Cool Dense Gas in Early-Type Galaxies

COOL DENSE GAS IN EARLY-TYPE GALAXIES

C. HENKELMax-Planck-Institut fur Radioastronomie, Auf dem Hugel 69, D-53121 Bonn, Germany

T. WIKLINDOnsala Space Observatory, S-43992 Onsala, Sweden

(Received 11 June, 1997)

Abstract. CO observations have shown that many lenticular and elliptical galaxies contain significantamounts of cool dense gas. This review summarizes the observational results related to the neutralgas phase and presents a systematic comparison with other interstellar and stellar data. The discoveryof very dense molecular gas in the nuclear regions of early-type galaxies, the possible existence of adust component neither seen optically nor in CO, internal inconsistencies of cooling flow scenarios,the origin of the cool gas, the presence of massive stars, aspects of galaxy evolution, and possibilitiesfor future research are discussed in the light of the new data.

With increasing distance, our knowledge fades,and fades rapidly.

(E. Hubble, The Realm of the Nebulae, 1936)

Table of Contents

1. Introduction2. Morphological Types3. The Molecular Data4. Early-Type Galaxies with Cool Dense Gas: General Results

4.1. Some Useful Quantities4.2. A Cautionary Note4.3. Lenticular Galaxies4.4. cD’s, Elliptical, and Dwarf Elliptical Systems4.5. The Very Dense Gas Component

5. Early-Type Galaxies with Cool Dense Gas: Individual Sources5.1. NGC 404, NGC 3593, NGC 4691, Three Active S0 Galaxies5.2. Two Dwarf Elliptical Companions of the Andromeda Galaxy5.3. The Polar Ring Lenticular NGC 26855.4. The Giant Elliptical Radio Galaxy Cen A5.5. NGC 1275, a Cluster Dominant cD Galaxy5.6. Joining...5.7. ... and Leaving the Family of Early-Type Galaxies

6. The Neutral Gas Phase: Impact and Relevance for a Better Understanding of Early-Type Galaxies6.1. Dust in Elliptical Galaxies

6.1.1. Extended Dust Lanes6.1.2. Dust from the Nuclear Regions

6.2. The Radio Continuum6.3. The Warm Gas Component6.4. Hot Versus Cool Gas: Cooling Flows

6.4.1. Individual Galaxies with X-Ray Halos

Space Science Reviews 81: 1–105, 1997.c 1997 Kluwer Academic Publishers. Printed in Belgium.

2 C. HENKEL AND T. WIKLIND

6.4.2. Cluster Cooling Flows6.5. On the Origin of the Cool Dense Gas in Early-Type Galaxies6.6. Rotation Curves and M=L Ratios6.7. Early-Type Galaxies and Massive Stars

7. OutlookA. Appendix

A.1. CO Beam Filling FactorsA.2. Star Formation RatesA.3. Acronyms

1. Introduction

Studying the interstellar gas of early-type galaxies has proved to be a fertile area forresearch into fundamental astrophysical processes. This includes the physical andchemical properties of the interstellar medium (ISM), star formation, the evolutionof galaxies, the triggering of nuclear activity, the interaction between various gascomponents, and the determination of dynamical masses.

For more than half a century, it has been known that at least some of the early-type galaxies are not gas-free inert stellar systems but contain cold dust and warmionized gas. While the nature of the nuclear jet in the central Virgo elliptical M 87(Curtis, 1918) was not directly understood, the existence of an ISM was inferred bythe presence of dust lanes (Hubble, 1936) and the detection of [O II]�3227 emission(Mayall, 1936, 1939). Other interstellar gas components were found much later.With instruments aboard an Aerobee rocket, Byram et al. (1966) detected X-rayemission associated with M 87. 21 cm line emission from neutral hydrogen (HI)was first reported by Lewis (1970) from the lenticular NGC 1291; the first HI

absorption profile was obtained by Roberts (1970) towards the compact nuclearradio continuum source of the giant elliptical Cen A. The last major interstellarcomponent, the molecular gas, was detected by Gardner and Whiteoak (1976a,b, 1979) observing absorption lines of hydroxyl (OH) and formaldehyde (H2CO)towards Cen A. Searches for the more widespread CO molecule date back tothe same decade (e.g., Johnson and Gottesman, 1979). The first successful COmeasurements were reported by Verter (1985) and Stark et al. (1986) for the early-type disk galaxies NGC 4438 and NGC 7371, by Wiklind and Rydbeck (1986) forthe dwarf elliptical system NGC 185, and by Phillips et al. (1987) for Cen A. TheseCO detections, the large number of follow-up observations, and the success of theInfrared Astronomical Satellite (IRAS) with its direct view of the interstellar dustcomponent (e.g. Jura et al., 1987; Bally and Thronson, 1989; Knapp et al., 1989)provide the motivation for this review.

Although comprising only a small percentage of the total mass of a galaxy, thecool ISM is of prime importance because it provides the fuel which is needed totrigger star formation and, in some cases, to feed an active nucleus. Roberts et al.(1991) compiled a list of 467 early-type galaxies (including Sa-spirals) with radio,infrared, optical, and X-ray properties. Lees et al. (1991) display global properties

COOL DENSE GAS IN EARLY-TYPE GALAXIES 3

of early-type galaxies in which CO has been searched for. A significant number ofthe CO detections discussed below (e.g., Wiklind et al., 1995; Knapp and Rupen,1996) are too recent to be included in these catalogs. Complementing reviewsmainly focused on stellar dynamics and optical and X-ray data (e.g., Athanassoulaand Bosma, 1985; Schweizer, 1987; Kormendy and Djorgowski, 1989; Sarazin,1990; de Zeeuw and Franx, 1991; Fabian et al., 1991; Capaccioli and Longo, 1994;Fabian, 1994; Ferguson and Binggeli, 1994), we summarize observational resultsrelated to the cool dense ISM. As a favor to the reader and as an attribute to thesize of the review, each section is kept self-contained; consecutive reading is notrequired. Acronyms used throughout the article are given at the end of the text(Section A.3).

After an introduction to morphological properties in Section 2, Section 3 pre-sents the CO data base that allows a quantitative analysis of the molecular compon-ent. In Section 4, we summarize general results obtained from studies of variousmolecular species, HI, and IRAS data. Selected individual objects are portrayedin Section 5, while Section 6 highlights the relevance of the cool dense gas forour general understanding of early-type galaxies. Sections 6.1–6.4 relate the cooldense ISM to radio continuum emission, optical extinction, optical line emission,and X-ray (‘cooling flow’) data. The latter is emphasized to acknowledge its impactfor neutral cloud physics, cooling flows, and other areas of research. The origin ofthe cool dense gas, measured rotation curves, and mass-to-light ratios are discussedin Sections 6.5 and 6.6. Arguments, favoring the sometimes questioned presence ofyoung massive stars, are outlined in Section 6.7. In the final section, we summarizesome of the most important results and suggest avenues for future research.

2. Morphological Types

Early-type galaxies include quite a heterogeneous group of objects: Lenticular orS0 galaxies possess both a dynamically hot stellar component, the bulge, and a coldcomponent, the disk, which does not contain spiral arms. Ignoring large diameter(� 10 kpc) low surface brightness (� 25min arc sec�2) galaxies, the nominallydiskless systems can be subdivided into ‘cD’ and ‘D’ galaxies (MB � �22M ),into ‘giant’, ‘intermediate’ (�20 :M5 < MB < �18 :M5), and ‘classical dwarf’,‘compact’ or ‘M 32 - type’ ellipticals (summarized hereafter by ‘E’), into ‘diffuse’or ‘bright’ dwarf ellipticals (hereafter ‘dE’), and into ‘extreme dwarf ellipticals’ or‘dwarf spheroidal systems’ (dSph, MB > �14M ; cf. Sandage and Binggeli, 1984;Bender, 1992; Schombert, 1992; Gallagher and Wyse, 1994). A few dS0 galaxies(e.g., Ferguson and Binggeli, 1994) are also known. For the light distributions,see, e.g., Young (1976), Jaffe (1983), Sandage and Binggeli (1984), Bender andMollenhoff (1987), Schombert (1992), Ferguson and Binggeli (1994), Graham et al.(1996), and Prugniel and Simien (1997). Elliptical and dwarf elliptical galaxiesare morphologically distinct (e.g., Kormendy and Djorgowski, 1989; Bender et al.,


Figure 1. Suggested morphological classification scheme for luminous early-type galaxies (Kormendyand Bender, 1996).

1992). With decreasing luminosity, elliptical galaxies show a higher optical surfacebrightness (they become more compact), while dE’s become more diffuse and donot follow the de Vaucouleurs r1=4 radial light distribution (e.g., Kormendy, 1977;Davies et al., 1983; Wirth and Gallagher, 1984; Binggeli et al., 1988; Fergusonand Binggeli, 1994). E’s can also be subdivided into coreless rotationally flattened‘disky’ ellipticals with seemingly low radio and X-ray luminosities for a givenLB (for their possible relation to S0 galaxies, see Figure 2), into optically moreluminous and predominantly unisotropic ‘boxy’ ellipticals with cuspy cores andnumerous globular clusters, and into ‘irregular’ and ‘elliptical’ E’s (Bender, 1992;Bender et al., 1992; Kormendy et al., 1996; Kissler-Patig, 1997). In this case theparametrization follows the isophotal shape parametera4 that gives the amplitude ofthe fourth cosine coefficient in the Fourier series describing the isophote deviationsfrom the best fitting ellipse.

E’s populate a ‘fundamental plane’ in the three-dimensional space defined bythe central projected velocity dispersion, a linear scale (e.g., the ‘core’, ‘half-light’or ‘effective’ radius), and a surface brightness (averaged over the area definedby the core, half-light or effective radius). Dwarf elliptical galaxies populate adifferent plane in this three-dimensional space, though there is some intersection(Bender et al., 1992; Peterson and Caldwell, 1993). Dwarf ellipticals appear tobe more closely related to dwarf irregulars (dIrr) than to ‘normal’ ellipticals orto the bulges of disk galaxies (e.g., Kormendy 1987a; but see Kritsuk, 1997). Themorphological transition from E to dE is atMB � �18M but there is overlap. Manyof the more luminous dE’s (MB <��16M ) are ‘nucleated’, possessing a central starcluster sometimes reaching Mv = �12M (e.g., Caldwell, 1996). Among systemsof low luminosity (MB � �17M ), there are truly intermediate types that may bein a transitional stage from dIrr to dE or vice versa (Sancisi et al., 1987; Priceand Gullixson, 1989; Carignan et al., 1991; Sandage and Hoffman, 1991; Lo et al.,1993; Sandage and Fomalont, 1993; Kormendy and Bender, 1994; Skillman, 1996).


Even though the basic properties of the various morphological types are welldefined, the classification of early-type systems remains ambiguous. This is illus-trated in Table I (see also Table 8 of Knapp et al., 1989), where Hubble types fromfour different catalogs are displayed. Perhaps the most prominent case is the proto-typical minor axis dust lane elliptical Cen A that is classified as a lenticular in theRC3 (de Vaucouleurs et al., 1991). Many effects can lead to contradictory classific-ations; some are mentioned here: (1) Spirals with faint disks or dominant bulges aresometimes taken as S0 or E galaxies. (2) The prejudice that early-type galaxies donot possess an ISM has led to a classification of some ellipticals as S0’s and of somelenticulars as spirals. (3) Many ellipticals contain weak stellar disks and have radioproperties similar to those of S0 galaxies (Bender et al., 1989, 1992). Disky E’s maythus only differ from lenticulars by, in the average, lower disk-to-bulge ratios (e.g.,Rix and White, 1990; Michard, 1994; Kormendy and Bender, 1996; Scorza andBender, 1996), while the boundary between both Hubble types is not well defined.The classification then strongly depends on inclination: face-on objects tend to beidentified as ellipticals while edge-on objects are more likely classified as S0 (e.g.,Bender et al., 1992; Jørgensen and Franx, 1994). (4) E’s having strong isophotaltwists and S0’s with small inner bars can show a similar morphology (Nieto andLagoute, 1992). (5) Some amorphous I0 systems, having accreted large amountsof gas, may hide former early-type galaxies (Price and Gullixson, 1989). (6) Themerging of gas-rich spirals can lead to the formation of elliptical galaxies which,initially, cannot be properly classified because large amounts of dust obscure thedirect view. Thus many ultraluminous IRAS galaxies may be direct precursorsof ellipticals (cf. Djorgovski and Santiago, 1992; Kormendy and Sanders, 1992;Mamon, 1992; Bender, 1996).

In the following, we focus our attention on lenticular and (dwarf) ellipticalgalaxies that are not ultraluminous in the infrared (LFIR <� 1011 L�). Cool gasproperties of distant mergers that may form giant ellipticals have been discussedelsewhere (e.g., Henkel et al., 1991; Solomon et al., 1992b). In critical cases withconflicting classifications, we have not only used the major catalogs but also morerecent information in the literature. Despite this cautious approach, a certain biasmay be unavoidable.

3. The Molecular Data

It is the large number of recent CO detections in the ISM of early-type galaxiesthat provides the starting point for this review. CO is the molecular tracer that ismost readily observed: It is one of the most abundant species ([CO]/[H2] � 10�4)in local galactic interstellar clouds; its electric dipole moment is extremely small(0.11 D); radiative rates are thus slow and competing collisonal excitation rates arefast enough to excite the CO mm-wave transitions at all relevant particle densities


Figure 2. Molecular emission from the lenticular galaxy NGC 4526. (a) A 12CO J = 1� 0 spectrumfrom the IRAM 30-m telescope obtained with an angular resolution of 2100 . (b) The correspondingspectrum from the NRAO 12-m telescope (Sage and Wrobel, 1989) with a beam size of 5500. Acomparison of line intensities indicates that the source size is < 2000.

(n (H2)>100 cm�3) to temperatures well above the level of the 2.7 K microwavebackground.

Including the central regions of Abell galaxy clusters, CO observations werereported for more than 200 early-type galaxies. The surveys are based on early-typesamples surpassing a minimum IRAS flux (Sage and Wrobel, 1989; Lees et al.,1991; Huchtmeier and Tammann, 1992; Wiklind et al., 1995; Knapp and Rupen,1996), on a combination of IRAS fluxes and HI intensities (Thronson et al., 1989),on a combination of IRAS fluxes and the presence of optical dust lanes (Wiklindand Henkel, 1989), on the absence of FIR and the presence of X-ray emission(Braine et al., 1997), on proximity and estimated flow rate of an associated clustercooling flow (Braine and Dupraz, 1994), on X-ray absorbing column density andsource size (Antonucci and Barvainis, 1994) or on optical criteria (Bregman et al.,1992). A CO survey of polar ring galaxies (Galletta et al., 1997) has also beenpublished. Among the IRAS selected objects, the IRAS flux density catalog ofKnapp et al. (1989) plays a major role. Only the Antonucci and Barvainis (1994),Braine and Dupraz (1994), and the Braine et al. (1997) samples did not lead to COdetections. Some of the surveys contain sources that do not meet the predefinedcriteria and the understanding of what an early-type galaxy is differs from authorto author. Thus the entire sample of observed sources is quite heterogeneous andincomplete. Not surprisingly, this also holds for the subsample of �70 detected


galaxies (38% of the reported targets with cz <7000 km s�1). These sources havebeen selected after carefully examining the Hubble type (there are approximately30 E and 40 S0 galaxies) and the significance of each reported CO detection. TablesI and II display some properties of those�70 early-type galaxies in which CO hasbeen detected. Tentative detections and negative results are also listed (footnotes(f) and (g) of Table I). Following the table captions, source selection criteria andcontroversial CO data are discussed in more detail. Main results are summarizedin Section 4.

The columns of Table I contain the following information:

Column 1: Source name. Entries are identified, in order of priority, by ‘NGC’,‘UGC’, or ‘IC’ designations.

Column 2: The Hubble type favored by us (mostly taken from Knapp et al., 1989;otherwise, see footnote (a)). The index ‘p’ was omitted, since all displayed early-type galaxies contain significant amounts of molecular gas and are thus ‘peculiar’according to the traditional point of view.

Column 3: Hubble types taken from Sandage and Tammann (1987; RSA),de Vaucouleurs et al. (1991; RC3), Burstein et al. (1987), and Nilson (1973) orLauberts (1982).

Columns 4 and 5: R.A. and Dec. coordinates (epoch: 1950.0), mostly takenfrom the RC3.

Column 6: Radial velocities (cz) in km s�1. Whenever possible, velocities frompublished CO profiles were taken.

Column 7: +: detection towards one or a few positions; M: map.Column 8: References to molecular data which are given in footnote (b).

Table II displays molecular masses and surface densities, HI masses, S100�m

/S60�m flux ratios, and infrared, blue, and X-ray luminosities for a sample ofsources that is identical with that of Table I. Since deviations between valuesobtained by different authors often surpass given statistical errors, we did notattempt to provide standard deviations. To investigate differences in independentresults, see the references given in Table I for the molecular data and Huchtmeierand Richter (1989) for HI. The columns contain the following information:

Column 1: Source name. See Table I for details.Column 2: Distance in Mpc (DMpc). Most distances were estimated correcting

for the solar motion relative to the Local Group, for the peculiar motion of theLocal Group, and for the Virgocentric flow according to model 3.1 of Aaronsonet al. (1982). H0 = 75 km s�1Mpc�1. Galaxies belonging to the Virgo and Fornaxclusters were assumed to have distances of 18 and 20 Mpc, respectively. ForNGC 185 and NGC 205, we took the distance to M 31, D = 0:7 Mpc (e.g.,Sandage and Tammann, 1987). For NGC 404, NGC 3077, and NGC 3593, and


Cen A, D = 10, 3, 7, and 4 Mpc. (see Becker et al., 1989; Wiklind and Henkel,1990, 1992a; Israel, 1992; Rydbeck et al., 1993).

Column 3: The logarithm of the total molecular mass, MH2 , in units of M�,obtained from Equation (2) in Section 4.1. In those cases, where maps have beenobtained, MH2 may be larger than the values derived from Columns 2, 4, and 5.

Column 4: The molecular surface density, �H2, in units of M� pc�2, for thecentral region of the galaxy. The surface densities were calculated with �H2 =� ICO (1�0) and � = 3:02M� pc�2/K km s�1. ICO (1�0) is the integrated CO J =

1� 0 line intensity in units of K km s�1.Column 5: Angular size, �B (in arcsec), of the region for which a molecular

surface density is given in Column 4.Column 6: Logarithm of the HI mass in units of M�.Column 7: Dust color temperatures derived from S100�m and S60�m data (Knapp

et al., 1989; IRAS, 1989). The spectral energy distribution of a one temperaturecomponent is described by a modified Planck function, S� = �� B� (Td), with theemissivity of the dust grain population, �� , the Planck function, B� , and the dusttemperature, Td, at frequency �. �� / �� , with � � 1 for � <� 100�m and � � 2 for� >� 1000�m (e.g., Wiklind and Henkel, 1995). For � = 1 (as assumed in TableII), ratios of S100�m/S60�m= 6:05, 3.20, 2.04, 1.49, 1.13, and 0.93 correspond todust color temperatures of 25, 30, 35, 40, 45, and 50 K, respectively.

Column 8: Logarithm of the infrared luminosity (LFIR/L�), derived from theIRAS fluxes (Knapp et al., 1989; IRAS, 1989) for the interval 6�m< � <400�m,extrapolating the emission beyond 12 and 100�m and assuming a grain emissivityproportional to � (i.e., � = 1; cf. Wouterloot and Walmsley, 1986).

Column 9: Logarithm of the blue luminosity (LB /L�) obtained from the RC3Bo

T magnitudes that account for galactic and internal absorption and for redshift (deVaucouleurs et al., 1991). Only in the case of NGC 5953 the blue luminosity is takenfrom Wiklind and Henkel (1989). The absolute blue solar magnitude is 5 :M47. Toconvert log (LB /L�) to absolute magnitudes, use MB = �2:5 log (LB /L�) + 5.47mag. Blue absolute magnitudes separating different classes of elliptical galaxies(–18 :M0, –18 :M5, –20 :M5, –22 :M0; see Section 2) correspond to log (LB /L�) valuesof 9.4, 9.6, 10.4, and 11.0, respectively.

Column 10: Logarithm of the X-ray luminosity (LX /L�), taken from Robertset al. (1991) and Fabbiano et al. (1992).

Column 11: References for Columns 3–6 and 10.


Details of source selection relevant for Tables I and II

The number of reported detections is larger than the number of sources displayedin Tables I and II. Our definition of ‘(tentative) detection’ is more conservative thanthat of most authors reporting new results (cf. Wiklind and Henkel, 1989; Sofueand Wakamatsu, 1993; Knapp and Rupen, 1996). The polar ring galaxies detectedby Galletta et al. (1997) are not included because of uncertainties in the Hubbleclassification. Tentative detections providing useful targets for future surveys aregiven in footnote (f) of Table I. Non-detections are summarized in footnote (g). Inthe following, controversial results are briefly described.

Wiklind et al. (1995), Knapp and Rupen (1996) or Braine et al. (1997) do notconfirm reported detections of NGC 1052, NGC 2534, and NGC 4472, (Huchtmei-er et al., 1988, 1994; Wang et al., 1992a). Sage and Wrobel (1989), Thronson et al.(1989), Lees et al. (1991), and Young et al. (1995) do not confirm the detectionof NGC 5363 reported by Sanders and Mirabel (1985). Because of possible clas-sification errors (compare Jackson et al., 1989; Gordon, 1990; Li et al., 1993) orconflicting results (Lees et al., 1991; Huchtmeier and Tammann, 1992), NGC 3928and IC 3370 have also no entries in the main parts of Tables I and II. The COspectrum displayed by Sage and Wrobel (1989) towards NGC 4526 (for a negativeresult, see Thronson et al., 1989) has been confirmed by us. Both spectra are shownin Figure 2. The discrepancy in CO radial velocity of the S0 galaxy NGC 5866(Wiklind and Henkel (1989) report 400 km s�1� vLSR�750 km s�1; Taniguchi et al.(1994) find 700 km s�1� vLSR�1000 km s�1) is resolved by a third measurement:The data of Wiklind and Henkel mainly show the approaching, the spectrum ofTaniguchi et al., displays the receding part of the galaxy’s disk and bulge. A similareffect may explain the discrepant velocities obtained for NGC 3168 (see Taniguchiet al., 1994). A few Blue Compact Dwarf galaxies with elliptical isophotes andsome amorphous I0 galaxies are also part of the sample. These are either dwarfellipticals in a very active phase of their evolution or early-type galaxies havingaccreted large amounts of interstellar or intergalactic matter.


