:y lim t > 1 hour where SiN is the required signal-to-noise ratio. D is the mirror dia meter, 8 is the angular size of the image, fsky is the sky background per unit solid angle and t is the total integration time. Between 3500 and 7000 R the sky background in space is lower by a factor of 2-4 than the one of a good night on ground. This roughly compensates for the loss in D over the largest ground telescopes. The increased resolution will allow observations to go some 3 magnitudes fainter in the optical region. Further improvement is expected because of the lower UV and IR sky background and because of the larger wavelength range available. As an example it has been calculated that the FOC

will be able to observe a star of mli = 28 with SiN = 7.5 in la hours (Reference I). It will then be possible to do photometry on more dis tant and/or intrinsically fainter stars in dwarf galaxies (for varia bility, HR diagrams, etc.). b) By the same arguments no basic improvements are expected on the limiting surface brightness in the optical region. Since, moreover, the avail able fields of view for imaging with ST are small (see Table 2), and the ob~erving time difficult to get, no advantages are foreseen in the search for dwarf galaxies. Nevertheless the possibility of detecting dwarf galaxies in serendipity images is open. c) Resolution of galaxies into stars will be possible la times further than from ground. This will allow a better study of known objects to be made. d) The nuclei of dwarf galaxies will be resolved la times further (Virgo cluster ?). e) Low- and high-resolution UV spectroscopy will allow us to study in more detail the chemical composition and the kinematics of the gas and dust in dwarf galaxies. - 184 - f) A further improvement in the detailed study of dwarf galaxies (including intergalactic HII regions) will be provided by the high angular resolu tion spectroscopic facility of the FOC (Reference I).

In conclusion, the review of 5T capabilities opens the possibility of im proving our understanding on known dwarf galaxies, but also suggest obser vational work to be done now from the ground•.

REFERENCE.

\. F. Macchetto, "The Faint Object Camera", in the Proceedings of the E50 Workshop on 'Two dimensional photometry', Noordwijkerhout, 21-23 November 1979. - 185 -

TABLE 1

MAIN CHARACTERISTICS OF THE SPACE TELESCOPE

f/number 24

Diameter 2.4m

Optical design Ritchey-Chretien Cassegrain 2 Useful collecting area 40000 cm

Scale 3.S8 arcsec/nnn

Field of view diameter 18 arcmin

Spectral range I ISO ~ -I nnn

Wavefront error A/13.S at 6328 ~

(design goal: A/20)

Radius of 70% encircled 0.1 arcsec

energy from a point source at 6328 ~

Stray light 2 in directions with 23 TIlv/arcsec

-> SOO from the sun > 70 0 from the earth limb > ISO from the full moon

Pointing accuracy 0.01 arcsec

Pointing stability 0.007 arcsec

Orbit altitude SOO km

inclination 28~8

period 9S min. TABLE2

MAINCHARACTERISTICSOF THE FIRST SCIENTIFIC INSTRUMENTSFOR THE SPACETELESCOPE

Parameter Unit WIDEFIELD/ FAINT OBJECT HIGH RESOLUTIONI HIGH SPEED FAINT OBJECT PLANETARYCAMERA SPECTROGRAPH SPECTROGRAPH PHOTOMETER CAMERA f/number 12.9 / 30 N.A. N.A. N.A. 96 / 48 / 288

Picture format pixel 1600 x 1600 I x 512 I x 512 N.A. 512 x 512

Pixel size arcsec O. I/ 0.043 N.A. N.A. N.A. 0.022/0.044/0.007

Field of view arcsec 2 1 160 x 160/68 x 68 o 0.1 - 4.3 1 0.25xO.25 - 2 x 21 0 0.4-1-10 l1xl1/22x22/3.8x3.8

Wavelength 2 4 3 ~ A/!J."A N.A. 10 - 103 105-2xl0 -2xl0 3 N.A. 50 - 100 - 2 x 10 resolution CP (]' Spectral range R 1150 - I 1000 I ISO - 9000 I 100 - 3200 1150 - 9000 1150 - 7000 m Dynamic range v 8.5-28 / 7.5-27.5 6 - 25 - 17 4 - 24 21-29/22.5-29/I 8.5-28 Photometric accuracy % '" 1 0.2 I '" I

Shortest 0.05 25 0.016 I 100 exposure time msec 100

Principal J.A. Westphal R.J. Harms J.C. Brandt R.C. Bless F. Macchetto Investigator and California Inst. Univ. of Calif. Goddard Space Univ. of (Project Scientist) Lead Institution of Technology at San Diego Flight Center Wisconsin European Space Agency