Table IEarly-type galaxies detected in CO (see Section 3 for further information)

(1) (2)a (3) (4) (5) (6) (7) (8)Source Hubble Type �1950 �1950 vLSR CO Ref.b

NGC 83 E -, E, E, E 00 18 46.6 +22 09 30 6350 + 43, 45NGC 185 dE dE, E, -, dE 00 36 12.0 +48 03 50 �250 M 3, 32, 47NGC 205 dE S0/E, E, E, E 00 37 38.7 +41 24 44 �250 + 12, 46NGC 404 S0 S0, S0, E, E/S0 01 06 39.3 +35 27 10 �50 M 12, 13, 14

18, 20, 23, 44NGC 759 E -, E, E, E 01 54 52.8 +36 05 57 4650 M 43, 47UGC 1503 E -, E, -, E 01 58 24.2 +33 05 19 5100 + 43NGC 807 E -, E, -, E 02 02 03.2 +28 44 57 4700 + 43NGC 855 dE -, E, E, E 02 11 10.8 +27 38 39 600 + 23, 43IC 1830 SB0a -, S0, -, SB0a 02 36 52.0 �27 39 36 1400 + 24UGC 2456 SB0 -, SB0, -, SB0 02 56 49.8 +36 37 21 3600 + 7, 14NGC 1291c SB0a SBa, SB0, -, SB0a 03 15 28.0 �41 17 24 800 + 24, 45NGC 1275 E E, p, -, p 03 16 29.9 +41 19 55 5250 M 9, 10, 34, 38, 44NGC 1316 E Sap, S0, E, S0 03 20 47.0 �37 23 06 1700 + 14, 36NGC 1326 SB0a SBa, SB0, -, S0a 03 22 01.8 �36 38 28 1350 M 14, 24, 42NGC 1386 S0a Sa, SB0, -, S0a 03 34 52.0 �36 09 48 900 + 14, 42UGC 2836 S0 -, S0, -, E/S0 03 40 38.2 +39 08 16 4950 + 14NGC 1482 S0a -, S0, -, Sa 03 52 27.0 �20 38 54 1850 + 44NGC 1546d S0a Sc, S0?, -, Sa? 04 13 32.1 �56 11 06 1250 M 32NGC 1819 SB0 -, SB0, -, SB0 05 09 06.5 +05 08 26 4450 + 13, 44NGC 1947 E S0, S0, -, S0 05 26 28.0 �63 48 06 1200 + 36NGC 2328 E/S0 -, S0?, -, E/S0 07 01 00.4 �41 59 37 1100 + 23, 43NGC 2320 E -, E, E, E 07 01 49.5 +50 39 24 5550 + 43NGC 2685 S0 S0, SB0, -, p 08 51 41.3 +58 55 30 900 + 13, 19, 39, 40, 44NGC 2768 E S0, E, E, E/S0 09 07 45.2 +60 14 40 1350 + 23, 43, 45NGC 3032 S0 Sa, S0, -, S0 09 49 14.1 +29 28 20 1550 + 12, 13, 44NGC 3077 I0 I0, I0, -, I0 09 59 21.9 +68 58 33 0 M 6, 44NGC 3166 S0a Sa, S0, -, S0a 10 11 09.3 +03 40 25 1000 + 14, 39NGC 3265 E -, E, -, E? 10 28 19.1 +29 03 13 1430 + 21, 23UGC 5720 Im/E -, Im, -, p 10 29 22.9 +54 39 23 1050 M 29, 45NGC 3593 S0a Sa, S0a, -, S0 11 12 00.2 +13 05 24 650 M 13, 14, 31, 44NGC 3597 E -, S0, -, S0 11 12 13.4 �23 27 19 3500 + 43NGC 3656 E -, I0, -, p 11 20 50.5 +54 07 08 2900 + 29, 43NGC 3665 E S0, S0, -, E/S0 11 22 00.9 +39 02 16 2100 + 12NGC 3682 S0a -, S0a, -, S0a 11 24 46.2 +66 51 56 1600 + 14UGC 6877 dE/S0 -, S0, -, - 11 51 38.5 +00 24 53 1150 + 28NGC 4125 E -, E, E, E 12 05 34.8 +65 27 18 1300 + 43NGC 4138 S0 Sa, S0, -, S0 12 06 59.3 +43 57 57 1050 + 13, 44NGC 4293 S0a Sa, SB0a, -, SB0a 12 18 41.1 +18 39 36 950 + 14NGC 4369 S0a Sc, S0, -, S0a 12 22 07.8 +39 39 33 1000 M 8, 14, 48NGC 4383 S0a I0?, Sa, -, p 12 22 53.9 +16 44 49 1700 + 12, 13, 44NGC 4385 SB0 SBbc, SB0, -, SB0 12 23 09.0 +00 50 53 2150 + 24, 44NGC 4438 S0a Sb, S0a, -, S 12 25 13.5 +13 17 11 250 + 2NGC 4457 S0a Sb, S0a, -, S0a 12 26 26.1 +03 50 51 850 + 13, 44


Table I (continued)

(1) (2)a (3) (4) (5) (6) (7) (8)Source Hubble Type �1950 �1950 vLSR CO Ref.b

NGC 4476 E E, S0, E, S0 12 27 26.8 +12 37 27 2000 + 23, 43NGC 4526 S0 S0, S0, -, S0 12 31 30.4 +07 58 33 600 + 12, 49NGC 4649 E S0, E, E, E 12 41 09.0 +11 49 23 1000 + 12NGC 4691 SB0 I0, SB0a, -, - 12 45 38.9 �03 03 37 1150 M 24, 32, 37, 44NGC 4710 S0 S0, S0, -, S0a 12 47 09.1 +15 26 15 1100 M 12, 13, 18, 33, 41, 44NGC 4940 S0a -, Sa, -, S0 13 02 06.2 �46 58 07 5000 M 48NGC 5128 E S0+S, S0 13 22 31.6 �42 45 32 550 M 4, 16, 17

-, E+Sb? 22, 26, 27, 35NGC 5195 I0 SB0, I0, -, I0 13 27 52.5 +47 31 48 600 M 11, 12, 18, 44NGC 5253 I0 I0, Im, -, E/S0 13 37 05.2 �31 23 21 400 + 8, 14NGC 5266 E S0, S0, -, S0 13 39 56.0 �47 54 49 3000 + 36NGC 5666 E -, S?, -, - 14 30 43.4 +10 43 47 2200 + 21, 23, 43NGC 5866 S0 S0, S0, -, S0 15 05 07.8 +55 57 16 700 + 13, 14, 39, 44, 48NGC 5953 S0 -, Sa, -, S0 15 32 13.2 +15 21 40 2000 + 8, 14, 44, 48NGC 6014 S0 -, S0, -, S0 15 53 29.4 +23 05 01 2350 + 45NGC 6524 E/S0 -, S0, -, E/S0 17 57 50.0 +45 53 23 5600 + 45IC 5063 E S0a, S0, -, S0a 20 48 12.7 �57 15 32 3350 + 24, 43NGC 7013 S0a -, S0a, -, Sa 21 01 26.1 +29 41 51 800 + 13, 14, 39, 44NGC 7077 dE/S0 -, S0?, -, E? 21 27 27.5 +02 11 39 1150 + 28NGC 7176e E -, E, E, E 21 59 13.8 �32 13 48 2750 + 14, 25NGC 7213 S0a Sa, Sa, -, S0 22 06 09.0 �47 25 42 1750 M 14, 32NGC 7233 S0a -, S0a, -, SBa 22 12 44.0 �46 05 48 1900 + 24NGC 7252 E merger, S0a, -, mult. 22 17 58.0 �24 55 48 4750 + 15, 30, 44NGC 7371 S0a SBa, S0a, -, - 22 43 25.2 �11 16 00 2850 + 1, 5NGC 7465 SB0 -, SB0a, -, SB0 22 59 31.9 +15 41 47 1950 + 13, 23, 44NGC 7632 SB0a -, SB0a, -, S0 23 19 17.0 �42 45 12 1550 + 14NGC 7679 S0 Sc/Sa, SB0, -, S0 23 26 13.6 +03 14 09 5150 + 14

a Hubble types not taken from Table 2 of Knapp et al. (1989): NGC 205: Freedman (1992); NGC 404:Wiklind and Henkel (1990); NGC 855: Walsh et al. (1990); UGC 2456: Nilson (1973); NGC 1316:Schweizer (1980); UGC 2836: Lees et al. (1991); NGC 1482: Lauberts (1982), de Vaucouleurs et al.(1991); NGC 1546: Lauberts (1982), de Vaucouleurs et al. (1991); NGC 1947: Bertola et al. (1992b);NGC 3077: Price and Gullixson (1989); NGC 3597: Taniguchi et al. (1990b); UGC 5720: Looseand Thuan (1986); NGC 3656: Balcells and Stanford (1990), Mollenhoff et al. (1992); NGC 3665:Gregorini et al. (1989); UGC 6877: Sage et al. (1992), de Vaucouleurs et al. (1991); NGC 4318:Nilson (1973), de Vaucouleurs et al. (1991); NGC 4369: Nilson (1973); NGC 4438: de Vaucouleurset al. (1991); NGC 4457: Nilson (1973), de Vaucouleurs et al. (1991); NGC 4476: Prugniel et al.(1987); NGC 4691: Wiklind et al. (1993); NGC5266: Varnas et al. (1987); NGC 5953: Nilson (1973),de Vaucouleurs et al. (1991); IC 5063: Colina et al. (1991); NGC 7013: de Vaucouleurs et al. (1991);NGC 7077: Nilson (1973), Sage et al. (1992); NGC 7233: de Vaucouleurs et al. (1991); NGC 7252:Dupraz et al. (1990); NGC 7371: de Vaucouleurs et al. (1991); NGC 7632: de Vaucouleurs et al.(1991); NGC 7679: Nilson (1973).


b References to molecular data: (1) Verter (1985); (2) Stark et al. (1986); (3) Wiklind and Rydbeck(1986); (4) Phillips et al. (1987); (5) Verter (1987); (6) Becker et al. (1989); (7) Heckman et al.(1989b); (8) Jackson et al. (1989); (9) Lazareff et al. (1989); (10) Mirabel et al. (1989); (11) Sage(1989); (12) Sage and Wrobel (1989); (13) Thronson et al. (1989); (14) Wiklind and Henkel (1989);(15) Dupraz et al. (1990); (16) Eckart et al. (1990b); (17) Israel et al. (1990); (18) Sage (1990);(19) Taniguchi et al. (1990a); (20) Wiklind and Henkel (1990); (21) Gordon (1991); (22) Israel et al.(1991); (23) Lees et al. (1991); (24) Tacconi et al. (1991); (25) Huchtmeier and Tammann (1992); (26)Israel (1992); (27) Quillen et al. (1992); (28) Sage et al. (1992); (29) Wang et al. (1992a); (30) Wanget al. (1992b); (31) Wiklind and Henkel (1992a); (32) Wiklind and Henkel (1992b); (33) Wrobel andKenney (1992); (34) Reuter et al. (1993); (35) Rydbeck et al. (1993); (36) Sage and Galletta (1993);(37) Wiklind et al. (1993); (38) McNamara and Jaffe (1994); (39) Taniguchi et al. (1994); (40) Watsonet al. (1994); (41) Welch and Mitchell (1994); (42) Horellou et al. (1995); (43) Wiklind et al. (1995);(44) Young et al. (1995); (45) Knapp and Rupen (1996); (46) Young and Lo (1996); (47) Welch et al.(1996); (48) Wiklind and Henkel (1997); (49) This review.c The HI linewidth of �75 km s�1 (van Driel et al., 1988) and the CO linewidth of 285 km s�1

(Tacconi et al., 1991) are discrepant. CO should be remeasured.d The spectrum shown in Wiklind and Henkel (1989) was obtained towards a position that is offsetfrom the center of the galaxy.e Accounting for the different treatment of the spectra, the Wiklind and Henkel (1989) and Huchtmeierand Tammann (1992) data are consistent.f Tentative or controversial detections, requiring follow-up observations:NGC 1400 (Lees et al., 1991; Bregman et al., 1992), NGC 2217 (Tacconi et al., 1991), NGC 2831(Wiklind et al., 1995), NGC 4261 (Jaffe and McNamara, 1994; this paper, Sect. 6.2.1), IC 3370(Lees et al., 1991), NGC 4429 (Sage and Wrobel, 1989), NGC 4459 (Sage and Wrobel, 1989;Thronson et al., 1989; Young et al., 1995), NGC 4546 (Sage and Galletta, 1994), NGC 4550 (Wiklindand Henkel 1997), NGC 4650A (Watson et al., 1994), NGC 4697 (Sofue and Wakamatsu, 1993;Knapp and Rupen, 1996), NGC 5084 (Bregman et al., 1992), NGC 6278 (Knapp and Rupen, 1996),NGC 6987 (Knapp and Rupen, 1996), 235–G49 (Knapp and Rupen, 1996), NGC 7457 (Taniguchiet al., 1994), NGC 7722 (Wang et al., 1992a).g Negative results:Nearby galaxies (cz < 7000 km s�1): NGC 216 (Lees et al., 1991), NGC 315 (Braine et al., 1997),NGC 632 (Jackson et al., 1989), NGC 708 (Grabelski and Ulmer, 1990; Sofue and Wakamatsu, 1993;Braine and Dupraz, 1994), NGC 720 (Braine et al., 1997), NGC 741 (Wiklind et al., 1995), NGC 770(Young et al., 1995), NGC 1023 (Verter, 1985; Bregman et al., 1992), NGC 1052 (Lees et al., 1991;Wang et al., 1992a; Wiklind et al., 1995; Knapp and Rupen, 1996), NGC 1060 (Knapp and Rupen,1996), NGC 1153 (Knapp and Rupen, 1996), NGC 1380 (Horellou et al., 1995), NGC 1387 (Horellouet al., 1995), IC 2006 (Lees et al., 1991), IC 2040 (Horellou et al., 1995), NGC 1600 (Sofue andWakamatsu, 1993), NGC 1638 (Bregman et al., 1992), NGC 2110 (Heckman et al., 1989b), IC 450(Heckman et al., 1989b), NGC 2314 (Verter, 1985, 1987), NGC 2502 (Knapp and Rupen, 1996),NGC 2534 (Wang et al., 1992a; Wiklind et al., 1995), NGC 2783 (Wiklind et al., 1995), NGC 2832(Knapp and Rupen, 1996), NGC 2974 (Bregman et al., 1992; Sofue and Wakamatsu, 1993; Knappand Rupen, 1996), NGC 2983 (Jackson et al., 1989), NGC 3011 (Jackson et al., 1989), NGC 3065(Bregman et al., 1992), NGC 3081 (Heckman et al., 1989b), NGC 3090 (Knapp and Rupen, 1996),NGC 3226 (Verter, 1985; Thronson et al., 1989; Young et al., 1995), NGC 3245 (Sage and Wrobel,1989), NGC 3273 (Lees et al., 1991), NGC 3311 (Grabelski and Ulmer, 1990; McNamara and Jaffe,1994), NGC 3379 (Sofue and Wakamatsu, 1993), NGC 3413 (Thronson et al., 1989; Young et al.,1995), NGC 3516 (Heckman et al., 1989b), NGC 3773 (Jackson et al., 1989), NGC 3805 (Knappand Rupen, 1996), NGC 3837 (Wiklind et al., 1995; Knapp and Rupen, 1996), NGC 3841 (Knappand Rupen, 1996), NGC 3842 (Wiklind et al., 1995; Knapp and Rupen, 1996), NGC 3870 (Verter,1985), NGC 4111 (Bregman et al., 1992), NGC 4203 (Bregman et al., 1992), NGC 4278 (Lees et al.,1991; Bregman et al., 1992), NGC 4335 (Knapp and Rupen, 1996), NGC 4365 (Braine et al., 1997),NGC 4374 (Sofue and Wakamatsu 1993; Wiklind et al., 1995; Knapp and Rupen, 1996), NGC 4406(Bregman and Hogg 1988; Wiklind et al., 1995; Braine et al., 1997), NGC 4435 (Sage and Wrobel1989; Bregman et al., 1992), NGC 4472 (Bregman and Hogg, 1988; Huchtmeier et al., 1988, 1994;Braine et al., 1997), NGC 4486/M 87/Vir A (Braine and Wiklind, 1993; Sofue and Wakamatsu,1993), NGC 4503 (Young et al., 1995), NGC 4515 (Knapp and Rupen, 1996), NGC 4564 (Bregmanet al., 1992), NGC 4581 (Lees et al., 1991; Wiklind et al., 1995; Knapp and Rupen, 1996), NGC 4589(Wang et al., 1992a; Sofue and Wakamatsu, 1993; Wiklind et al., 1995), NGC 4596 (Young et al.,


1995), NGC 4621 (Bregman et al., 1992), NGC 4636 (Braine et al., 1997), NGC 4645 (Knapp andRupen, 1996), NGC 4684 (Bregman et al., 1992), NGC 4696 (O’Dea et al., 1994b), 507–G25 (Knappand Rupen, 1996), NGC 4742 (Lees et al., 1991; Knapp and Rupen, 1996), NGC 4753 (Thronsonet al., 1989; Knapp and Rupen 1996), NGC 4936 (Knapp and Rupen, 1996), NGC 4984 (Younget al., 1995), NGC 5018 (Lees et al., 1991), NGC 5087 (Bregman et al., 1992), NGC 5101 (Tacconiet al., 1991), NGC 5102 (Tacconi et al., 1991; Bregman et al., 1992), NGC 5273 (Heckman et al.,1989b), NGC 5322 (Knapp and Rupen, 1996), NGC 5363 (Sanders and Mirabel, 1985; Sage andWrobel, 1989; Thronson et al., 1989; Lees et al., 1991; Young et al., 1995), 511–G23 (Knapp andRupen, 1996), UGC 9473 (Knapp and Rupen, 1996), NGC 5820 (Bregman et al., 1992), NGC 5838(Knapp and Rupen, 1996), NGC 6654 (Lees et al., 1991; Knapp and Rupen, 1996), NGC 6702(Wang et al., 1992a; Wiklind et al., 1995; Knapp and Rupen, 1996), NGC 6703 (Knapp and Rupen,1996), NGC 6861 (Knapp and Rupen, 1996), NGC 7052 (Wang et al., 1992a; Wiklind et al., 1995),NGC 7075 (Knapp and Rupen, 1996), NGC 7332 (Verter, 1985), NGC 7426 (Knapp and Rupen,1996), IC 1459 (Lees et al., 1991), NGC 7464 (Lees et al., 1991; Wiklind et al., 1995), NGC 7468(Verter, 1985; Lees et al., 1991; Wiklind et al., 1995), NGC 7617 (Knapp and Rupen, 1996), Mrk 328(Sage et al., 1992).Distant galaxies (cz > 7000 km s�1): For CO observations, mostly from the central regions of Abellclusters, see Bregman and Hogg (1988), Mirabel et al., (1989), Grabelski and Ulmer (1990), Antonucciand Barvainis (1994), Barvainis and Antonucci (1994), Braine and Dupraz (1994), McNamara andJaffe (1994), and O’Dea et al. (1994b).

4. Early-Type Galaxies with Cool Dense Gas: General Results

4.1. SOME USEFUL QUANTITIES

Having presented the molecular data in Section 3, we still have to define thosephysical parameters that are needed for a careful interpretation. Usually, two rota-tional CO transitions are measured, the J = 1 � 0 line at � = 2:6 mm and theJ = 2 � 1 line at 1.3 mm. Their relative intensity provides a rough measure ofphysical conditions. J = 2 � 1=J = 1 � 0 line intensity ratios well in excess ofunity imply warm (>20 K) gas of small column density. If the ratio is smaller thanone, the gas is likely cool (<20 K) and optically thick in the CO lines. Sometimes12CO (the main species) and 13CO (a rare isotopic species) are both observed.Assuming that the transitions of 13CO are optically thin and that the 12CO/13COabundance ratio is of order 50, we can obtain an approximate estimate of lineopacities from the measured 12CO/13CO line intensity ratios. This also allows thedetermination of a ‘beam filling factor’, specifying that fraction of the radio beamfrom where the bulk of the emission arises (see the Appendix; Section A.1).

One of the most important parameters is the molecular gas mass. Although12CO lines tend to be saturated, there may exist an approximately linear correl-ation between integrated CO line intensity, ICO, and the column density of themost abundant molecule, NH2 (e.g., Henkel et al., 1991; Solomon and Barrett,1991; Young and Scoville, 1991; Mauersberger and Henkel, 1993). The correlationobtained from the galactic disk (Strong et al., 1988), accounting for a re-evaluationof the calibration scheme by Bronfman et al. (1988), is

NH2=ICO = 1:9� 1020 cm�2 (K km s�1)�1; (1)


leading to

MH2

M�

= 3:0 1049�mH2

M�

� �DMpc tg

��b

2

��2 �NH2

ICO

�ICO (2)

(�b: beam diameter; DMpc: distance in Mpc). The NH2=ICO conversion factorshould generally be applicable in regions with virialized molecular clouds of solarmetallicity emitting optically thick subthermally excited CO lines and correspondsto an H2 surface density of �3 M� pc�2 (K km s�1)�1. While small galaxies oflow metallicity show a tendency for larger conversion factors (e.g., Ohta et al.,1993; Rubio et al., 1993; Becker et al., 1995; Wilson, 1995; Arimoto et al., 1996;Sakamoto, 1996), recent studies of 12CO, C18O, and dust emission indicate lowervalues (by factors of�3–5) towards the central regions of the Milky Way, the nearbystarburst spirals NGC 253 and NGC 4945, and distant ultraluminous infraredgalaxies (cf. Sodrowski et al., 1995; Dahmen et al., 1996; Mauersberger et al.,1996a, b; Solomon et al., 1997). Furthermore, galaxies with small S60�m/S100�m

dust color temperatures seem to have gas to dust mass ratios well below�700, thecommon value only accounting for the warm dust seen by IRAS (Wiklind et al.,1995). An interpretation in terms of an underestimated molecular mass is one ofseveral possibilities. All this implies that molecular masses and surface densitiesgiven in Table II are accurate to only �0.7 on a logarithmic scale.

Other useful quantities include the star formation efficiency (SFE) and its recip-rocal value, the gas consumption time. The SFE is defined by the ratio between thestar formation rate (SFR) and the molecular gas mass, thus providing an estimateof the efficiency by which the molecular ‘fuel’ is converted into stars. In order tohave a physical meaning, the SFE has to be defined over a specified time interval,that is often implicitly assumed to be the interval over which the SFR is measured.The critical time scale is of order 107 yr, if the SFR is deduced from the far infraredor H� luminosity. The blue luminosity, LB , is usually taken as a measure of thelonger term (� 109 yr) star formation rate (Gallagher et al., 1984; Young et al.,1985; but see Sage and Solomon, 1989). The Appendix (Section A.2) providescommonly used equations relating measured luminosities to star formation rates.

LFIR provides upper limits to the SFR since it is not certain that all the infraredemission is due to dust in star forming regions (Helou, 1986). With LH� and LB ,lower limits can be obtained, since LH� and even more so LB may suffer fromextinction. The gas consumption time,MH2/SFR, is then the time needed to convertall the molecular gas into stars, assuming that no molecular material is dispersed,that no gas is added, and that star formation rate and initial mass function remainconstant. While these assumptions may not be realistic (cf. Section 6.5), SFE andgas consumption time are useful measures of the current star formation activity. Agas consumption time of �gas = 5� 108 yr as determined for the galactic disk (cf.Henkel and Mauersberger, 1993) and a stellar mass loss rate of _M = 0:015M� yr�1

(LB /109 L�) for an old stellar population (Faber and Gallagher, 1976) then yields


Table IICharacteristic properties of early-type galaxies detected in CO (Tdust is given in units of K; MH2 ,MHI, LFIR, LB , and LX values in units of M� or L� are displayed on a logarithmic scale; seeSection 3 for further information)

(1) (2) (3) (4) (5) (6) (7) (8) (9) (10) (11)Source DMpc MH2 �H2 �b MHI Tdust LFIR LB LX Ref.a

NGC 83 85 9.0 14 21 37 10.4 10.6 45NGC 185 0.7 4.9 7 33 5.2 26 6.1 8.2 13, 45NGC 205 0.7 5.1 5 55 5.6 31 6.5 8.4 3.3 14, 22, 32, 45NGC 404b 10 7.8 35 21 8.8 38 8.9 9.8 24, 29NGC 759c 65 9.1 35 21 <9.4 33 10.2 10.5 45UGC 1503 70 8.9 18 21 9.2 30 10.0 10.2 45NGC 807 65 9.0 22 21 9.8 28 10.2 10.5 45NGC 855 9 6.3 2 21 8.0 34 8.7 9.0 45IC 1830 17 7.7 5 43 9.3 39 9.5 9.5 33UGC 2456 50 9.3 35 33 <8.9 42 9.5 10.4 9, 24NGC 1291 9 7.3 7 43 8.6 30 8.9 10.4 33, 43NGC 1275 70 9.6 30 21 10.0 46 11.0 11.2 20, 26NGC 1316d 20 8.5 25 43 <8.1 34 9.7 11.1 7.2 10, 32, 39, 45NGC 1326 20 8.3 20 43 9.6 39 10.0 10.0 33NGC 1386 20 8.3 14 44 <7.6 39 9.9 9.7 6.9 34, 44UGC 2836 70 9.7 50 33 35 11.0 11.0 24NGC 1482 25 8.9 33 45 41 10.8 9.8 46NGC 1546 13 8.5 60 43 30 9.8 9.6 50NGC 1819 65 9.3 12 45 9.5 38 11.0 10.7 23NGC 1947 14 8.1 7 43 28 9.2 9.9 6.1 32, 39, 45NGC 2328 15 7.9 11 43 41 9.3 9.4 45NGC 2320 80 8.6 7 21 26 10.3 10.8 45NGC 2685 18 7.9 45 17 9.3 27 8.9 10.0 5.7 6, 10, 28, 32NGC 2768 25 7.1 3 21 <8.2 30 9.2 10.7 45NGC 3032 25 8.0 4 45 8.4 35 9.7 9.8 23NGC 3077 3 7.0 120 21 9.0 38 8.7 9.0 <4.9 18, 34, 41NGC 3166e 17 8.5 120 21 8.3 34 9.8 10.2 6.4 7, 24, 32NGC 3265 25 7.7 4 32 8.3 41 9.7 9.5 30, 31, 45UGC 5720 20 7.8 19 21 8.7 45 9.7 9.4 37NGC 3593 7 8.7 190 21 8.3 36 9.5 9.6 5.5 10, 23, 32, 38NGC 3597 50 9.3 21 43 <10.4 42 10.8 10.2 45NGC 3656 45 9.1 70 21 9.1 34 10.3 10.2 45NGC 3665 35 8.7 8 55 29 10.1 10.5 22NGC 3682 25 8.4 21 33 35 10.0 9.7 24UGC 6877 25 7.2 3 21 8.3 8.4 9.3 37NGC 4125 25 7.7 8 21 <9.0 35 9.2 10.7 45NGC 4138 20 8.0 5 45 9.2 9.9 3, 23, 47NGC 4293 18 8.8 210 21 9.0 38 9.6 10.3 1, 24NGC 4369 18 8.3 70 21 8.6 36 9.8 9.8 19, 24NGC 4383 18 7.9 7 45 9.3 41 9.9 9.7 23NGC 4385 18 7.9 7 43 8.4 42 9.7 9.6 <6.9 33, 34NGC 4438 18 8.8 10 100 8.9 32 9.7 10.5 6.9 4, 12, 34


Table II (continued)

(1) (2) (3) (4) (5) (6) (7) (8) (9) (10) (11)Source DMpc MH2 �H2 �b MHI Tdust LFIR LB LX Ref.a

NGC 4457 18 8.2 14 45 8.4 35 9.8 10.0 23NGC 4476 18 7.7 16 21 <8.3 33 9.0 9.5 <6.6 32, 45NGC 4526 18 8.3 11 55 9.3 32 9.9 10.5 6.2 1, 22, 32NGC 4649 18 7.6 2 55 <8.0 44 9.1 10.8 7.5 10, 22, 32, 45NGC 4691 18 8.9 25 43 9.0 37 10.1 10.0 33, 40NGC 4710 18 8.2 13 45 7.4 34 9.9 10.0 23NGC 4940 70 9.5 17 43 32 10.7 10.7 11, 50NGC 5128 4 8.3 80 43 8.6 36 10.0 10.5 6.3 10, 32, 36, 45NGC 5195 10 8.6 30 55 <7.8 9.5 10.0 5.8 2, 22, 27, 32NGC 5253 3 6.1 4 43 8.1 50 8.9 9.0 4.8 1, 24, 34NGC 5266 40 8.7 8 43 9.9 31 10.0 10.8 39, 45, 49NGC 5666 35 8.6 40 21 9.0 37 9.9 9.8 45NGC 5866 13 7.8 40 21 <8.5 30 9.7 10.1 6.0 10, 23, 24, 32NGC 5953 35 9.1 45 33 8.9 36 10.5 10.0 16, 24NGC 6014 35 8.4 6 30 �8.4 41 9.8 10.0 15, 48NGC 6524 75 8.6 2 30 35 10.9 48IC 5063 40 8.3 3 43 9.6 67 10.6 10.4 7.0 32, 45NGC 7013 13 7.5 23 21 8.9 34 9.1 10.1 23, 24NGC 7077 15 6.6 2 21 8.2 9.0 37NGC 7176f 35 9.2 4 43 33 10.2 10.4 35, 45NGC 7213 20 8.8 45 43 9.6 30 9.8 10.4 9.0 8, 24, 32NGC 7233 20 8.1 8 43 9.3 38 9.8 9.6 11, 33NGC 7252 60 8.7 17 21 9.5 39 10.7 10.7 7.9 25, 34, 42NGC 7371 35 8.7 2 100 9.7 10.4 17NGC 7465 25 8.1 7 45 9.5 36 10.0 9.8 23NGC 7632 20 8.5 30 43 38 9.6 9.6 24NGC 7679 65 9.3 25 33 9.2 42 11.0 10.6 8.1 5, 11, 34

a (1) Davies and Lewis (1973); (2) Weliachew and Gottesman (1973); (3) Grewing and Mebold(1975); (4) Huchtmeier and Bohnenstengel (1975); (5) Balick et al. (1976); (6) Shane (1980); (7)Haynes (1981); (8) Reif et al. (1982); (9) Mirabel and Wilson (1984); (10) Forman et al. (1985);(11) IRAS (1989); (12) Stark et al. (1986); (13) Wiklind and Rydbeck (1986); (14) Canizareset al. (1987); (15) Chamaraux et al. (1987); (16) Garwood et al. (1987); (17) Verter (1987); (18)Becker et al. (1989); (19) Jackson et al. (1989); (20) Lazareff et al. (1989); (21) Sage (1989);(22) Sage and Wrobel (1989); (23) Thronson et al. (1989); (24) Wiklind and Henkel (1989);(25) Dupraz et al. (1990); (26) Jaffe (1990); (27) Sage (1990); (28) Taniguchi et al. (1990a);(29) Wiklind and Henkel (1990); (30) Gordon (1991), ICO(2�1)=ICO(1�0) = 0:5; (31) Lees et al.(1991), ICO(2�1)=ICO(1�0) = 0:5; (32) Roberts et al. (1991); (33) Tacconi et al. (1991); (34)Fabbiano et al. (1992); (35) Huchtmeier and Tammann (1992); (36) Israel (1992); (37) Sage et al.(1992); (38) Wiklind and Henkel (1992a); (39) Sage and Galletta (1993); (40) Wiklind et al.(1993); (41) Bi et al. (1994); (42) Hibbard et al. (1994); (43) Bregman et al. (1995); (44) Horellouet al. (1995); (45) Wiklind et al. (1995); (46) Young et al. (1995); (47) Jore et al. (1996); (48)Knapp and Rupen (1996), ICO(2�1)=ICO(1�0) = 0:5; (49) Morganti et al. (1997); (50) Wiklindand Henkel (1997).b �H2 refers to the (000,1000) offset position (see Wiklind and Henkel, 1990).c �H2 � 750M� pc�2 at r �650 (Wiklind et al., 1997).d �H2 refers to the (�2100; 2100) offset position (see Sage and Galletta, 1993).e MHI is uncertain (Haynes, 1981).f LFIR was obtained from the combined emission of NGC 7173 and NGC 7176.


Figure 3. HI column density distribution for nine early-type disk galaxies (from van Driel and vanWoerden 1991). Each galaxy has been rotated in such a way that the HI kinematical axis is in verticaldirection. Dashed ellipses: 25m arc sec�2 optical contour. Linear scale: 10 kpc. Hatched ellipse: HIbeam.

an equilibrium cool gas mass of �gas _M <� 108 M� for an average giant ellipticalgalaxy.

4.2. A CAUTIONARY NOTE

As already mentioned (Section 3), the observed sample of sources is not complete.The large fraction of reported nondetections (see Table I), the small number ofdetected sources (� 70), and uncertainties in the morphological classification mayblur any emergent picture. Nevertheless the results presented here provide a usefulstarting point for a future complete and sensitive volume limited survey.


4.3. LENTICULAR GALAXIES

In spiral galaxies, far infrared fluxes are positively correlated with integrated COintensities (e.g., Sanders, 1991). Thus, IRAS selected samples of early-type galaxieswere used to maximize detection rates. The three most comprehensive CO linesurveys, those of Sage and Wrobel (1989), Thronson et al. (1989), and Wiklind andHenkel (1989) highlight different properties; there is however significant overlapand, in these cases, the agreement is good. The main results, also accounting forthe studies of Lees et al. (1991), Bregman et al. (1992), and Hogg et al. (1993) areas follows:

(1) Average H2 gas masses and star formation rates are smaller by an order ofmagnitude compared to spirals of the same blue luminosity.

(2) HI and CO emission mostly arises from rotating disks. In one case, themolecular disk is known to be tilted relative to the stellar disk (Wiklind 1991).While HI is often distributed in inner and outer rings (Figure 3), the latter beinglocated outside the main body of the galaxy (e.g., van Driel and van Woerden,1991), the molecular gas appears to be more centrally concentrated than in spiralgalaxies. To give an example: A comparison of 12-m NRAO and 30-m IRAMCO data from NGC 4526 (Figure 2) shows almost identical lineshapes and theline temperature ratio is approximately 1:6 as expected for a source of size <2000

(<1.6 kpc). Characteristically, much of the molecular emission is contained ina single telescope beam, so that comparisons between total molecular and farinfrared emission can be made without applying corrections. Only NGC 5195, thecompanion of the ‘whirlpool’ galaxy M 51, is known to show emission over severalarcmin (Sage, 1989). CO beam filling factors of a few percent are common for thecentral kiloparsec (e.g., Wiklind and Henkel, 1992a).

(3) The J = 2 � 1=J = 1 � 0 CO and the 12CO/13CO line intensity ratios,0.5�1.0 and 5�15, respectively, are consistent with those observed in molecularclouds of the galactic disk and other spiral galaxies. The data indicate opticallythick 12CO emission from predominantly cool (<20 K) gas.

(4) MH2 /Mdust mass ratios, S100�m/S60�m flux ratios, and star formation effi-ciencies are similar in lenticular and spiral galaxies.

(5) Assuming that gas and dust temperatures differ by a constant factor or areequal and normalizing the CO and far infrared luminosities according to theirdependence on temperature (Tdust and T 5

dust, respectively), Wiklind and Henkel(1989) obtain an approximately linear correlation,

MH2

Tdust= (1:8� 0:5)� 105

LIR

T 5dust

!0:90�0:05

; (3)

(cf. Figures 4 and 5). The correlation between H2 mass and blue luminosity,

MH2 / L1:0�0:1B ; (4)


Figure 4. The molecular mass (from Equations (1) and (2)) as a function of infrared luminosityfor the galaxies presented in Tables I and II. Open circles: dwarf elliptical galaxies; filled circles:elliptical galaxies; asterisks: E/S0 galaxies; open squares: lenticular galaxies; crosses: S0a galaxies;horizontal lines: I0 galaxies. For comparison, some spirals and ultraluminous IRAS galaxies (fromYoung et al., 1989) are also displayed (small triangles).

also appears to be linear (see Figure 6; MH2 and Tdust are given in M� and K,the luminosities in L�). The correlations (3) and (4) hold for lenticular and spiralgalaxies and are valid over 3�4 orders of magnitude in MH2, LB , (MH2=Tdust) and(LFIR/T 5

dust). The MH2 –LB correlation is not quite as well defined as in spirals. Atthe low end, the MH2/LB ratios are smaller than for any spiral, while the ratioscan be as high as for the most gas-rich spirals (MH2/LB�0.1 in solar units; cf.Lees et al., 1991). S0a galaxies are similar to spirals in the sense that they have awell defined most likely mean value of the MHI/LB ratio (�0.05M�/L�). In S0galaxies, the MHI/LB ratios are consistent with a power-law (Wardle and Knapp,1986).


Figure 5. Same as Figure 4, but with H2 masses and far infrared luminosities (in solar units) dividedby Tdust and T 5

dust (in K), respectively.

(6) From Hubble types S0a to E/S0, the fraction of objects detected in HI or inthe far infrared gradually decreases. Mass and surface density of the neutral gascomponent also drops, both in absolute terms and after normalization to the blueluminosity (Hogg et al., 1993).

(7) The normalized HI mass MHI/LB is reduced in the cores of large clusters(e.g., Haynes et al., 1984). While the presence of bars does not strongly affectMHI/LB or LFIR/LB ratios, MHI and LFIR are enhanced in galaxies with nuclearactivity (Eskridge and Pogge, 1991).

(8) Individual MH2/MHI ratios can vary by at least two orders of magnitude.Average values for those galaxies where CO is detected were reported to be signi-ficantly larger than one (Young and Knezek, 1989), slightly larger than one or oforder one (Thronson et al., 1989), and consistent with one (Lees et al., 1991).


(9) Many polar ring galaxies are lenticulars (see Sections 5.3 and 6.5). Theyare characterized by large HI masses relative to their blue and infrared luminosities(MHI/LB�MHI/LFIR� 1M�/L�), suggesting that most of the gas in the rings isin stable orbits, rather than flowing inwards triggering star formation in the innergalaxy (Richter et al., 1994; Cox and Sparke, 1996; Huchtmeier, 1997). Moreinterferometric data (see e.g., van Gorkom et al., 1987) are desirable to confirmthis result.

(10) A detailed comparison ofLFIR,LH�,LB , ICO, and the estimated mass returnrate from evolved stars indicates that at least some of the galaxies undergo ‘bursts’of star formation on time scales �108 yr (see Section 5.1 for specific examples).During the ‘bursts’, star formation activity reaches levels which are typical for thedisk of the Milky Way. Thus lenticulars, like spirals, can show variations in starformation rates over short periods of time. S0 ‘bursts’ are however less impressiveand the level of activity during quiescent periods is lower.

(11) The average LFIR/LB ratio, measuring the current versus the long-termaverage star formation rate, is similar for lenticular and spiral galaxies. This indic-ates that the star formation rate averaged over time intervals � 108 yr remainsapproximately constant.

(12) The presence of large molecular clouds with normal star formation efficien-cies in systems without spiral arms demonstrates that the H2 mass, not the spiralstructure, is determining the global star formation rate; spiral arms are largelyirrelevant to the SFE and the gas consumption time. For the formation of largemolecular clouds, spiral arms are also not required.

A consequence of an almost constant star formation rate per unit molecular mass(see Figure 4) is that the rate of star formation varies linearly with surface density(at least above a certain threshold value; cf. Kennicutt, 1989). The similarity inthe J = 2 � 1/J = 1 � 0 CO and 12CO/13CO line intensity ratios, MH2=Mdust

mass ratios, and MH2 -LB and MH2-Tdust correlations indicates that molecular gasproperties do not greatly differ between lenticulars and spirals. It is thus possible thatthe galactic disk N (H2)/ICO conversion ratio (see Section 4.1) is indeed applicableto lenticular galaxies as well. The similarS100�m/S60�m,MH2=Mdust, andMH2=Tdust

ratios of lenticular and spiral galaxies suggest that the bulk of the FIR emissionarises from dust heated by stars.

While many of the summarized results are quite general and of interest for alltypes of disk galaxies, we still have to account for systematic effects caused bythe selection of sources (cf. Sections 3 and 4.2). Selecting IRAS galaxies implies abias in favor of a moderately massive ISM and a higher than average molecular gascontent. Hence, on the average, the molecular gas mass and star formation rate inlenticulars may be more than an order of magnitude smaller than in spiral galaxies.Because of this bias, results suggesting MH2/MHI� 1 are probably not significantand MH2/MHI�1 as found for spirals (e.g., Sage, 1993) is likely the better value. Abias in favor of a more massive molecular ISM also explains another phenomenon:In most galaxies with CO detections, the star formation rate substantially surpasses


Figure 6. The molecular mass (from Equations (1) and (2) as a function of blue luminosity for thegalaxies presented in Tables I and II. Open circles: dwarf elliptical galaxies; filled circles: ellipticalgalaxies; asterisks: E/S0 galaxies; open squares: lenticular galaxies; crosses: S0a galaxies; horizontallines: I0 galaxies; small triangles: spirals and ultraluminous IRAS galaxies taken from Young et al.(1989).

the stellar mass return rate. This cannot be typical for all lenticular galaxies as canbe inferred from the large number of molecular non-detections (Lees et al., 1991;see also the discussion of NGC 404 in Section 5.1).

4.4. CD’S, ELLIPTICAL, AND DWARF ELLIPTICAL SYSTEMS

The family of nominally diskless systems is quite heterogeneous. A few nearbytransition type dE/dI dSph systems are known to contain HI (e.g., Da Costa, 1992;Skillman, 1996). Molecular data from compact ellipticals (e.g., M 32) and dwarfspheroidals (e.g., the Fornax system) have not yet been published. HI and IRASdata do not suggest the presence of a significant cool dense gas component in


Figure 7. CO J = 1 � 0 (upper panel) and 2–1 (lower panel) 30-m IRAM spectra of NGC 759,obtained with angular resolutions of 2100 and 1300 (Wiklind et al., 1995). Although the source remainsspatially unresolved, the double peaked lines indicate the presence of a rotating ring.

compact ellipticals which are also affected by massive nearby companions (cf.Burkert, 1994; for M 32, an upper CO 1–0 3� limit of Tmb = 0.015 K (channelwidth: 6 km s�1) was recently obtained with a 5000 beam; L. J. Sage, priv. comm.).Thus the following summary is confined to cD, giant, intermediate, and dwarfelliptical galaxies.

Most of the molecular data have been obtained by Lees et al. (1991) and Knappand Rupen (1996) in the CO J = 2 � 1 transition and by Wiklind et al. (1995)in the CO J = 1 � 0 and 2 � 1 transitions. Blue compact dwarfs with ellipticalisophotes have also been observed (Sage et al., 1992). As for the lenticulars, mostof the observed sources were selected from IRAS data; while this implies a biasin favor of a moderately massive ISM, searches for CO in FIR weak ellipticalswere unsuccessful (Braine et al., 1997). The following main results have to bementioned:

(1) 70% of the lenticulars withmB < 14m are detected by the IRAS satellite; thecorresponding percentage for bright ellipticals is only 45% (Knapp et al., 1989).While the detection rates are larger than for optical dust lanes (e.g., Veron-Cettyand Veron, 1988), the relative numbers (70 versus 45%) are consistent with the


Figure 8. Spatial distribution of the integrated CO emission from NGC 759, observed with the Plateaude Bure Interferometer. The map includes all channels between �237 and +185 km s�1 relative tocz = 4665 km s�1, with a contour spacing of 0.8 Jy km s�1. The hatched ellipse in the lower leftcorner shows the beam size; the cross denotes the optical (RC3) position of the galaxy. Epoch 2000coordinates are displayed, while Table I gives epoch 1950 coordinates (from Wiklind et al., 1997).

optical data and with the observational experience that it is more difficult to detectCO in elliptical than in lenticular galaxies.

(2) While it is too early to define a ‘typical’ molecular morphology, disk orring-like structures may be common. Although signal-to-noise ratios are small,filled-aperture measurements indicate that there are at least four studied galaxieswithout detectable nuclear CO emission (Dupraz et al., 1990; Wang et al., 1992b;Sage and Galletta, 1993; Wiklind et al., 1995, 1997). This might imply that, inellipticals, the molecular gas is less centrally condensed than in lenticulars. At leastin some cases, however, CO spectra obtained towards the nuclear regions includethe major part of the emission (e.g., Wiklind et al., 1995, 1997). Towards NGC 759


Figure 9. The distance distribution of 65 early-type galaxies with detected and 108 early-typegalaxies with undetected CO emission (for details, see Section 3 and Table I). Limiting redshift:cz = 6750 km s�1. The ordinate denotes the number of galaxies per bin. A significant differencebetween detected and undetected sources is not apparent in the case of the S0 and S0a galaxies. Forthe elliptical galaxies, the detection rate is particularly small in the 10–20 Mpc bin that is dominatedby Virgo cluster galaxies. Nearby ellipticals (D < 10 Mpc) show detection rates that are larger thanaverage.

(see Figures 7 and 8), filled-aperture and interferometric CO maps provide evidencefor a ring-like molecular structure with an average galactocentric radius of�650 pc(2 :002) and a surface density of several 100M� pc�2 (Wiklind et al., 1995, 1997;Henkel et al., 1996). In the case of Cen A, both an extended dust lane and a nuclearsource contribute to the CO emission (see Section 5.4).

HI emission often arises from rings that are tilted relative to the stellar body.Warped elongated features are also seen but the sample of investigated sources issmall (e.g., Raimond et al., 1981; van Gorkom et al., 1986). Some radio galaxiesshow HI absorption components that are redshifted w.r.t. the galaxy’s systemicvelocity, indicating non-circular motion (see Section 6.4.2).


(3) Most of the few measured J = 2 � 1/J = 1 � 0 line intensity ratios aresmaller than unity, thus providing some evidence for the presence of subthermallyexcited cool gas with optically thick CO lines (Wiklind et al., 1995, 1997).

(4) The correlation between the far infrared luminosity (LFIR) and the molecularmass (MH2) deduced from CO via the normal N (H2)/ICO conversion factor (seeSection 4.1) is in most cases consistent with that obtained for other types of galaxies(Figure 4). The far infrared colors determined from 60 and 100�m fluxes are similarto those found in morphologically ‘later’ galaxies. This implies that most of theinfrared emission arises from cool interstellar dust grains (cf. Section 6.1.1) and thatstar formation efficiencies, gas consumption times, and possibly also N (H2)/ICO

conversion factors are similar to those of disk galaxies.(5) In elliptical galaxies, MH2 and LB are less well correlated than in disk

galaxies (Figure 6). The progression from Sa to E galaxies introduces a gradualchange from a gaussian distribution of MH2 /LB to a power-law distribution (Leeset al., 1991). Such a distribution has also been determined forMHI/LB (Knapp et al.,1985a; Forbes, 1991). Van Gorkom (1992) proposesN (MHI/LB)/ (MHI/LB)�3=2

and N (MH2 /LB) / (MH2 /LB)�1=3, with N denoting the number of sources. Theaverage molecular gas content relative to LB progressively decreases from Sa/S0to elliptical galaxies. As for isolated spirals, no correlation between LFIR/MH2 andLB is obtained.

(6) The average MH2/MHI ratio is either slightly above (Lees et al., 1991)or below (Huchtmeier et al., 1995; Wiklind et al., 1995) unity. The latter result,although still referring to a small minority of nearby ellipticals, is based on a largersample and should be preferred. HI and H2 masses are poorly correlated. HI appearsto be depleted in galaxies residing in large clusters (e.g., Kumar and Thonnard,1983). If a significant portion of the H2 gas is concentrated at small galactocentricradii, the hot winds proposed for ellipticals like NGC 3265 and NGC 5666 (Lakeet al., 1987) must be weaker than anticipated to allow the molecules to survive (cf.Fabbiano, 1989; Smith et al., 1993).

(7) At least a few of the Blue Compact Dwarfs with elliptical isophotes in theirouter regions are dE systems, undergoing a violent phase of star formation. Theyappear more readily detectable (in CO) than their irregular counterparts (Sage et al.,1992).

(8) Field ellipticals seem to be overrepresented among those E galaxies that areknown to show CO emission (Wiklind et al., 1995). This is illustrated in Figure 9where the number of detected and undetected sources is displayed as a function ofdistance. The Virgo cluster bin shows a particularly small detection rate.

(9) Lees et al. (1991) suggested that molecular gas is preferentially detectedin the bluer, lower luminosity (dwarf) ellipticals, because these have more thantwice the detection rates of the brighter galaxies (their Figures 12 and 13). Now,having access to a larger set of data, Figure 10 displays the number of detected andundetected sources up to cz = 7000 km s�1 as a function of LB . No significantdifference in the distributions of detected and undetected sources is seen, although


Figure 10. The LB distribution of 63 early-type galaxies with detected and 107 early-type galaxieswith undetected CO emission (from Table I). The limiting redshift is cz = 7000 km s�1. The abscissais log (LB /L�), the ordinate is the number of galaxies in each bin. A significant difference betweenthe samples of detected and undetected sources is not apparent, although high luminosity lenticularsand low luminosity ellipticals seem to be slightly overabundant in the distribution of CO detectedgalaxies.

the undetected samples show a certain deficiency of high luminosity S0 and lowluminosity E galaxies. The latter result is consistent with a high detection rate ofnearby ellipticals (see Figure 9) that is (in part) caused by the fact that the twomost luminous Local Group dwarf ellipticals, NGC 185 and NGC 205, are knownto contain molecular gas (Section 5.2). Whether this is a general property of dE’sat the upper end of their luminosity range remains to be seen.

(10) Is the CO detection rate correlated with the intrinsic shape of the observedelliptical galaxies? CO detections alone are not numerous enough to answer thisquestion; since LCO and LFIR are correlated (Figures 4 and 5) and since the sampleof IRAS observed galaxies is relatively large and complete, we have taken theIRAS data to check whether the presence of a detectable cool ISM is relatedto the isophotal shape parameter a4 (Section 2) that describes deviations from


Figure 11. Far infrared luminosities as a function of the isophotal shape parameter a4 for the‘nearly complete’ sample of RSA ellipticals with BT < 12 :m4 and � < �10� (Bender et al., 1988).a4 < 0 describes a boxy, a4 > 0 denotes a disky elliptical. Using a (the semi-major axis length)for normalization, a4=a changes little with galactocentric radius. Sources with upper FIR luminositylimits are displayed as arrows.

a ‘perfect’ ellliptical. ‘Boxy’ E’s are characterized by a4 < 0, ‘disky’ E’s showa4 > 0. Combining the IRAS data presented by Knapp et al. (1989) with the almostcomplete sample of RSA ellipticals with BT < 12 :m4 and � > �10� (Bender et al.,1988), a correlation is not apparent (Figure 11). Taking instead the sample ofsources analysed by Bender et al. (1989) leads to a similar result.

4.5. THE VERY DENSE GAS COMPONENT

From detailed studies of the Galaxy it is well established that stars are formingout of compact molecular clumps; these clumps are denser than the surroundinginterclump gas that emits the bulk of the CO emission. Unlike CO that tracesn (H2)densities down to a few 100 cm�3, most other molecules requiren (H2)> 104 cm�3

to be detected in emission against the microwave background. Among these high


Figure 12. HCN spectra from (a) NGC 3593 and (b) NGC 4691. The 2500 resolution data wereobtained with the 30-m IRAM telescope in 1993. Note that the NGC 4691 detection is tentative.For further HCN data and corresponding CO spectra, see Israel (1992), Wiklind and Henkel (1989,1992a), and Wiklind et al. (1993).

density tracers, the � = 3 mm J = 1� 0 lines of HCN and HCO+ are prominent. Acomparison of the CO, HCN, and HCO+ line intensities thus provides informationon the density structure of the cool gas component. To date, three early-type galaxiesmay have been observed in HCN (one detection is tentative) and one is also observedin HCO+(see Figure 12 for data not published elsewhere). The line intensityratios are displayed in Table III. The star formation activity (see Section 4.1)and molecular line ratios from the central kpc of the lenticulars NGC 3593 andNGC 4691 are similar to the corresponding values from the galactic disk. Theemission from the nuclear region of Cen A has, however, all the molecular lineproperties of a nuclear starburst (see also Section 5.4 and Henkel et al., 1991).

Since the detection of an intense water-vapor ‘megamaser’ in the starburstgalaxy NGC 4945 (Dos Santos and Lepine, 1979), more than a dozen of similarsources were found in systematic searches for � = 1:3 cm JKaKc = 616 � 523 H2Oemission (Braatz et al., 1996b; Greenhill et al., 1997; Hagiwara et al., 1997). Theemission probably traces gas with 100 K < Tkin< 1000 K and n (H2)>107 cm�3


Table IIIEarly-type galaxies known to contain cool (Tkin < 1500 K) very high density (n (H2)>104 cm�3) molecular gas

Source Type CO/HCNa CO/HCO+a LH2O/L�b

NGC 1052, probably centerc E � � 200NGC 1386, probably centerd S0a � � 120NGC 3593, center S0 70e � �

NGC 4691, center S0 80e � �

Cen A, center E 24f 22f �

TXFS 2226–184, probably centerg E/S0 � � 6100Milky Way, disk Sb 100h � �

a Line intensity ratios from the J = 1 � 0 emission lines. Uncertainties are of order�20%.b Assuming isotropic emission.c H2O megamaser emission detected by Braatz et al. (1994).d H2O megamaser emission detected by Braatz et al. (1996b).e This review; the HCN detection in NGC 4691 is tentative so that the given ratio is aminimum value.f Israel (1992)g H2O gigamaser emission detected by Koekemoer et al., (1995).h Solomon et al. (1992a).

(e.g., Neufeld et al., 1994) located in the inner few pc of the parent galaxy (Claussenet al., 1986; Haschick et al., 1990; Miyoshi et al., 1995). Three of the known H2Omegamasers are associated with early-type galaxies, the E4 NGC 1052, the len-ticular NGC 1386, and the E/S0 TXFS 2226–184 (Braatz et al., 1994; Koekemoeret al., 1995).

NGC 1052 and TXFS 2226–184 (Table III) are particularly interesting.NGC 1052 is a boxy elliptical that may contain a stellar cuspy core (see Sec-tion 2). The maser in TXFS 2226–184 is the most luminous known and may alsobe termed ‘gigamaser’. The measured line profiles do not show the discrete narrowfeatures of width 0.3–10 km s�1 characterizing megamaser emission from spiralgalaxies (Figure 13). Instead, a single broad feature of width 100 km s�1 is seen.Both masers are marginally bright enough to be studied with interferometers atsub-milliarcsecond resolution. It was suspected that the observed lineshapes arehighlighting systematic differences in the nuclear gas distribution between early-and late-type galaxies (Koekemoer et al., 1995). This would be consistent withschemes attempting to unify distinct classes of active galactic nuclei (cf. Falckeet al., 1995). The lack of CO emission from NGC 1052 infers an H2O line bright-ness temperature and source size of>1000 K and<0 :008, respectively (Braatz et al.,1994). A preliminary analysis of NGC 1052 VLBA data (Braatz et al., 1996a)shows that the masers are aligned along the same axis as the radio continuum jet.Hence an association with an accretion disk is not likely. The masers are primarilyconcentrated in two clumps with angular and radial velocity separations of 0 :000004


Figure 13. 22 GHz H2O megamaser spectra: Upper panel: A spectrum from the spiral galaxyNGC 4258, containing the apparently brightest megamaser. Lower panel: A spectrum from theearly-type galaxy TXFS 2226–184, containing the intrinsically most luminous H2O maser. Whilemegamaser emission from spirals shows a superposition of narrow spectral features, filled-apertureprofiles from the two known E/S0 megamaser sources are characterized by a single broad component.

and 50 km s�1. The velocity ‘drift’ detected in single-dish spectra may thus becaused by variatons in the relative intensities of these two maser clumps.

5. Early-Type Galaxies with Cool Dense Gas: Individual Sources

Before discussing observational results in a broader perspective, we first intro-duce some of the most prominent sources in more detail. Very few objects havebeen mapped so far. These include three ‘normal’ but CO rich lenticular galaxiesthat are less massive than the Milky Way and a few elliptical systems, some ofthem interacting. The detection of CO in polar rings is another result worth to bementioned.

5.1. NGC 404, NGC 3593, NGC 4691, THREE ACTIVE S0 GALAXIES

NGC 404 has often been classified as a nearby dwarf elliptical system seen face-on. Optical surface photometry and other data favor however an S0 classification


Figure 14. A comparison between the radio continuum and 12CO distributions in the core ofNGC 3593. The 2.7 GHz radio continuum contours (Jenkins, 1982) superposed on (a) the COJ = 1 � 0 and (b) the J = 2 � 1 distributions. The 1.49 GHz radio continuum contours (Condonet al., 1990) superposed on (c) the CO J = 1�0 and (d) the 2�1 distributions. The maps were alignedwithin the absolute pointing uncertainties of a few arcsec in the CO data, assuming that the mainradio continuum peaks are situated symmetrically on opposite sides of the nucleus. The continuumand CO distributions are displayed with solid and dashed lines, respectively. For the contour levels,see Jenkins (1982), Condon (1990), and Wiklind and Henkel (1992a).

and a larger (�10 Mpc) distance (see Wiklind and Henkel, 1990). NGC 3593 andNGC4691 also appear to be lenticulars. Each of them has been mapped in theCO J = 1� 0 and 2�1 lines with high angular resolution (Wiklind and Henkel,1990, 1992a; Wiklind et al., 1993). In NGC 404, the CO emission arises from aprominent dust lane which may be part of an annular structure around the nucleus.In contrast to NGC 404 and NGC4691, that have almost face-on disks, NGC 3593has a large inclination. This latter object contains a primary outer disk with ascale length of 4000 and a counterrotating secondary inner disk with a scale lengthof 1000 (Bertola et al., 1996), identified through a decomposition of the opticalsurface brightness distribution. Both the warm ionized and the cool molecular gasbelong kinematically to the inner secondary disk. The bulk of the molecular gasis distributed in a ring-like structure of radius 200�350 pc located at the turningpoint of the rotation curve. The diameter of the J = 1 � 0 ring is larger thanthe J = 2 � 1 ring and the angular separation of the two associated peaks ofradio continuum emission (Figure 14). This infers that either there is optically


Figure 15. The distribution of HI and CO in NGC 404. The dashed circle is the optical extent of thegalaxy (from Wiklind, 1991).

thin CO near the inner boundary of the ring or that there is a radial gradient indensity or temperature with highest values at small galactocentric radii. The SB0galaxy NGC 4691 contains several molecular cloud complexes aligned with thestellar bar (see Combes and Elmegreen, 1993, for numerical simulations). Weakemission in a direction perpendicular to the bar is also observed. Optical andnear infrared data reveal four central spatially unresolved hotspots. One or two ofthese exhibit H� emission with a full width to half power (FWHP) linewidth of�V1=2�1200 km s�1, that is blueshifted by �500 km s�1 relative to the systemicvelocity (Garcıa-Barretto et al., 1995).

In Table IV, we compare some of the global parameters of NGC 404, NGC 3593,and NGC 4691. All entries except the absolute blue magnitudes, the molecularmasses, and the star formation rates are independent of the adopted distance.NGC 3593 and NGC 4691 are similar in most of their properties. NGC 404 hasa smaller molecular mass but total neutral gas masses (MH2 + MHI) and neutralgas masses per unit blue luminosity are similar. Thus NGC 404 contains ten timesmore atomic than molecular gas (for the spatial distributions, see Figure 15), whileHI and H2 gas masses are similar for NGC 3593 and NGC 4691. All three galaxies


Table IVA comparison of the lenticular galaxies NGC 404, NGC 3593, andNGC 4691a

NGC 404 NGC 3593 NGC 4691

Distance Mpc 10 7 18M b

B Mag. �19.0 �18.6 �19.6SFRc M� yr�1 0.2 2 8MH2 108 M� 0.7 4 9MH2=MHI 0.1 2.5 1.0MH2 /LB M�/L� 0.01 0.1 0.08Mgas=M

dtot 0.003 0.016 0.017

LFIR/LB 0.1 0.7 1.0LFIR/MH2 L�/M� 8 5 12

a Wiklind et al. (1993); otherwise, see Wiklind and Henkel (1990,1992a).b Corrected for extinction and inclination; MB�=5 :M47.c Derived from the far infrared (FIR) luminosity, see Section 4.1.d Mgas=MH2 +MHI. Mtot derived assuming Mtot/LB= 10M�/L�.

exhibit optical emission lines indicating ongoing or recent star formation. The ratioof star formation rates derived from the far infrared and the blue luminosities, i.e.the current versus average SFR, is of order 20 for NGC 3593 and NGC 4691 butonly 2 for NGC 404. It thus appears that NGC 404 has experienced a period ofenhanced star formation activity but that the rate of star formation has droppedquite recently (� 109 yr). The current SFR of NGC 404 deduced from the infraredluminosity, �0.2 M� yr�1, may be similar to the mass return rate from late-typestars.

5.2. TWO DWARF ELLIPTICAL COMPANIONS OF THE ANDROMEDA GALAXY

The three largest Local Group ellipticals M 32, NGC 185, and NGC205, all com-panions of the Andromeda galaxy, have experienced episodes of star formationwell after their formation (Hodge, 1989). While the compact elliptical M 32 doesnot show a significant ISM, the dwarf ellipticals NGC 185 and NGC 205 stillcontain dust clouds, detectable amounts of HI and CO, and young stars (Baade,1951; Johnson and Gottesman, 1983; Wiklind and Henkel, 1992b; Welch et al.,1996; Young and Lo, 1996, 1997). With their small distances (DMpc � 0.7) bothgalaxies serve as ideal targets deepening our understanding of the star-gas cycle inspheroidal systems.

NGC 185 was the first dwarf elliptical with a detected molecular component(Wiklind and Rydbeck, 1986). It contains a star cluster of age 2–4 107 yrs and showstwo major nuclear dust lanes that are a few 1000 off the nucleus (e.g., Hodge, 1963).HI and CO show no sign of rotation and are spatially separated (Welch et al., 1996),with CO aggregating closer to the bottom of the galaxy’s potential well where the


Figure 16. Red continuum (including H�) image of NGC 205 with HI contours (levels: 5%, 10%,20%, ... 90% of the peak column density of 4� 1020 cm�2). The irregular dark areas near the centerare the dust lanes observed by Hodge (1973). Filled triangles: CO detections; open triangles: COnon-detections. The position of the center is also marked (reproduced from Young and Lo, 1996).

gas pressure may be highest (cf. Spergel and Blitz, 1992). Linewidths are �V1=2 �

30 km s�1; rotation of the gas is not observed. In contrast to the results presented byWelch et al. (1996), the CO data of Wiklind and Rydbeck (1986) and Wiklind andHenkel (1992b) reveal distinct kinematical features. The –200 km s�1 componentappears to be associated with a prominent dust cloud seen in optical images. Weakfeatures at –300 km s�1 and –400 km s�1 may also be present. The intensity of theUV radiation (typically a few percent of that measured for the solar neighborhood) islow even inside the central group of young stars; molecular clouds can thus surviveat very modest H2 densities and HI column densities (n (H2) � 3 � 102 cm�3;N (HI) � 1020 cm�2; cf. Welch et al., 1996; Young and Lo, 1997). Because of theweak UV radiation field and the small dust opacities expected in a low metallicityenvironment, shielding of stellar continuum radiation by dust does not play therole it plays in the Galaxy (see also Sage et al., 1992). If supernova blast wavesremoved the interstellar medium during the last major episode of star formation,


then stellar evolution can account for the presently observed interstellar mass onlyif such small scale ‘starbursts’ have been infrequent. Alternatively, modest starformation events that are sufficiently frequent (every 107�8 yr ) could also provideequilibrium with the returning gas (Welch et al., 1996).

NGC 205 (Figure 16) is nucleated, containing a central star cluster of age�107 yr (e.g., Young and Lo, 1996). Hodge (1973) finds a dozen of dark nebulaewithin 10000 of the nucleus. The highly clumped HI is distributed in an elongated,bent structure that is 900 pc long and 300 pc wide (Young and Lo, 1996, 1997). Incontrast to NGC 185, there is rough agreement in the distributions of HI, CO, andthe dust lanes; the kinematics of the CO and HI emitting gas are similar on scales> 90 pc. Linewidths are or order 10–20 km s�1, with �V1=2;HI >�V1=2;CO for thenorthern and�V1=2;HI <�V1=2;CO for the southern cloud. While HI might envelopethe northern molecular cloud, HI and CO linewidths do not suggest a similarscenario for the southern region. A velocity gradient of order 0.4 km s�1 pc�1 isobserved along the major axis, suggestive of rotation. The global HI profile is offsetby 20 km s�1 from the stellar velocity of the galaxy, suggesting that the gas andthe stars may form two distinct dynamical systems. As for NGC 185, H2 massesbased on a standard conversion factor (see Table II and Equations (1) and (2)) maybe too small by a factor of up to 5; a new ‘mini-starburst’ may occur at any time inthe near future.

5.3. THE POLAR RING LENTICULAR NGC 2685

Polar ring galaxies possess large quantities of luminous matter surrounding thegalaxy’s main body in nearly polar orbits. Most of these galaxies are of type S0.Several sources have confirmed polar rings in the sense that the kinematics ofthe ring and the underlying galaxy have been determined, revealing a large anglebetween the rotation axes. Many other galaxies featuring a similar morphology andincluding >0.5% of the nearby lenticulars (Whitmore et al., 1990) are consideredcandidates. Polar rings contain significant amounts of neutral gas (Richter et al.,1994) and are believed to be formed during galaxy interactions (see Section 6.5).

Following the detection of CO in the polar ring of the (possibly) spiral galaxyNGC 660 (Combes et al., 1992; van Driel et al., 1995), Watson et al., (1994) detectedCO in the polar ring of the lenticular NGC 2685. NGC 2685, the Helix galaxy,has two rings: An inner polar ring (angular radius �2000) with colors consistentwith a normal initial mass function and a photometric age of 2–5 Gyr (Peletier andChristodoulou, 1993) and an outer non-polar ring (angular radius a few arcmin)rotating around the kinematical axis of the central parts of the galaxy. Gas massesare MH2 � 108 M� for the central 1500 (Taniguchi et al., 1990a), MH2 � 3 108 M�

� MHI for the nuclear region and the inner (polar) ring (Watson et al., 1994),and MHI � 2� 109 M� for the entire galaxy (Shane et al., 1980). Most of the COemission observed seems to arise from the polar ring. With a total ofLFIR�109 L�,the galaxy is presently not very active.


The large amounts of neutral gas (Shane, 1980; Taniguchi et al., 1990a; Watsonet al., 1994; Schiminovich et al., 1995) with a rich molecular component makes anaccretion of molecule-poor material from a dwarf galaxy somewhat problematic. Aclose encounter with a gas-rich spiral also poses problems: This could be disruptivefor the entire system and no gas is seen ‘raining back’ into the center. NGC 2685thus provides an interesting case of galaxy interaction and detailed numericalsimulations are desirable.

5.4. THE GIANT ELLIPTICAL RADIO GALAXY CEN A

At a distance of only D � 4 Mpc, Cen A (NGC 5128; see Figure 17) is the nearestand therefore best studied minor-axis dust lane giant elliptical (e.g., Hui et al., 1995;Dobereiner et al., 1996; Kinzer et al., 1995; Schreier et al., 1996, Sparke, 1996;Kellermann et al., 1997). Cen A is suspected to be the product of the merging ofa small gas-rich system with a larger elliptical galaxy (Quillen et al., 1992, 1993).Being one of the brightest extragalactic IRAS point sources, Cen A provides aunique opportunity to observe the detailed dynamical, morphological, physical,and chemical state of its dense cool interstellar gas. A variety of molecular tracershave been observed, including H2, CO, CH, CH+, CN, CS, OH, HCN, HNC,HCO+, NH3, H2CO, HNCO, and C3H2 (see Gardner and Whiteoak, 1976a, b,1979; Whiteoak et al., 1980; Seaquist and Bell, 1986, 1990; Phillips et al., 1987;Bell and Seaquist, 1988; D’Odorico et al., 1989; Eckart et al., 1990a, b; Israel et al.,1990, 1991; Israel, 1992; Quillen et al., 1992; Rydbeck et al., 1993; van Langeveldeet al., 1995; Wiklind and Combes, 1997). A summary of chemical properties isgiven by Henkel and Mauersberger (1991).

Four major molecular gas components have been identified: (1) Narrow (�100km s�1) line emission from the optically visible dust lane is seen in CO andperhaps in H2CO, HNCO, and C3H2; (2) broad (>100 km s�1) line emission fromthe circumnuclear gas is observed in CO, CS (W. Wild, priv. comm), HCN, HNC,HCO+, C3H2, and perhaps in H2CO and HNCO (see Figure 18); vibrationallyexcited H2 is also detected; (3) narrow line absorption near the systemic velocity,at �550 km s�1, can be found in almost all detected molecular lines (Figure 18);exceptions are H2, HNCO, and CH+. In a few transitions, the component is alsoobserved in emission; (4) red-shifted absorption lines at �600 km s�1 are observedin CO, HCO+, OH, and perhaps H2CO. C3H2 was also detected (J. B. Whiteoak,priv. comm.). HI is associated with components (1) and (3) and is also seen in theouter shells of the galaxy (van Gorkom et al., 1990; Schiminovich et al., 1994). Inthe following, we address the four molecular components separately.

(1) The optically visible dust lane containsMHI � 7 108 M� (van Gorkom et al.,1990) and MH2 � 108 M� (e.g., Eckart et al., 1990b). The kinetic temperaturederived from CO spectra is � 10–15 K. While the intrinsic velocity dispersion is� < 10 km s�1, the density might be as high as n (H2) = 104 cm�3 (e.g., Quillenet al., 1992; Wild et al., 1996). The globally averaged ratio between far infrared


Figure 17. Overlay of a numerically simulated warped disk in a prolate potential on an optical imageof Cen A (Quillen et al., 1992). The model distribution provides a good fit to the observed morphologyand kinematics of the dust lane.

luminosity and molecular mass inferred from CO is of order 50L�/M� (see TableII) which is larger than in the galactic disk (�4L�/M�; e.g., Sanders 1991). The farinfrared emission appears to be more extended than the CO emission, indicating thepresence of diffuse clouds (Eckart et al., 1990b). Morphology and kinematics of themolecular gas are similar to that of the ionized material seen in H� (Quillen et al.,1992). The profile shapes are well fitted by a warped disk in a prolate gravitationalpotential with r � 3 kpc (see Figure 17). This model does not only reproduce theCO distribution but is also consistent with HI and near infrared data and with theoptical appearance of the dust lane (Quillen et al., 1992).

(2) The circumnuclear cloud, first detected by Israel et al. (1990), forms amolecular ring which is seen almost edge-on and which is oriented perpendicularto the axis of the nuclear radio continuum jet (Figure 19). Analyzing the gas withvelocities differing by more than 170 km s�1 from the systemic velocity, Rydbecket al. (1993) find r � 70� 120 pc, vrot �220 km s�1, and vLSR�560 km s�1 for thecloud, centered on the compact nuclear radio continuum source. Dust emission isalso detected from this region (Eckart et al., 1990b; Hawarden et al., 1993). WithMH2 � 107 M�, the molecular mass is about 1% of the dynamical mass. Spatially


Figure 18. Cen A HCN and HCO+ spectra, showing broad nuclear emission and narrow foregroundabsorption (from Israel, 1992).

unresolved H2 emission (�400) may arise from vibrationally excited moleculescoating the inner edge of the circumnuclear structure. Av � 27m towards thenucleus of which� 20m may be due to to the circumnuclear molecular gas (Israelet al., 1990).

Most of the line intensity ratios found in the circumnuclear gas are, within afactor of two, similar to those observed in other active nuclear regions of highinfrared luminosity (cf. Table 3 of Israel, 1992). The J = 1 � 0 12CO/13CO lineintensity ratio,�10, is similar to those measured in M 83 and in the starburst galaxyNGC 253. HCO+/13CO, HCN/13CO, and HNC/ 13CO line intensity ratios are con-sistent with those measured in NGC 253, demonstrating that high density clumps(n (H2)> 104 cm�3) must be quite widespread. The HCN/HNC line intensityratio, �2, requires Tkin>10 K (e.g., Schilke et al., 1992). The CO line intensityratios do not strongly constrain the admissable Tkin range but are consistent withtemperatures well above 10 K.

(3) If sensitivity and frequency resolution are high enough, three absorptionfeatures near the systemic velocity can be distinguished. Their velocities are vLSR =539, 544, and 550 km s�1. These are also seen in the 21 cm HI absorption profile(van der Hulst et al., 1983). In all 9 cm CH transitions (Whiteoak et al., 1980) andin some of the 18 cm OH transitions (van Langevelde et al., 1995) the componentsare observed in emission. From an analysis of HCN hyperfine components, Eckartet al. (1990a) propose that the absorption arises from cool, dense, and clumped


Figure 19. Deconvolved CO J = 2�1 emission line intensities towards the Cen A nuclear continuumsource. The angular resolution of the map was improved by the application of a Statistical ImageAnalysis technique. Radio contours are given as solid lines (from Rydbeck et al., 1993.)

gas (NH2 � 3 � 5 � 1021 cm�2 for [CO]/[H2] = 10�4; n (H2)� 104 cm�3; Tkin<�20 K) with properties similar to those of the extended CO clouds in the opticallyvisible dust lane. Seaquist and Bell (1990) and Wiklind and Combes (1997) favordiffuse lower density gas. Based on the HCN/HNC J = 1 � 0 line intensity ratio(Wiklind and Combes, 1997), Tkin > 10 K. While this may support an associationwith the nuclear region, the small line widths and the almost systemic velocitiessupport an association with the optical dust lane at larger galactocentric radii. Acomparison of the�550 km s�1 component column densities with those of cloudsin the galactic disk and other galaxies indicates normal molecular abundances(Eckart et al., 1990b; Wiklind and Combes, 1997). Towards the inner part of theradio continuum jet, � 2000 NE of the nucleus, 110 � 111 (6 cm) H2CO absorptionwas also detected (Seaquist and Bell, 1990). Along the lines of sight to the nuclearcontinuum source and the inner jet, H2CO column densities appear to be similar.

HCO+ spectra, separated by seven years, show no significant change. Thisallows to constrain the size of the nuclear radio source at � = 3 mm to > 500 AU(Wiklind and Combes, 1997).

(4) The redshifted absorption features, also seen in HI (van der Hulst et al.,1983), are observed in the range vLSR = 570� 620 km s�1. HCN/HNC J = 1� 0line intensity ratios suggest that the gas is cool (Tkin <� 10 K; Wiklind and Combes,1997). While the temperature can be used as an argument against an association withthe nuclear region, the large velocity range of the absorbing gas, the multitude ofvelocity components, and the deviation from systemic velocity are in favor of suchan association, involving circumnuclear gas falling towards the nucleus. HCO+/CO


and HCO+/HCN J = 1�0 line intensity ratios are larger than towards the systemicabsorption features. With � = �ln (1+Tb=F Tc) and

R� dv / N Aul �

�3 / N �20

(Tb (negative) and Tc (positive) denote line and continuum brightness temperature,F is the continuum source covering factor of the absorbing cloud,N is the columndensity, andAul is the Einstein coefficient for spontaneous emission) the opacity fora given column density depends critically on the molecule’s electric dipole moment,�0. While CO tends to be much more abundant than other species (like HCN), thisis almost compensated by its �300 times smaller electric dipole moment. HCO+

and HCN have similar rotational constants and electric dipole moments, so that acomparison of these two species is relatively easy. The outstanding HCO+/HCNline intensity ratios (�1; Eckart et al., 1990a; Wiklind and Combes, 1997) imply anHCO+ enhancement caused either by cosmic ray ionization or by shock inducedprocesses (Mitchell and Deveau, 1983; Farquhar et al., 1994). With respect tothe HCO+/HCN line intensity ratio, the south-western HCO+ hotspot in M 82(Carlstrom, 1988) and clouds near sites of massive star formation in the Magellanicclouds (Johansson et al., 1994; Chin et al., 1996, 1997; Heikkila et al., 1997) appearto be chemically related.

5.5. NGC 1275, A CLUSTER DOMINANT CD GALAXY

The cD NGC 1275 (Per A, 3C 84), the dominant galaxy in the Perseus cluster(A 426; D � 70 Mpc), is an outstanding object: The X-ray emission of the clusteris commonly interpreted as the signature of a cooling flow with an inflow rateof _Mx � 300M� yr�1 (see Section 6.4.2). The Seyfert nucleus of NGC 1275 isassociated with a powerful nonthermal radio source with FR I jet and (perhaps)counterjet (e.g., Vermeulen et al., 1994; Walker et al., 1994; Levinson et al., 1995);the central region is a strong source of thermal far infrared emission (Lester etal., 1995) and contains a population of compact luminous objects (possibly youngglobular clusters; e.g., Holtzman et al., 1992; Richer et al., 1993; Nørgaard-Nielsenet al., 1993; Ferruit and Pecontal, 1994; Zepf et al., 1995; Kaisler et al., 1996).The most outstanding property of NGC 1275 is, however, the presence of twooptically identified kinematical components (Minkowski, 1957). Optical emissionis observed at the systemic recessional velocity of 5250 km s�1 and at 8200 km s�1

(hereafter ‘LV’ and ‘HV’ system). The HV system may be moving through thewestern part of the nuclear region at large velocity.

Neutral hydrogen absorption was detected in both velocity components (deYoung et al., 1973; Crane et al., 1982; Jaffe, 1990). A few narrow (�V1=2 �

5 km s�1) HV features and a single broad LV feature (a few 100 km s�1 wide) wereobserved (for the line profiles, see Section 6.4.2). A large scale line of sight magneticfield component of –9.3�2.7�G was tentatively derived from HI measurements ofthe HV system (Kazes et al., 1991).

Molecular gas from the LV component was first detected by Fischer et al.,(1987) in H2, showing a compact core and faint emission possibly extending in


Figure 20. Distribution of integrated CO J = 1–0 emission towards the cluster dominant cD galaxyNGC 1275. Beam size: �b = 5 :001 � 3 :006. Contour levels are 20, 30, 40...70 mJy per beam. Negativecontours are indicated by dashed lines (taken from Inoue et al., 1996).

east-west direction up to 900 from the nucleus (Kawara and Taniguchi, 1993; Inoueet al., 1996). CO emission from the LV component was reported by Lazareff et al.(1989), Mirabel et al. (1989), Reuter et al. (1993), Braine et al. (1995), and Inoueet al. (1996) with a spatial extent similar to that of the 100�m emission (a few1000; Lester et al., 1995). The total molecular mass is MH2 � 3 � 109 M� whichis remarkable for a centrally located cluster dominant galaxy. Reuter et al., (1993)find a central, an eastern, and a western component, the latter with � 1500 (4.5 kpc)offsets from the nucleus. The interferometic measurements of Braine et al. (1995)with their high spatial resolution (1 :007) and narrow velocity channels (6.6 km s�1)only led to a marginal CO detection. A better choice of observing parameters(Inoue et al., 1996; see Figure 20) resulted in a confirmation of an overall east-west


orientation of the cool gas. CO peaks are found 3–400 south-east and north-west ofthe nucleus, suggesting rotation with Vrot � 200 km s�1 around the nucleus.

5.6. JOINING...

Galaxy mergers are spectacular cosmic events involving huge stellar masses andintense radiation fields. They are often accompanied by high concentrations ofmolecular gas and dust and by enhanced rates of star formation. The remnantgalaxy after the merger may differ significantly from those of the progenitors.

NGC 7252 (Arp 226; D � 60 Mpc), the ‘Atoms for Peace’ galaxy, providesan important link between galaxy mergers and the formation of elliptical galaxies(see Figure 21). NGC 7252 shows a bright central region with delicate ripples,faint surrounding loops, and ‘E + A’ or ‘post-starburst’ spectral signatures. Twolong tidal arms point towards the east and northwest and project to �100 kpc fromthe center. This is the characteristic signature of a merger of two similarly massivedisk galaxies (Schweizer, 1982; Borne and Richstone, 1991; Hernquist and Barnes,1991; Mihos et al., 1993; Hibbard and Mihos, 1995). With outer optical isophotesthat are (in spite of the tidal arms) almost elliptical and with an outer enveloperadial light distribution following the r1=4 law (cf. Schweizer, 1982; Stanford andBushhouse, 1991), the system can be classified as a young elliptical merger remnantwith a part of the tidal arm material returning to the central region (Hibbard andMihos, 1995). The classification as a merger is further supported by the complexkinematics of the ionized gas which reverses its sense of motion at radii of �1000

(�3 kpc; Schweizer, 1982). The system still contains huge amounts of cool neutralgas (MH2 � MHI � 109 M�; Dupraz et al., 1990; Wang et al., 1992b; Hibbardet al., 1994). HI is concentrated exclusively in the outer tidal region of the galaxy,in the two arms, and in the ‘Western Loop’ next to the central body of the galaxy.Near the end of each tidal arm there is a giant HII region associated with largeconcentrations of gas and stars which approach the mass and linear extent of dwarfgalaxies. The CO emission is confined to an inner ring of radius 0.3–1.5 kpc. Thisinner ring has an H� counterpart and both components counterrotate w.r.t. theionized gas of the outer galaxy. The H� emission as well as the high far infraredluminosity (see Table II) indicate ongoing star formation at rates that are muchhigher than in most other elliptical galaxies (cf. Fritze-v. Alvensleben and Gerhard,1994). During the next stage of evolution the molecular gas may be converted intostars. Since the gas mass is small relative to the total mass, an inconspicious stellardisk may form and the future object might evolve to a disk-shaped giant ellipticalgalaxy.

5.7. ... AND LEAVING THE FAMILY OF EARLY-TYPE GALAXIES

While NGC 7252 and NGC 3597 (Dupraz et al., 1990; Taniguchi et al., 1990b)are examples of galaxies joining the family of early-type galaxies (for further


Figure 21. A blue photograph of NGC 7252, reproduced from Schweizer (1982). The length of thebar corresponds to 10 (17.5 kpc).

candidates, see Vigroux et al., 1985; Sancisi et al., 1987), the amorphous irregularNGC 3077 demonstrates that some galaxies evolve into the opposite direction. Along tidal HI arm connecting NGC 3077 with M 81 shows that the galaxies areinteracting (e.g., van der Hulst, 1979; Thomasson and Donner, 1993). Optically,NGC 3077 does not appear to be an early-type galaxy. However, near infrareddata show a luminosity distribution similar to that of the dwarf elliptical galaxyNGC 185 (Price and Gullixson, 1989). In addition to the HI gas from the tidalarm, there exists a large molecular complex with MH2 � 107 M� (D � 3 Mpc;Becker et al., 1989) providing the fuel for star formation at a rate of�0.1M� yr�1.Per surface area, this rate is two orders of magnitude higher than in disks of spiralgalaxies (Thronson et al., 1991). During the last 108 yrs, the newly formed stars andthe dust associated with molecular clouds have altered the optical morphology ofNGC 3077; in the future, the star forming process will also change the morphologyin the near infrared, thus eliminating the signatures of its previous stage of evolutionand making an assessment of the galaxy’s history more and more difficult.


6. The Neutral Gas Phase: Impact and Relevance for a BetterUnderstanding of Early-Type Galaxies

Focusing our attention on the properties of the cool dense ISM with its HI, molecu-lar, and far infrared emission, we have not yet discusssed important implicationsrelated to phenomena mainly observed in other parts of the electromagnetic spec-trum. In the following we present a more systematic comparison between the coolgas phase and other interstellar gas components. Firstly, correlations with optic-al (and submillimeter) measurements of the dust are analysed. Then we addressaspects related to the radio continuum emission, to the warmer ionized gas with itsoptical emission lines, and to the hot gas in the X-ray halos of early-type galaxies(Sections 6.1–6.4). The latter is discussed in greatest detail, not so much because itis the most complex aspect, but because it provides fresh views on neutral gas cloudsin an exotic environment, because it severely constrains the fate of the cooling gas,and because it is relevant to galaxy evolution and to estimates of the baryonic massof the universe. The origin of the cool neutral gas, rotation curves, mass-to-lightratios, and the presence of massive stars are discussed in Sections 6.5–6.7.

6.1. DUST IN ELLIPTICAL GALAXIES

6.1.1. Extended Dust LanesIn Section 4, molecular data were compared with IRAS flux densities. However,IRAS alone does not provide a complete picture of the dust component; the cooldust (T < 30 K), usually responsible for the bulk of the dust mass, is not seenand the angular resolution (a few arcmin) is too coarse for a detailed analysis.Hence a complementary discussion of optical data is needed to establish the finescale distribution of the dust. Optically, we are confronted with serious selectioneffects. Optical methods detecting dust in galaxies take advantage of the smoothdistribution of the stellar light. Thus only dust distributions that are distinct fromthe light distribution of the stellar background can be observed. In addition, theprobability to optically detect dusty disks, lanes, patches or wisps depends heavilyon orientation effects and is biased towards edge-on distributions.

First systematic studies of the cold dust date back almost two decades (e.g.,Bertola and Galletta, 1978; Hawarden et al., 1981; Ebneter and Balick, 1985;Sadler and Gerhard, 1985; Sparks et al., 1985; Ebneter et al., 1988; Veron-Cettyand Veron, 1988) and showed that dust lanes often lie parallel to the opticalmajor axis in seemingly oblate (e.g., NGC 3665) and parallel to the minor axisin seemingly prolate (e.g., Cen A) elliptical galaxies; a minority of ‘skew’ caseswith the dust at an intermediate angle also exists. Furthermore it was shown thatdust lanes are common, a result consistent with the IRAS observations. Studies of3CR sources (radio galaxies of the third Cambridge Catalogue; see Laing et al.,1983; Spinrad et al., 1985) indicate that there is no change in the fraction of largeelliptical galaxies showing dust out to z = 0:5 (Baum et al., 1995a). While datafrom the (sub)millimeter wave region are still scarce (see Fich and Hodge, 1993;


Wiklind and Henkel, 1995, for first observational results), a detailed comparisonwith IRAS data using the technique of co-adding survey scans (Knapp et al., 1989)is possible.

The most detailed optical study, involving deep narrow- and broad-band imagingof a magnitude limited (Bo

T <12 :m0) sample of 56 RSA ellipticals (Sandage andTammann, 1987) was presented by Goudfrooij et al., (1994a, b, c) and Goufrooijand de Jong (1995). In general, extinction curves as a function of wavelength arefound to run parallel to the galactic extinction curve. The ratio of total to selectiveextinction, Rv = Av=E (B � V ), falls into the range 2.1� Rv �3.3 and is thussmaller on average than the canonical galactic value, 3.1, also found in other spiralgalaxies (e.g., Elmegreen, 1980; Knapen et al., 1991). A smaller Rv for ellipticals(e.g., Warren-Smith and Berry, 1983; Brosch et al., 1985; Brosch and Loinger,1991; Pinkney et al., 1996) hints at smaller dust grain sizes. Those galaxies withapparently small characteristic grain sizes have smooth regularly distributed dustlanes whereas the dust lanes with normal grain sizes appear to be patchy or irregular.This can be explained by an aging effect: Grain sizes slowly decline in a hostileenvironment (e.g., Tsai and Mathews, 1995; Voit and Donahue, 1995), while thedust and the associated gas are kinematically relaxing (see also Section 6.4).

A comparison of total amounts of dust deduced from optical data and far infraredfluxes yields a important inconsistency: Goudfrooij and de Jong (1995) obtain ‘FIRdust masses’ (typically 105�6 M�) that are larger by an order of magnitude thanthose deduced from optical data. Not all galaxies are showing this effect (Wiseand Silva, 1996), a part of the reddening may be compensated by scattering (Wittet al., 1992), and dust properties can differ from those encountered in the Galaxy(as outlined above). Nevertheless the discrepancy between optical and far infrareddata is not a sensitive function of the inclination of the dust lane and has to be takenseriously. To explain the observational results, Goudfrooij and de Jong postulatean extended diffuse dust component which is undetectable using optical methods,but that is responsible for (a part) of the observed radial color gradient in ellipticals(see also, e.g., Peletier, 1987; Eckart et al., 1990b; Nørgaard-Nielsen et al., 1993;Reid et al., 1994; Wise and Silva, 1996)

A detailed comparison of the dust, HI, and molecular gas distribution has not yetbeen made: the angular extent of most galaxies is too small to spatially resolve the� = 21 cm HI and 2.6 or 1.3 mm CO emission with the beams usually employed.Interferometric data resolve out and thus loose smooth and extended HI or COcomponents; even if the source is exceptionally extended (for Cen A, see vanGorkom et al., 1990; Eckart et al., 1990b), large filled-aperture HI beams andthe weakness of individual CO features makes complete fully sampled mappingdifficult. For NGC 1275, a comparison of the optical extinction of the patchydust (Nørgaard-Nielsen et al., 1993) with the 100�m dust emission (Lester et al.,1995) and the molecular gas (Reuter et al., 1993) indicates a similar extent for allthree components (�3000). At least in this case, the hypothesized extended dustcomponent is not associated with significant amounts of molecular gas.


Figure 22. V � I color index image of NGC 4261 with the dark reddened dust lane extending fromthe center 1500 to the north (from Mollenhoff and Bender, 1987).

6.1.2. Dust from the Nuclear RegionsWe have already mentioned NGC 1052, NGC 1386, Cen A, and TXFS 2226–184(Sections 3.4 and 4.3), where the molecular emission from the nuclear region mustbe accompanied by large amounts of dust. Towards some ellipticals, Knapp et al.(1990) find a thermal dust component with color temperatures well in excess of30 K; they assume that the emission is arising from the central regions of thosegalaxies, where an active nucleus provides a significant input of energy. HSTobservations of early-type galaxies at D = 10 � 20 Mpc, observed with a linearresolution of a few pc only, reveal previously unknown central disks and patchydust clouds that seem not always to be dynamically relaxed (Jaffe et al., 1994; vanden Bosch et al., 1994; Forbes et al., 1995; Lauer et al., 1995; van Dokkum andFranx, 1995; Byun et al., 1996; de Juan et al., 1996; Carollo et al., 1997). Detection


rates in radio galaxies appear to be higher than in radio quiet sources. Towards theS0 galaxy NGC 3998, Knapp et al. (1997) find both extended 4.5–15�m emissionfrom the bulge stars and strong point source emission from the nucleus. The innerregion of NGC 3998 appears to contain a few 1000M� of warm (Td � 200 K) dust,possibly associated with an AGN. Another source of interest is the Virgo E2 radiogalaxy NGC 4261 (3C 270), suspected to contain a massive black hole at its center.Mollenhoff and Bender (1987), Jaffe et al. (1993, 1996), and Mahabal et al. (1996)find a dust lane elongated almost perpendicular to the double radio jet (Figures 22and 23). The dust lane may also be seen in HI (Jaffe and McNamara, 1994). Thenucleus is slightly displaced from the very center of the inner dust disk (1 :006 �0 :007) and both disk and nucleus are slightly displaced w.r.t. the isophotal center,indicating a merging event (Ferrarese et al., 1995). Only a tentative detection ofCO has so far been reported for the J = 2� 1 line (Jaffe and McNamara, 1994);sensitive CO J = 1� 0 measurements with the 30-m IRAM telescope (C. Henkeland T. Wiklind, unpublished) do not confirm this detection.

6.2. THE RADIO CONTINUUM

The radio continuum in early-type galaxies is often dominated by non-thermalradiation related to a compact active nucleus (for a classification of the observedphenomena, see Fanarov and Riley, 1974; Baum et al., 1995b; Colina and de Juan,1995; Zirbel, 1996, 1997). Table I displays three classical radio galaxies that aredetected in CO: Per A (NGC 1275), For A (NGC 1316), and Cen A (NGC 5128).For a more distant radio galaxy (4C 12.50 at czCO=36:650 km s�1) that is alsoultraluminous in the far infrared (log (LFIR/L�) = 12.1), see Mirabel et al. (1989).Other prominent radio galaxies also show evidence for opaque nuclear dust lanes(see, e.g., Carilli and Barthel, 1996, for Cyg A), but Hyd A, Vir A, and Cyg Aremain undetected in CO (Heckman et al., 1989; Mirabel et al., 1989; Braineand Wiklind, 1993; Sofue and Wakamatsu, 1993; Barvainis and Antonucci, 1994;Braine and Dupraz, 1994; McNamara and Jaffe, 1994; O’Dea et al., 1994b). Is therea connection between phenomena in the nuclear region and the cool dense gas? Arethere indications that the gas is feeding a central engine? While lenticular galaxiesoften are weak radio continuum emitters (e.g., Bieging and Biermann, 1977), giantellipticals can be strong radio sources, showing a double lobed structure extendingfar beyond the optical size of the galaxy. IRAS data indicate that giant ellipticalswith strong radio continuum emission have higher infrared luminosities at a givenoptical luminosity than ‘normal’ early-type galaxies (Impey et al., 1990; Knappet al., 1990). There also exists a correlation between the radio continuum powerand the presence of optically visible dust lanes (e.g., Buson et al., 1993). Whilein spirals the non-thermal radio continuum is well correlated with the far infraredluminosity (e.g., Wunderlich et al., 1987), such a linear correlation does not holdfor ellipticals and, to a lesser degree, for lenticulars. The lack of a good correlation


Figure 23. HST greyscale image of NGC 4261. Size 3 :0035� 3 :0055. North is towards the upper right(P.A. = –20�). Four isophotal ellipses are drawn. For the innermost and outermost ellipses, majorand minor axes are also displayed. The crossing points, the ‘isophotal centers’, are slightly displaced(from Jaffe et al., 1996), possibly indicating an accretion or merging event.

may be due to radio emission mainly arising from the nucleus (Bregman et al.,1992).

HI absorption has not been detected towards strong, presumably beamed, coredominated radio sources. This result is consistent with nuclear HI absorption pref-erentially arising in a plane perpendicular to the radio jet. z < 0:1 HI absorptionhas been observed towards a total of 17 radio galaxies (cf. van Gorkom et al., 1989;McNamara et al., 1990; Jaffe, 1992; O’Dea et al., 1994a; Conway, 1997). Detec-tions were obtained towards several FR I galaxies but rarely towards FR II sources(for an exceptional case, the FR II Cyg A, see Conway and Blanco, 1995). The lackof FR II detections may be ascribed to their high redshifts and weak radio cores thatmake observations difficult. A surprisingly large fraction of HI absorbers has been


detected among the rare Compact Symmetric Objects (CSO’s) and the related SteepSpectrum Core (SSC) sources (Conway, 1997). The CSO’s show the same double(or triple) morphology as classical radio sources but they are smaller by a factor of1000 or more. The SSC sources have radio lobes more than 100 kpc away from acentral steep spectrum core that shows CSO morphology (cf. Barthel et al., 1985;Wilkinson et al., 1994). The CSO’s might be in an early stage of evolution towardsa classical radio source; SSC objects may represent old radio sources that havestarted, after a quiescent period, another phase of activity. The high incidence of HI

absorption in CSO and SSC sources, if significant (there are 5–10 sources detectedso far), can be taken as evidence that these radio morphologies are preferentiallyassociated with cool gas or that these sources form ideal background objects inwhich to see absorption. Since the sizes of CSOs are comparable to those of 100 pcscale circumnuclear disks (see Section 6.1.2), we may speculate that in a largefraction of these galaxies the entire radio mini-lobe will be covered by the HI disk.Preliminary spectral line VLBI measurements (Conway, 1997) are not inconsistentwith this hypothesis.

Is there an obvious connection between outer dust lanes and nuclear cores?High luminosity early-type galaxies (LB > 3 � 109 L�) known to contain HI

likely possess a nuclear radio source. While HI mass and radio power are notproportional (Knapp, 1987), the correlation nevertheless suggests that parts ofthe cool gas component are feeding the central engine. The observed redshiftedabsorption features towards Cen A and other sources (van Gorkom et al., 1989;see also Sections 5.4 and 6.4.2) might arise from such infalling gas. The linkbetween the cool gas component and nuclear activity is further strengthened by thecorrelation between dust lane and radio axis orientation (Figure 24), first noted byKotanyi and Ekers (1979; see also Mollenhoff et al., 1992; Baum et al., 1995a; vanDokkum and Franx, 1995; for a similar correlation with warm ionized gas disks,see Bertola, 1992): The gaseous rings or disks associated with the dust lanes areoriented almost perpendicular to the radio jets or to the line connecting the radiolobes. This holds whether the dust lane is aligned to the optical major axis (as, e.g.,in the case of NGC 3665) or whether the dust lane is perpendicular to the opticalmajor axis (Cen A). A correlation between optical (major) axis and radio (jet) axisis not observed (e.g., Birkinshaw and Davies, 1985; Sansom et al., 1987). From a3CR sample containing 32 sources with z < 0:5, Baum et al. (1995a) conclude thatthe optical axis aligns best with the radio axis when the dust is oriented along theminor optical axis (such as in Cen A), suggesting a connection between alignmenteffect and dust distribution. The 3CR sample shows a significant dispersion in theradio to dust alignment: two of the 32 sources (3C 305, 3C 443) show jets directlypointing into the dust disk. In spite of these two sources, however, statisticalevidence strongly favors a link between the angular momentum of the dust laneat kpc scales and the nuclear region that determines the jet axis. Apparently, theexistence of a dust lane is an effective trigger of an active nucleus.


Figure 24. Distribution of the difference between the position angle of dust lanes and radio jets for29 sources. The different classes mark different detection qualities of dust and radio features. For theclear detections (class I) there is a strong tendency for orthogonal configurations (from Mollenhoffet al., 1992).

6.3. THE WARM GAS COMPONENT

In the Milky Way, the warm ionized gas component (T � 104 K) comprises atotal mass of several 107 M� (Miller and Cox, 1993). Being mainly heated bystars more massive than the Sun, the galactic mass ratio of warm to cool gasis � 0:01. In early-type galaxies, the warm ionized gas component is even lessconspicuous compared to the combined mass of the cool and hot gas components.Almost 50% of the 300 early-type galaxies studied by Caldwell (1984) and Phillipset al. (1986) were detected in either [O II] �3727 or [N II] �6584, with typicalionized gas masses of 102 � 104 M� and minimum equivalent widths of 1 and


0.8 or 0.4 A, respectively. The detection rates for elliptical and lenticular galaxiesare similar and reach 70–85% when using longer exposures (Macchetto et al.,1996). In radio-weak early-type galaxies the emission line luminosity is correlatedwith the absolute magnitude of the parent galaxy. Luminous galaxies are morelikely to show emission than smaller galaxies. 70% of the early-type galaxies withMB < �20M have an ionized gas mass of at least 103 M�. The linear correlationbetween emission line and stellar luminosities in radio-weak (and thus often lowluminosity) early-type galaxies suggests that the dominant ionizing source must bea component of the stellar population, either old evolved or massive young stars (cf.Burstein et al., 1988; Greggio and Renzini, 1990). High relative velocities betweenthe interstellar and circumstellar gas components provide additional heating (e.g.,Brighenti and Mathews, 1996; Renzini, 1996).

Radio galaxies are more likely to show emission then galaxies of the same stellarluminosity without a strong radio continuum source (e.g., Morganti 1992). Theionized gas mass in the nearby galaxies studied by Buson et al. (1993) is typically10–100 times that of a ‘normal’ early-type galaxy of similar stellar luminosity.These galaxies show significantly enhanced radio continuum fluxes (typically byfactors of 10–15) and appear to have an unusually high fraction of cold gas beingionized rather than being exceptionally gas rich. The extra ionizing source maybe related to an active nucleus (for the warm ionized gas component reaching107 M� in even more luminous radio galaxies, see, e.g., Robinson et al., 1987,and Section 6.4.2). While a correlation between warm gas and HI masses couldnot be established (Balkowski et al., 1986), nearly all the observed dusty galaxiesshow emission from the warm ionized gas component. In most sources, the warmgas is concentrated in the inner few kpc. The observed rapidly rotating HI disksor rings (for an exceptional case, see Schweizer et al., 1989) and the presence ofdust within the optically visible emission regions favor an association between thewarm ionized and the cool neutral gas.

While luminous ellipticals show LINER spectra (e.g., LH� < L[NII]), dwarfellipticals show spectra resembling those of galactic H,II regions (LH� > L[NII];Phillips et al., 1986; Sadler, 1987). The gas is visible over 0.5–1 kpc and oftenshows no sign for rotation (cf. Phillips et al., 1986; Bothun and Mould, 1988; Heldand Mould, 1994). The extent of optical line emission appears to be similar to thatof CO in those dwarf galaxies, where molecular line emission is detected (e.g.,Sage et al., 1992). The number of such dwarfs is small, however.

6.4. HOT VERSUS COOL GAS: COOLING FLOWS

In giant elliptical and cD galaxies, the hot gas dominates the interstellar medium(MX > MHI+H2 +MHII). In spite of this situation, data from the cool neutral gascomponent are essential to determine the fate of the cooling X-ray emitting gas.Optical and X-ray luminosities are correlated, with a large dispersion, by LX/LB

� (1 � � � 7 with increasing exponent for larger LB), so that the hot coronal


gas contributes most of the X-ray emission in luminous (LB>1010 L�) ellipticals.In less luminous objects discrete stellar sources may provide a significant fractionof the total X-ray flux (e.g., Canizares et al., 1987; Valentijn, 1988; Fabbiano et al.,1989; Donelly et al., 1990; Sarazin, 1990; White and Sarazin, 1991; Bregman et al.,1992; Eskridge et al., 1995; Fabbiano, 1996a). Most luminous X-ray emission,sometimes in excess ofLX=1011 L�, is associated with a high (>30%) percentageof galaxy clusters and is presently observed out to z � 0:5 (Donahue and Stocke,1995). The extended X-ray halos associated with individual galaxies, groups orclusters emit thermal emission from diffuse gas at temperatures of 107 � 108 K(Sarazin, 1990; Fabian, 1994).

Most models of the gaseous halos involve ‘cooling flows’. Such a cooling flowforms when the cooling time of the gas is smaller than the age of the system. Asthe gas begins to cool, its density increases and its volume shrinks, so that the gaspressure continues to support the weight of the outer quasi-hydrostatic layers. In anhomogeneous and spherical coronal halo, the cooled gas then moves towards thecenter of the associated galaxy or cluster. To establish a radial temperature gradientor to cool clouds within the hot gas, electric conductivity must be reduced by a feworders of magnitude. This may be caused by the presence of tangled magnetic fieldsor by plasma microturbulence (e.g., Jafelice, 1992; Fabian, 1994; Pistinner et al.,1996; Pistinner and Shaviv, 1996). An often adopted model (for recent reviews,see Sarazin, 1990; Fabian et al., 1991; Fabian, 1994) is based on the ‘standard betalaw’

SX (r) = So

�1 + (

r

rx)2��3�+1=2

(5)

and

n(r) = no

"1 +

�r

rx

�2#�3�=2

; (6)

where So and no are the fluxes and number densities at the center of the flow, rx

is the core radius, and � (fits to the X-ray surface brightness give � � 0:6; Jonesand Forman, 1984; Fabbiano, 1989) is an adjustable parameter which is unity inthe case of ideal hydrostatic equilibrium. The standard beta law fits the observedazimuthally averaged X-ray surface brightness SX reasonably well, except in thecentral regions, where anLX excess is sometimes observed (e.g., Jones and Forman,1984).

What is the fate of the cooling gas? A linear correlation extending over fiveorders of magnitude exists between LX and the luminosity limit of the radiocontinuum, where a given fraction of the galaxies is detected (e.g., Valentijn andBiljeveld, 1983; Valentijn, 1988; Eskridge et al., 1995). This and the link betweenthe angular momentum vector of dust lanes at kpc scales and the radio continuumjet axes (Section 6.2) suggests that at least some of the gas is accreted to the nuclear


source, where it may feed an active nucleus. For the bulk of the cooling gas, twopossibilities appear to be particularly compelling:

(1) Most of the material is deposited at the center of the flow where the gas iscool enough to form stars.

(2) Most of the gas is converted into stars well before it reaches the centralregion, leading to an extended stellar halo that should be detectable in some form.

If the gas cools to form stars or substellar objects, then a cool neutral gas phasemust exist, where HI and H2 are abundant. The boundary conditions for such neutralclouds w.r.t. X-ray heating, cosmic-ray flux, and dust grain properties must be quitedifferent from those observed in the Milky Way. Thus a thorough discussion of suchscenarios is relevant not only for a better understanding of early-type galaxies andcooling flows; is is also fundamental for questions related to physics, chemistry, andevolution of molecular clouds. Furthermore, the reported high cooling rates bearsignificant cosmological implications in regard to the baryonic mass fraction of theUniverse. Are dark galactic halos formed by star formation in cooling flows withan IMF extremely biased against high and intermediate mass (M > 0:5M�) starformation? Because of this wide impact and because reviews emphasizing the roleof the cool gas are missing, the (lack of a) neutral gas phase in cooling flows will bediscussed in greater detail than the warm gas and the radio continuum (Sections 6.2and 6.3). In the following, we first summarize the properties of the cool and hotgas associated with individual galaxies. The more thoroughly investigated clustercooling flows will be discussed in Section 6.4.2.

6.4.1. Individual Galaxies with X-Ray HalosModels of the X-ray emission in early-type galaxies distinguish several stages ofevolution: the wind phase, the outflow phase, and the inflow phase (e.g., Mathewsand Loewenstein, 1986; D’Ercole et al., 1989; Ciotti et al., 1991; Renzini, 1996).If the galaxies we observe are in the inflow phase (for the expected temperaturegradient, see Trinchieri et al., 1994; Kim and Fabbiano, 1995), the cooling timeof the hot coronal gas is shorter than a Hubble time over large parts of a galaxy’soptical image suggesting that the gas near the core has approximately settled into asteady state configuration (but see Kritsuk, 1992). X-ray luminosities of individualearly-type galaxies then indicate cooling rates of order 1M� yr�1, similar to therate of mass ejection from evolving stars (for the possible presence of cool X-rayabsorbing gas, see Fabbiano, 1992).

If this standard cooling scenario is valid, correlations may be found betweenLX ,LB , MHI, and MH2 (cf. Sarazin, 1990; Wiklind 1991; Eskridge et al., 1995). Thecorrelation between LX and LB can be explained by heat input from supernovaeand from collisions between particles of the interstellar gas and the outflowingmaterial from stars. Gains in cool gas mass should be compensated by lossesdue to star formation in the flow. LX and MH2 are then coupled, because LXcorrelates with the infalling and cooling gas mass, while MH2 is a measure of thestar formation rate. MHI and MH2 should then be related via the relative lifetimes


Table VHot and cool gas properties of individual galaxiesa

Source Db log _M fX SFRg SFRh SFRi

LcB Ld

FIR LX M eX

Mpc L� M� M� yr�1

NGC 315 65 11.1 9.6 8.0 10.1 2.25 2.71NGC 524 30 10.7 9.5 7.7 9.5 0.61 2.16NGC 1316 20 11.1 9.7 7.0 9.6 0.84 3.56 0.63NGC 1332 18 10.2 9.0 6.7 8.4 0.05 0.66NGC 1380 22 10.5 9.5 6.8 8.8 0.13 2.00NGC 1395 20 10.6 8.9 6.5 9.0 0.18 0.51NGC 1533 7 9.2 8.0 5.1 6.4 0.00 0.07NGC 1600 60 11.0 9.1 7.8 9.9 1.65 0.80NGC 2685 18 10.0 8.9 6.0 7.8 0.01 0.68 0.16 2.07NGC 2974 30 10.4 9.4 6.8 8.7 0.10 1.65NGC 3258 40 10.5 9.2 7.4 9.1 0.27 1.11NGC 3593 7 9.3 9.5 5.7 6.8 0.00 2.18 1.00 0.70NGC 4374 18 10.7 9.1 6.9 9.1 0.26 0.76NGC 4406 18 10.8 8.6 7.9 9.7 0.90 0.26NGC 4459 18 10.2 9.5 6.3 8.2 0.03 2.06NGC 4477 18 10.2 8.9 6.1 8.1 0.03 0.49NGC 4552 18 10.5 8.8 6.8 8.8 0.11 0.42NGC 4649 18 10.8 9.1 7.5 9.5 0.63 0.82 0.08NGC 4697 18 10.6 9.2 6.6 8.8 0.13 1.02NGC 4753 18 10.5 9.7 6.4 8.6 0.08 3.25NGC 5128 4 10.5 10.0 6.0 8. 4 0.05 6.72 0.40 0.60I 4296 45 10.9 9.0 7.6 9.7 1.12 0.65NGC 5866 13 10.1 9.7 6.0 7.9 0.02 3.04 0.80 1.40NGC 5898 30 10.4 9.0 6.6 8.6 0.08 0.70

a The sample of 24 S0a-E galaxies with classification and LX from Forman et al. (1985) that are detected byIRAS (Knapp et al., 1989). Two Sa galaxies from the Forman et al., sample, NGC 1316 and NGC 3593, areincluded (see the Hubble types given in Table I). NGC 6876 could not be included because the IRAS beamencompasses yet another nearby early-type galaxy (Knapp et al., 1989). CO, allowing estimates of MH2 , isonly observed towards six of the objects (see Tables I and II), while CO and HI were observed in four sources.For related data, see Table 3 of Braine et al. (1997).b Distances were determined in the same way as described in Section 3.c Obtained from the Bo

Tmagnitudes of the RC3 (de Vaucouleurs et al., 1991) accounting for galactic and

internal absorption and for redshift. Absolute blue solar magnitude: 5 :M47.d From Knapp et al. (1989) and IRAS (1989). The conversion of the fluxes into LFIR is done in the same wayas described by Wouterloot and Walmsley (1986).e From log(MX=M�) = �20:5 + 0:5 log LX [watt] + 1:2 log(LB=L�) (Roberts et al., 1991). MX valuesare upper limits in the case that a significant fraction of the X-ray emission is due to stellar sources.f From _MX [M� yr�1] = 2 � 10�10MX [L�] (Thronson et al., 1989; this corresponds to a cooling timescale of 5� 109 yrs for gas at Tkin= 107 K). _MX values are upper limits in the case that a significant fractionof the X-ray emission is due to stellar sources.g From SFR [M� yr�1]= 6:5�10�10 LFIR [L�] (Thronson and Telesco, 1986), which is valid for a SalpeterIMF with lower and upper cutoff masses of 0.1 and 100M�. If a significant fraction of the infrared emissionarises from diffuse gas not forming stars, the SFR is overestimated (Helou, 1986).h From SFR [M� yr�1] = 2 � 10�9 MH2 [M�]. The correlation is that for the galactic disk. For MH2 , seeTable II.i From SFR [M� yr�1] = 10�9 (MH2+MHI) [M�]. The correlation is that for the galactic disk. For MHIand MH2 , see Table II.


of neutral atomic and neutral molecular hydrogen. A short lifetime for dust grainsdue to sputtering in the strong X-ray radiation field (Draine and Salpeter, 1979)should minimize FIR luminosities.

In sharp contrast to this standard scenario, IRAS 100�m detections are numer-ous: Taking the sample of sources analysed by Forman et al. (1985), 37 (95%) outof a total of 39 S0a-E galaxies with X-ray detections have entries in the Knapp et al.(1989) IRAS data catalog. For those 24 (62%) that are detected in the far infrared,relevant parameters are displayed in Table V. These are the blue luminosity,LB , thetotal mass of the hot gas,MX , the hot gas inflow rate, _MX , the infrared luminosity,LFIR, and the star formation rate, SFR, as derived from the far infrared emission,the molecular, and the neutral gas mass.

_MX and SFRFIR values (with the SFR derived from the infrared luminosity) arepoorly correlated. FIR deduced star formation rates tend to be larger than the infallrate derived from X-ray data ( _MX ). This is a second (weaker) contradiction to thestandard picture.

Molecular gas has only been observed towards a few early-type galaxies withknown coronal X-ray emission. CO is not detected in X-ray luminous early-typegalaxies without significant far infrared emission (Braine et al., 1997). This isanother strong contradiction to the standard picture. CO profiles from an offsetposition 10 south of the center of the prominent Virgo elliptical NGC 4472 (Hucht-meier et al., 1988, 1994), showing an extended X-ray halo (e.g., Fabbiano, 1989),are not confirmed (Braine et al., 1997). For the six sources with CO detectionsin Table V, we can estimate the SFR dividing MH2 by 5 � 108 yr, the moleculargas consumption time of the galactic disk. For four sources, we can also divideMHI+MH2 by 109 yr, the total gas consumption time of the galactic disk (gas con-sumption timescales similar to those of the galactic disk are obtained with the LFIR

versus MH2 diagram for elliptical galaxies (Wiklind et al., 1995), assuming SFR� 1M� yr�1 for LFIR � 3� 109 L� (Section A.2) and MH2�MHI (Section 4.4)).The _MX/SFRCO ratios (with the SFR deduced from CO) tend to be smaller thanthose obtained using LFIR. Most _MX /SFRCO ratios are however still < 1, in spiteof the fact that _MX denotes the total deposited material which is not confined tothe CO beam or even to the limits of the optically visible galaxy. This is the fourthand weakest contradiction to the standard scenario, because it refers to a minorityof sources only.

We may follow several lines of reasoning to explain the large number of IRASdetections, the small _MX /SFR values, and the lack of correlations between thevarious observational parameters:

(1) The star formation rate is strongly varying with time. Inflowing gas isdeposited in an initially gas-poor central region. At this time, the star formation rateis well below the inflow rate. Once the amount of cool nuclear gas reaches a criticalconcentration (see, e.g., Kennicutt, 1989, for a threshold surface density in spirals),the star formation rate becomes larger than the inflow rate, thus depleting theavailable molecular fuel. A galaxy may undergo many such cycles. The weakness


of this scenario lies in the fact that there are too many IRAS galaxies among thenearby X-ray sources, and that ‘starburst ellipticals’ not undergoing interactionswith nearby galaxies are exceedingly rare (see Section 6.5).

(2) Effects of galactic rotation have not been incorporated in most models. Whilethere are no soft X-ray observations which indicate flattened distributions of thehot gas (e.g., Bregman et al., 1995; Kim and Fabbiano, 1995), it may be possiblethat galactic rotation, even if it is slow, is forcing the presumably almost dust freecooling material into a disk-like configuration (Kley and Mathews, 1995; Brighentiand Mathews, 1996). Gaseous or stellar disks observed in many ellipticals (e.g.,Rix and White, 1990) would then be formed during the cooling phase. The largenumber of IRAS detections and the correlation of CO with the far infrared, notwith the X-ray emission, does not favor such a scenario.

(3) The coronal gas is too metal poor to permit the detection of CO. This remainsan open question. Metallicities derived from ASCA are clearly subsolar (Awakiet al., 1994; Matsushita et al., 1994; Mushotzky et al., 1994; Arimoto, 1996a, b).Not being consistent with optical (stellar) data and theoretical expectations thatinclude the output of Type Ia supernovae, problems may be the limited accuracy ofatomic data related to the observed iron-L blends and the ‘low’ resolution (E=�E� 20 at 1.5 keV) of ASCA X-ray spectra (Arimoto, 1996a; Fabbiano, 1996b;Hwang et al., 1997).

(4) Star formation conditions in cooling flows may be fundamentally differentfrom those in the solar neighborhood, if the gas is almost devoid of dust (seeSection 6.4.2). In this case cool condensations would not provide the observed farinfrared emission and a comparison between cooling mass ( _MX ) and star formationrate (SFR, derived from LFIR) becomes meaningless.

With the large number of IRAS detections, SFR> _MX in most cases (Table V),and LFIR/MH2 values slightly larger than in the solar neighborhood, the infrared,mm-wave, and radio data imply that most of the observed cool dust and gas is notrelated to the hypothesized cooling flow.

There are also theoretical arguments which may imply that _MX is severelyoverestimated or that the flow is not the source of the neutral gas. If the galaxieswe observe are in the inflow phase and experience mass drop out, limits obtainedfrom optical and near infrared imaging require that up to 95% of the accreting gasis converted into low mass stars (mainly 0.5–1.0M�), which are primarily formedoutside the core (e.g., Schombert et al., 1993). Many individual early-type galaxieswith X-ray halos may however still be in the wind or outflow phase (Ciotti et al.,1991). In galaxies with nuclear activity even inflow may not necessarily lead todrastic cooling and high star formation rates. Instead, a centrally heated convectivecore of interstellar gas may form inside the volume occupied by the flow (Taborand Binney, 1993). Such a scenario does not require significant mass drop out anda highly reduced electric conductivity, but may be applicable to only a fraction ofstudied sources.


Information related to a systematic cooling flow mass drop out can also beobtained from the warm (� 104 K) gas (see Section 6.3). The presence of such gasis a characteristical property of dust lane ellipticals and a dust lane is detected invirtually all galaxies with warm gas (cf. Kim et al., 1988). Contrary to expectationsfrom cooling flow models, gas to dust mass ratios appear to be close to the galacticvalue (e.g., Sparks, 1992b). H� imaging of X-ray selected samples of galaxiesled Trinchieri and di Serego Alighieri (1991) and Macchetto and Sparks (1992) toconclude that LH�+[NII] may be correlated with LX , but in view of a large scatterit was also suggested that other parameters might play a role. Shields (1991) findsthat LH�+[NII] does not correlate with LX but with LFIR and thus with the coldinterstellar medium. Forbes (1991) concludes that dust and stars are decoupled(see also Section 6.5) and that cooling hot gas is not a significant source for dustin elliptical galaxies. According to Bregman et al., (1992), the cool gas is a diskphenomenon while the hot gas is a bulge phenomenon. There should be littleinteraction between the two components and the progression in galaxy type fromE to S0 is not only a sequence of decreasing stellar bulge to disk ratio but also asequence of decreasing hot-to-cold gas mass ratio. Although centrally peaked notrotationally supported HI may be present in the Virgo elliptical NGC 4406 (thisis expected in the standard cooling flow picture; see Bregman et al., 1988), noconvincing link between LX on the one hand and LH�+[NII], LHI, LCO, or LFIR onthe other hand has been established. In individual galaxies, cooling flow scenariosthus remain unconfirmed by data outside the X-ray window.

6.4.2. Cluster Cooling FlowsCluster cooling flows have recently been reviewed by Fabian et al. (1991) andFabian (1994), mainly focusing on the hot gas component. Here, we address suchcooling flow scenarios from an opposite perspective, emphasizing the cool neutralgas phase. Cluster X-ray halos have metallicities which are, within a factor of two,solar. The distribution of the hot gas mostly asssociated with clusters containing acentral dominant galaxy often shows a gradient with the coolest and densest gasbeing located close to the central dominating giant elliptical or cD galaxy. X-raydata suggest typical inflow rates of _MX = 50�100M� yr�1 and maximum valuesof up to 1000M� yr�1 (e.g., Johnstone et al., 1992; Allen et al., 1993, 1996; Edgeet al., 1994; Fabian and Crawford, 1995). Even a casual comparison with a normalspiral galaxy like the Milky Way having a star formation rate of a few M� yr�1

demonstrates that, if there is mass drop out, the initial mass function must be steeperthan in the galactic disk.

One way to search for the cooled down mass is to look for spatially extended lowtemperature gas. The X-ray luminositiy profile is often less peaked than it would beif all the gas were to flow into the center. The integral mass deposition rate appearsto be consistent with _MX(r) / r, so that the intracluster medium may consist ofseveral phases characterized by temperatures between 108 and 10�100 K, allowingthe formation of atomic and molecular neutral clouds. Radial profiles of temperature


Figure 25. Surface brightness and luminosities (assuming an emitting area of r = 50 kpc) for anabsorbing gas component withNH = 1021 cm�2, solar metallicity, and coronal ionization equilibriumat a pressure of 3� 105 K cm�3 (reproduced from Voit and Donahue, 1995).

and density inferred from X-ray imaging and spectroscopic observations show thatcooler (106 � 107 K) gas indeed exists. White et al. (1991) reported the discoveryof cool X-ray absorbing gas with substantial column densities (N � 3 1020 �1021 cm�2), masses (M � 1011 � 1012 M�), and covering factors (cf � 1) thatis distributed over the inner few 100 kpc and that appears to correlate with themagnitude of the cooling flow (see also Hu, 1992; Johnstone et al., 1992; Allenet al., 1993, 1995; Wang and Stocke, 1993; Nesci et al., 1995; Ragaranjan et al.,1995). The X-ray absorbing gas masses are comparable to the amount condensingout of a 100M� yr�1 flow over a Hubble time, providing a compelling argumentin favor of a direct connection to the cooling flow.

Clear evidence for gas cooling below 3 � 106 K would be an important stepin mapping the distribution of _MX(r); at these temperatures CNO nuclei areable to retain one or more electrons so that radiative cooling becomes efficient.O VIII luminosities in excess of 1043 erg s�1 that should arise from an absorbing


medium with Tkin � 106 K (see Figure 25) have not been seen in ASCA spectra(Fabian et al., 1994a; Voit and Donahue, 1995). Optical emission from Fe x, atracer of ionized gas at Tkin � 106 cm�3, may have been seen (Donahue andStocke, 1994). Patches of line emitting nebulosity are found in the inner few kpcof many cluster dominant galaxies (Heckman, 1981; Cowie et al., 1983; Hu et al.,1985; Johnstone et al., 1987; Heckman et al., 1989a; Sparks et al., 1989; Shieldsand Fillipenko, 1990; Allen et al., 1992; Crawford and Fabian, 1992; Donahueet al., 1992; Sarazin et al., 1992). Could this emission be related to the X-rayabsorbing layer? Optical emission lines are not detectable in all cluster cooling flowgalaxies and the H� surface brightness distribution does not match the expectedcooling rate, if the commonly assumed _MX(r) / r correlation holds (cf. Baum,1992; Fabian, 1994). Ongoing cluster mergers, turbulent mixing layers, magneticuplifting, magnetic reconnection, magnetic collimation or stirring produced by ahighly variable intermittent central radio source may disrupt accumulated clouds inthe core of a flow and may release enough energy to transiently power the opticalemission line nebulosity around the central galaxy (cf. Begelman and Fabian,1990; Crawford and Fabian, 1992; Daines et al., 1994; Pistinner and Shaviv, 1995;Jafelice and Friaca, 1996; Norman and Meiksin, 1996; Zoabi et al., 1996). Theoverall efficiency for producing relativistic particles might be small so that thethermal energy of the coronal gas and the energy arising from the active nucleusof the cluster dominant galaxy may be similar (of order 1061 erg). H� surfacebrightness limits several 10 kpc off the centers of cluster cooling flows are oforder 5 � 10�18 erg s�1 cm�2 arc sec�2 (Voit and Donahue, 1995) which shouldbe compared with � 10�20 erg s�1 cm�2 arc sec�2 for quiescently cooling gas(Baum, 1992). The observational limits are obtained from sky subtracted data sothat a smooth H� surface brightness distribution would not be seen.

Recently, vibrationally excited H2, tracing warm (Tkin� 1000� 2000 K) molecu-lar gas, was detected in half a dozen of cluster dominant galaxies with cooling flows(Elston and Maloney, 1994; Falcke et al., 1997; Jaffe and Bremer, 1997). The innerfew kpc of these galaxies emit strongly in the H2 (1� 0) S (1) line. The strength ofthe H2 emission is difficult to explain with current photon ionization models andrequires H2 excitation by either slow shocks or suprathermal secondary electronsproduced by X-ray photoionization of cold clouds. It is too early to assess theimportance of this discovery on overall cooling flow models, firstly because H2

column densities are not known and secondly because the spatial relation of the H2

emission to the observed H� filaments and radio jets has still to be determined.Cool neutral gas towards the central cluster dominant galaxies is also detected:

Broad HI absorption features (�V1=2> 100 km s�1) were reported from NGC 1275(Perseus cluster; see Figure 26), M 87 (Virgo cluster), Cyg A, and PKS 2322�123(Abell 2597), while narrow absorption has been observed towards the HV compon-ent of NGC 1275 (see Figure 26 and Section 5.5), towards the peculiar ellipticalgalaxy at the center of the 2A 0335+096 cluster (tentatively), towards Hydra A(Abell 780), and towards NGC 5920 (MKW 3) (Crane et al., 1982; Jaffe et al.,


Figure 26. LV and HV system HI absorption profiles towards NGC 1275. Upper panel: VLA (solidline) and NRAO interferometer (dotted line) LV absorption profiles (Crane et al., 1982). Lower panel:NRAO 140 ft HV absorption feature (de Young et al., 1973).

1988; Jaffe, 1990, 1992; McNamara et al., 1990; Dwarakanath et al., 1994; O’Deaat al., 1994a; Conway and Blanco, 1995). HI emission was reported from M 87(Jaffe, 1992). Towards NGC 1275 and M 87, HI was mapped. Adopting �HI < 1and an excitation or ‘spin’ temperature Tspin = 100 K, we obtain HI masses with

MHI = 104:6 V 23 �V1=2 Se;HI M� (7)

for 21 cm line emission and

MHI = 105:8 V 23 �V1=2

Sa;HI

ScM� (8)


for absorption (V3: velocity in units of 1000 km s�1, assuming H0 = 75 km s�1

Mpc�1; �V1=2: full width to half power linewidth in km s�1; Se;HI, Sa;HI, Sc: HI

emission, HI absorption, and 21 cm continuum fluxes in Jy;: source solid angle inarcsec2). The HI mass is of order 1010 M� towards the nuclear continuum source ofNGC 1275. Masses towards M 87, both for the emission and the absorption lines,are about one order of magnitude smaller. HI mass limits deduced from (the lackof) emission are also of order 109 M� (McNamara et al., 1990). If �HI > 1, thesemasses are lower limits. Towards the mapped sources NGC 1275 and M 87, theHI distribution is patchy, the velocities are systemic, and the velocity dispersionis large (�V1=2 � 400 km s�1). The narrow HI absorption features are mostlyredshifted relative to the systemic velocity by a few 100 km s�1 (cf. Figure 27).While the flow velocity of the hot gas, V � 6 [ _MX /100M� yr�1] [T /108 K]�1=6

[tcool=1010 yr]�1=3 (Fabian, 1994), may be much smaller, the difference appearsto be intuitively consistent with fast infall of cooling material towards the centralsource. If ram pressure balances the weight of such infalling clouds, terminalvelocities of order 100 km s�1 may indeed be reached (Daines et al., 1994); suchmotions create internal stresses that may rapidly lead to cloud breakup (e.g., Nulsen,1986).

Searches for CO in cluster cooling flow galaxies (Jaffe, 1987; Bregman andHogg, 1988; Heckman et al., 1989; Lazareff et al., 1989; Mirabel et al., 1989;Grabelski and Ulmer, 1990; Braine and Wiklind, 1993; Reuter et al., 1993; Sofueand Wakamatsu, 1993; Antonucci and Barvainis, 1994; Barvainis and Antonucci,1994; Braine and Dupraz, 1994; McNamara and Jaffe, 1994; O’Dea et al., 1994b)only led to a detection in NGC 1275 (cf. Section 5.5). In view of the number ofsurveys and the rather ‘normal’ metallicities of the hot gas, this is a remarkablypoor result.

If the hypothesized flow contains a multitude of cooling (and cooled) clouds atall radii, the tight HI and CO limits severely constrain the physical parameters ofsuch clouds. With the ratioMH2=

_MX denoting the average lifetime of the moleculargas, sensitive upper MH2 limits from the central regions of eight cluster coolingflows yield maximum cloud lifetimes of a few 106 yr (Braine and Dupraz, 1994).This is very small compared with typical molecular gas consumption times of3� 107 yr in a starburst (MH2 � 1010 M�; SFR� 300M� yr�1) and 5� 108 yr inthe galactic disk (MH2 � 2:5� 109 M�; SFR� 5M� yr�1). Only for NGC 1275,the cloud lifetime is with 1:5 � 107 yr ( _MX = 300M� yr�1) similar to thatexpected in a strong starburst. The characteristic lifetime of the HI clouds, definedas MHI= _MX , is also small, namely 3� 107 yr for NGC 1275 and smaller for mostof the other cluster cooling flows. The values and upper limits cannot be reconciledwith collapse time scales that include the dissipation of angular momentum, thepresence of magnetic fields, and the limited efficiency of the conversion of gas intostars. Cool infalling clumps may be disrupted, mixing into the hot gas component,thus reducing its temperature and cooling time and triggering the formation ofnew cool condensations (Fabian, 1994). Such a process might further prolong


Figure 27. HI absorption profiles towards four early-type radio galaxies. Optical depth is plotted asa function of heliocentric radial velocity. Velocities derived from the optical redshifts are indicatedby arrows. Note that the HI absorption peak is redshifted w.r.t. the optically determined systemicvelocity in three cases (reproduced from van Gorkom et al., 1989).

timescales for star formation. In analogy to Section 6.4.1, we thus have reason tosuspect that (1) the observationally derived HI and molecular masses underestimatethe true amount of cold neutral gas, that (2) the star formation process is different,proceeding faster and resulting in an initial mass function steeper than in molecularclouds of the solar neighborhood, or that (3) the cool gas is not related to the coolingflow. In the following, these possibilities are analysed:

Constraints from HI. Derivations of HI absorption column densities and thusmasses assume full coverage of the continuum source. If this does not hold or ifthe gas is very clumpy, higher column densities may be possible. In addition, theHI gas may not be confined to lines-of-sight towards the continuum source only,where the observational sensitivity (Sa;HI=Se;HI = Tcont=Tspin) may be highest. HI

emission from a region off the nuclear continuum source may contaminate theobserved absorption feature so that interferometric maps are required.


With a covering fraction of randomly placed clouds,

cf = 1� e�NclAcl=Acont ; (9)

and an absorbing optical depth of

< � >= Ncl (1� e�cl)�Vcl

�Vtot

Acl

Acont= (1� e�cl)

�Vcl

�Vtot(�ln[1� cf ]) (10)

(Ncl: number of clouds; �Vcl;�Vtot: full width to half power (FWHP) line-widthsof an individual cloud and of the ensemble of clouds in front of the nuclear radio con-tinuum source; Acl; Acont: projected area of an individual cloud and projected areaof the nuclear continuum source; cf. Braine et al., 1995) we obtain for NGC 1275with cf;HI = 0:9 and NHI = 1:3 1021 (White et al., 1991), < �HI >= 0:0021(Jaffe, 1990), and �Vtot;HI = 400 km s�1 (Jaffe, 1992) an individual cloud FWHPlinewidth of �Vcl;HI � 0.36 km s�1. The very low limit of �Vcl;HI is valid if�cl;HI > 1. We have assumed that the cool gas discovered by White et al. (1991)is primarily HI. This is suggested by numerical calculations of the thermal andchemical equilibrium of dust free clouds (Daines et al., 1994). With

NHI = 1:8 � 1018Z �

Tmb;HI

K

�dV cm�2; (11)

the lack of observed HI emission towards most X-ray luminous clusters and thehigh column densities suggested by White et al. (1991) can only be reconciled, ifsmall scale clumping is very pronounced, leading to �cl;HI > 100 (note that forcf � 1, clouds only have to overlap along the line of sight, not necessarily invelocity). Correcting for saturation broadening, this yields an intrinsic cloud velo-city linewidth �Vcl;HI < 0.1 km s�1, including thermal and turbulent broadening.A moderately smaller covering factor, cf � 0:5, would not significantly alter thisresult. For thermal broadening alone, assuming Tkin= Tspin (a reasonable approx-imation in view of the slow radiative rates involved), we find FWHP linewidthsof

�Vcl;thermal;HI = (8 ln2 kTkin=mHI)1=2 �

�Tspin

22 K

�1=2

km s�1: (12)

Even if turbulent broadening were negligible in these clouds, a combination ofEquation (12) with the intrinsic �Vcl;HI value requires Tspin < 1 K, a temperaturethat is incompatible with excitation from the microwave background, radiativeheating from other sources, and collisional excitation. It is also incompatible withTspin > 20 K found throughout the Galaxy (e.g., Kulkarni and Heiles, 1988).Estimates for other clusters of galaxies lead to the same internal contradiction.Only a small covering fraction, cf � 0:1 or less, is consistent with the HI data.

The molecular gas. Characteristic column densities of 3�1020�1021 cm�2 (e.g.,White et al., 1991) are only barely sufficient to provide the degree of self-shielding


that is required to form large amounts of CO in the Galaxy (e.g., Frerking et al.,1982; Elmegreen, 1985; Duvert et al., 1986). A slightly smaller metallicity and alack of dust will increase required column densities further. If there exist manyclouds with individually smaller column densities than the total values suggestedby X-ray data, CO emission at detectable levels can be firmly excluded. On theother hand, if stars are forming at rates of approximately _MX and if the thresholdfor self-gravitation is indeed similar to the threshold for efficient CO shielding(Elmegreen, 1985), CO emission has to be seen: For _MX � 100M� yr�1 and a(rather short) molecular timescale of 108 yr, MH2 � 1010 M�. In spite of thesebasic problems and ignoring the suggestion by Daines et al. (1994) that most of thecool gas should be in the form of HI, we now discuss CO absorption in the sameway as before HI.

CO absorption against the nuclear continuum sources should be detectableirrespective of theNH2=ICO conversion factor and the kinetic temperature of the gas.With a CO excitation temperature that is small relative to the continuum backgroundradiation temperature (otherwise CO would be dissociated), line intensities aredetermined by the continuum flux, the line opacity, and the background sourcecovering factor. For predominantly molecular gas we find for NGC 1275 withEquation (10) cf;H2 � 0:9, �Vcl;H2 � 400 km s�1, and < �CO 1�0 ><0.0021(Braine et al., 1995)

�Vcl;CO < 0:36 (1� e��cl;CO 1�0)�1 km s�1: (13)

Because of

NH2 � 104NCO = 2:3 1018 Tex;CO (1� e�5:53=Tex;CO)�1 R �CO 1�0 dV cm�2

(14)

(Tex;CO in K), high opacities (�CO 1�0 > 1) are required if column densities arelarge enough to keep the gas molecular. The small CO electric dipole moment(see Section 3) then ensures that Tex;CO 1�0 � Tkin. The clouds have however tobe optically thin in the CO J = 1� 0 line, of order �cl;CO < 0.3, so that �Vcl;CO

(see Equation (13)) can approach those particularly small values observed in tinyquiescent galactic molecular clouds (e.g., Dickman, 1975). Optically thin CO thenrequires, for self-shielding, NHI > NH2 as long as metallicities are roughly solarand [CO]/[H2] abundance ratios are not far below galactic values. This leads backto the previously discussed (and rejected) scenario, where HI is dominating thecf � 1 cool gas component. Since gas phase chemical reactions involving themolecular ion H+

3 provide a rapid route for the formation of molecules like CO, itis difficult to see how these clouds could have most of their hydrogen in molecularform and yet have very small abundances in other molecules (like CO). Transientheating of grains by X-rays leads to CO desorption rates that are larger than thoseinferred by cosmic-ray heating of grains in galactic molecular clouds (Voit andDonahue, 1995); CO abundances vary little with cosmic-ray flux (Farquhar et al.,


1994). If the gas is dust free so that formation of molecules (including H2) proceedsslowly, we would end up again with a medium dominated by HI.

Sometimes it is claimed (e.g., Fabian, 1994) that the pressure of the hot clustergas is higher than that of the galactic ISM and that this may alter cloud propertiesand star formation conditions. Even the highest such values in the inner 10 kpc,P = nT � 106 � 107 K cm�3, are commonly encountered in molecular cloudsof the galactic center region (n (H2)� 104 � 105 cm�3; Tkin� 100 K; e.g., Gustenand Henkel 1983; Walmsley et al., 1986; Mauersberger et al., 1986), where theNH2/ICO conversion factor is smaller than in the galactic disk (Sodroski et al.,1995; Dahmen et al., 1996). The application of a ‘normal’ conversion factor wouldthus yield molecular masses which are too large, while the opposite is needed toexplain standard cooling flow scenarios.

As already mentioned, the lack of CO emission in all observed sources (exceptNGC 1275) excludes warm (Tex� 5 K; cf � 1) H2 clouds. Often it was assumed(e.g., Fabian et al., 1991, 1994c; Daines et al., 1994; Fabian, 1994; Ferlandet al., 1994) that molecular clumps should cool to temperatures near 2.7 K. WithTmb = (J�;Tex � J�;2:7)(1 � e�� ) (J : Planck function) the main beam brightnesstemperatureTmb goes to zero when the excitation temperatureTex approaches 2.7K.While absorption towards the nuclear continuum source should then be seen, suchabsorption is not observed. While this is difficult to reconcile with our argumentsconcerning column densities and linewidths, it is nevertheless useful to calculateequilibrium temperatures and to see whether very cool temperatures are, in prin-ciple, consistent with heating and cooling rates. Estimating the energy absorbedthrough X-ray heating (see O’Dea et al., 1994b; Braine et al., 1995), one obtainsan expression like

nH2� � 2:6 10�26�

Acl

(5 1017 cm)2

� �rcluster

100 kpc

��2 � Lx;tot

1045 erg s�1

� �nH2

104 cm�2

��fx

0:05

� �fh

0:5

� �fabs

0:5

�erg cm�2 s�1; (15)

where rcluster is the distance from the center of the flow, fx Lx;tot = LX fromr < rcluster; fh is the fraction of energy that goes into heating and fabs is thefraction of the incident flux absorbed (fabs � 0.5 for column densities in excess of3� 1022 cm�2 and Tx � 5� 107 K).

Neglecting any other heating mechanism, we then find with the cooling function� (cf. Goldsmith and Langer, 1978; O’Dea et al., 1994b)

nH2� � n2H2� =

7� 10�27 [nH2=103cm�3]0:4 T2:2[nH2=103cm�3]0:06

kin erg cm�3 s�1 (16)

an equilibrium temperature Tkin>�10 K (cf. Braine et al., 1995). Based on moresophisticated models, O’Dea et al. (1994b) find Tkin = 20 � 30 K, a result thatis almost independent of the metallicity of the gas; Voit and Donahue (1995)


calculate Tkin = 15 � 50 K, using cooling functions of Hollenbach and McKee(1979) and Goldsmith and Langer (1978) and including gas-grain interactions. Thecooling functions were obtained with the escape probablility approximation. Thisapproximation allows the direct determination of local cooling rates, because in thiscase only nearby parts of a cloud are radiatively interacting. Other approximations,e.g., the microturbulent approach, lead to similar results (e.g., Gerola and Glassgold,1978).

Ambipolar diffusion may heat cores with low column density and high spatialdensities (Braine et al., 1995). This requires that Larson’s relation n (H2)L �const. (L: cloud diameter; Larson, 1981) and the correlation between magneticfield B and H2 density observed in clouds of the galactic disk, B � B� n (H2)1=2,approximately hold. Genzel (1992) then gives an expression equivalent to

n2H2�amb � 6 10�35

�B�

�G

�4 � xc

10�7

��1 � NH2

5 1023cm�2

��2

n2H2

erg cm�3 s�1;

(17)

with xc denoting the ionization fraction of the gas.Cloud cores with large column densities (NH2

>� 1025 cm�2) might cool downto small (�2.7 K) temperatures because of a lack of X-ray and cosmic ray heating.However, the brightness temperature of an optically thick CO line is close to thegas temperature at which the line becomes opaque and will arise from the moreextended outer molecular layers that are also warmer. Photons will diffuse from thecloud’s warmer outer layers into the cool interior and gravitational contraction andprototstellar activity may further heat the central core. Thus even in this extremecase physical conditions still yield detectable CO emission. Heating to very hightemperatures (Tex � 100 K) by the nuclear nonthermal radio continuum (Maloneyet al., 1994) could reduce opacities in the observed lower CO rotational transitions.This effect, however, is obviouly not inhibiting the detection of CO in NGC 1275and should be relevant for the nuclear torus, not for a region of at least several kpcin size.

Neutral clouds with dust. So far, we have discussed clouds ignoring the (presum-ably absent) dust component. While dust can, in principle, cool a cloud much morerapidly than the gas, it is also an efficient absorber of stellar radiation. Dust emis-sion, if present, should thus have color temperatures of several 10 K (Voshchinnikovand Khersonskij, 1984; Dwek et al., 1990; Braine et al., 1995), mainly emitting inthe IRAS bands. The fact that warm dust is barely seen in � 20 sources (LFIR �1010 L�; e.g., Hu, 1988; Bregman et al., 1990; Grabelski and Ulmer, 1990; Knappand Patten, 1991; Lester et al., 1995; the non-detections at submm-wavelength byAnnis and Jewitt (1993) are just consistent with the IRAS data) reduces its opticalopacity to Av � 0 :m5 (see also Boyle et al., 1988; Romani and Maoz, 1992; Fer-guson, 1993, and references therein). Overall, the transfer of energy between dustand gas (if dust is present) will avoid a gas component with Tkin� 3 K. A variety


of further possibilities to explain the lack of CO emission can also be shown to beunconvincing (cf. Braine et al., 1995; Voit and Donahue, 1995): Photodissociationby UV photons is not likely in view of the weak UV radiation field of most ellipticalgalaxies. Very large CO linewidths (�Vtot � 1000 km s�1) can also be excluded.Drag from the intracluster medium is slowing down fast clouds to speeds similarto the turbulent velocities of the ambient gas (<� 1000 km s�1) within 108 yr. TheCO linewidths observed towards NGC 1275 as well as the HI data, sometimesobserved with bandwidths of several 1000 km s�1, do not provide a hint for verybroad features. Antonucci and Barvainis (1994), although providing less sensitivelimits than other groups, would have detected CO with �Vtot up to 2000 km s�1.

Summary. The contradictions between the standard model and the HI and COdata, the sensitive HI and CO limits obtained with filled-aperture and interferometricantennas, the small NH2/ICO conversion factor of the highly pressurized gas in thegalactic center region, the color temperatures of the dust far infrared emission, andthe high equilibrium temperatures of the cooling clumps all demonstrate that thereis no significant sink for very cool (� 3 K) gas and dust. We thus conclude that HI

and CO column densities are not underestimated.

The star formation process. If star formation is taking place at the rates inferredfrom X-ray observations and with a ‘normal’ IMF, the accreting galaxies should beremarkably blue even if _MX were substantially overestimated. Photometric workat optical and near infrared wavelengths have confirmed that galaxies with largeflows tend to exhibit bluer central colors than non-accreting cD galaxies; there mayexist a correlation between color and _MX (Thuan and Puschell, 1989; McNamaraand O’Connell, 1992). In general, the blue anomalies imply that only a few percentof the accreted gas may be forming stars if the IMF is ‘normal’ (cf. Wirth et al.,1983; Bertola et al., 1986; Crawford et al., 1989). Mg2 and D4000 indices also seemto be correlated with _MX . This supports the idea that cluster cooling flows haveongoing star formation at a level SFR� 0.1 _MX , if the IMF is normal (Cardiel et al.,1995). With SFR� 0.1 _MX , the extremely short molecular and neutral atomic gastime scale limits (MH2 / _MX andMHI/ _MX) would increase by an order of magnitudeand become somewhat more credible; however, the time scales remain significantlybelow those of the Galaxy and most of those ellipticals detected in CO. Furthermore,the optically traced ongoing star formation does not follow the mass depositionprofiles inferred from X-ray data: The star formation is more concentrated towardsthe central region of the galaxies. Furthermore, observational data and numericalresults rule out cooling flows as a mechanism for forming significant numbers ofglobular clusters in giant E and cD galaxies (Bridges et al., 1996).

At least five reasons have been mentioned why star formation conditions incooling flows may differ from those encountered in the Galaxy: (1) the gas pressureis high and the clouds may be extremely cool (Tkin � 3 K), (2) the clouds movingthrough a hot ambient medium may be small, (3) there may be a lack of dust grains,


and (4) relative velocities between the gas and stellar components may be large(>100 km s�1).

(1) Small kinetic temperatures and high pressure yield small Jeans masses (MJ �3 [T /10 K]2 [P=105 cm�3]�1=2 M�� 0.1M�) above which a cloud collapses underits own weight if magnetic forces and turbulence can be neglected. This has beenused as an argument in favor of an IMF heavily biased towards low mass stars(Fabian, 1994; Ferland et al., 1994). Neglecting magnetic forces and turbulenceis, however, not justified in view of observational results from galactic clouds (cf.Larson, 1981; Heiles et al., 1993); as mentioned above, Tkin � 3 K is inconsistentwith equilibrium calculations yielding Tkin >� 10 K.

(2) Dense gas clouds are disrupted by viscous drag when ‘falling’ throughthe hot medium. This may be an efficient way to keep the clouds small althoughmagnetic fields might prevent break up (cf. Hattori and Habe, 1990; Loewenstein,1990; Balbus, 1991; Yoshida et al., 1991).

(3) In the galactic disk, dust grains are formed in the circumstellar shells oflate-type stars and presumably also in supernova and nova ejecta (e.g., Seab et al.,1987). In the hot gas phase dust grains are efficiently destroyed at a rate tsputter � 107

(nH=10�3)�1 (a/10�2�m) yr (a: grain size; see Draine and Salpeter, 1979; Baum,1992) that is shorter than the cooling rate. Although infall of gas-rich companiongalaxies may inject fresh dust-rich material (cf. Hu, 1988, 1992; Sparks et al., 1989;McNamara and O’Connell, 1992), while large grains (a � 0.5 �m) may survive inthe outer regions (Voit and Donahue, 1995), the newly formed cold clouds are likelyalmost free of dust. Photodissociation by X-ray photons and long H2 formationtimescales caused by a lack of grain surfaces should yield a high ionization fractionof the molecular gas. This may couple ionized and neutral particles, thus inhibitingeffects related to ambipolar diffusion and reducing low mass star formation ratessubstantially.

(4) Large relative velocities between the inflowing gas and the mass losing starsof the central elliptical galaxy or its satellites might heat the gas, thus inhibitingstar formation (cf. Mathews, 1990; Sarazin, 1990).

The differences outlined above might indeed be significant enough to affect theinitial mass function. Denser clumps moving through the hot gas may be shreddedinto smaller condensations, leading perhaps to the formation of extended envelopeslike those reported from 1E 111.9�3754 and the cD galaxy in A 3284 (Maccagniet al., 1988; Molinari et al., 1994). While, however, an IMF biased in favor oflow mass stars would yield a better agreement between optically determined starformation rates and _MX values, the conflict with small MH2/ _MX and MHI/ _MX

lifetimes of the neutral gas components remains severe.

Cool gas and cooling flows: Are they related? It is remarkable that CO has onlybeen detected towards NGC 1275 but not towards any other source of its class. Whyis NGC 1275 so peculiar? The FIR and CO emission from NGC 1275 resemblesthat of gas-rich spirals undergoing a burst of star formation (cf. Section 5.5), while


other properties (apart from the peculiar HV system) are commonly observed inother cluster dominant galaxies as well (e.g., McNamara et al., 1996). Instead ofarising from a cooling flow, the gas and dust may have been accreted from one ormore gas rich galaxies; galaxy interaction may also be responsible for the centralswarm of star clusters. The non-detection of CO emission in other cluster dominantgalaxies with cooling flows then indicates that such events are rare. It is interestingto note that it is much easier to explain the CO emission from NGC 1275 withoutinvoking the presence of a cooling flow; otherwise, we face the problem to explainthe lack of detectable CO emission in other cluster dominant galaxies.

To summarize. A direct link between HI, CO, and X-ray data in cluster dominantgalaxies has not been found. From an observational point of view, the problemsare enormous. High angular resolution data average over large linear scales (inNGC 1275, 100 corresponds to 300 pc). Small-scale density and temperature fluc-tuations cannot be observed. For the tiny distant cooling clumps in cluster coolingflows, there are few observational constraints which can guide a theoretical study.Essential parameters like mass spectrum, density and temperature distribution,intrinsic velocity dispersion, magnetic field strength, and kinematical propertieswill not be measured directly in the foreseeable future.

From X-ray measurements, consistent _MX values are derived by several appar-ently independent methods (cf. Arnaud, 1988; Canizares et al., 1988; Mushotzkyand Szymkowiak, 1988). Evidence for cooling flow material below 106 K ishowever totally missing. Contrary to many previous claims, there are severe prob-lems with the standard cooling flow scenario. Very broad lines or small excitationtemperatures (Tex� 3 K) appear not to be feasible means of reconciling largemolecular or atomic gas masses with the few detections and tight upper limitsobtained at radio wavelengths. It is difficult to reconcile the high covering fractionof the X-ray absorbing gas with the presence of a ubiquitous neutral gas component,observationally limited (if it exists at all) to cf � 0.1. The X-ray absorbing columndensities appear to be too small to build up molecular clouds massive enoughto explain the star formation rates obtained from spectroscopic studies at opticalwavelengths and the optically traced star formation does not follow the radial dis-tribution of the X-ray mass deposition rates. Assuming that the X-ray absorbinggas is hot (T > 104 K), evolutionary scenarios with a substantial contraction of thesurface area (by at least a factor of 10; cf. Norman and Meiksin, 1996; Allen andFabian, 1997) and significant radial infall of newly formed clumps prior to starformation only solve a part of the intrinsic inconsistencies of the model.

Searching for alternative scenarios for individual galaxies and cluster coolingflows, effects caused by the cluster’s gravitational potential, reheating of the cool-ing gas by conduction and magnetic reconnection, drag heating by galaxies andsupernovae, cosmic ray heating, and heating by hydrodynamic waves have beensuggested (e.g., Miller, 1986; Bregman and David 1988; Binney, 1988; Bohringerand Morfill, 1988; Meiksin, 1988, 1990; Pringle 1989; Rosner and Tucker, 1989;


Balbus and Soker, 1990; Mathews, 1990; Loewenstein et al., 1991; Murray andBalbus, 1992; Tabor and Binney, 1993; see Fabian et al., 1994b, for the modelof Sparks, 1992a; Norman and Meiksin, 1996). While the thermal energy of thecoronal gas is larger than the energy residing even in the most prominent radiolobes of cluster dominant galaxies, it may well be that relativistic particles onlymake up a small part of the nuclear energy ejected into interstellar space. Yet anoth-er scenario (Voit and Donahue, 1995) makes use of the fact that large grains cansurvive a Hubble time in a hot plasma (a > 1�m, nH < 10�4 cm�3). Chances forsurvival become better with lower gas density and larger grain size. This impliesimproved chances for survival at very large radii (r >100 kpc). Large grains arecharacterized by low optical extinction efficiencies per unit volume and fairly neut-ral extinction with wavelength; thus mass is locked up rather inconspicuously andmight be extended enough (r > 0:5 Mpc), that IRAS detections are inhibited.The mass column density for marginal optical thickness at optical and soft X-raywavelengths, � 10�4 g cm�2, appears however too large to fulfill the criteria ofhigh covering factor (cf � 1) and large extent (� 100 kpc).

The introduction of a cooling flow model is compelling, because it allows asystem to reach a state of lower energy. The standard scenario corresponds to ahighly ordered non-interacting system. Potential complexities, e.g., in the under-lying gravitational potential, are difficult to assess. The ‘deconvolution method’to derive _MX from the X-ray continuum is sensitive to the maximum radius ofsteady state cooling, while estimates based on spectroscopic data are sensitive tothe time interval the observed gas will emit a given line. It remains open, whether asort of energy input, leading to overestimates of _MX with both methods (cf. Sumi,1988), can reconcile radio and X-ray data. In view of the stringent observation-al limits and the not well understood internal heating sources, models requiringdrastically reduced amounts of cool deposited material are becoming more andmore attractive and present day cooling flows, if not so massive, will loose muchof their cosmological importance.

6.5. ON THE ORIGIN OF THE COOL DENSE GAS IN EARLY-TYPE GALAXIES

As we have seen (Section 6.4), there is little direct evidence for the formation of vastamounts of cool dense gas in cooling flows. What other processes may play a role?In order to discuss the origin of the cool ISM, it is important to obtain a vague idea ofthe ages of the various early-type objects. De Vaucouleurs r1=4 luminosity profilesare successfully reproduced by N-body simulations, following the dissipationlesscollapse and violent relaxation of an initially clumpy distribution of stars; tidallystripped galaxies can also develop a r1=4 distribution (e.g., van Albada 1982;Aguilar and White, 1986; Hjorth and Madsen, 1991). Dissipation may have playedan important role during the collapse of low luminosity early-type objects (e.g.,Kormendy, 1989; Burstein et al., 1992; Barnes, 1996). Observational evidencesupports the view that mergers and hyperbolic encounters produce galaxies with an


Figure 28. The HI distribution in the elliptical galaxy NGC 1052 (center) and in the spiral companionNGC 1042 (van Gorkom et al., 1986). The map shows that HI can be observed beyond the opticalradius of an early-type galaxy and that a gas-rich spiral is occasionally seen nearby.

r1=4 profile (see Section 5.6 for an example). While stellar disks in disky E’s may beprimordial (e.g., Scorza and Bender, 1995), careful observations of light profiles andkinematical properties often reveal shells, loops, ripples or twists (e.g., Capaccioli,1987; Quinn and Hernquist 1987; Balcells, 1992; Forbes and Thomson, 1992;Hernquist and Spergel, 1992; Nieto et al., 1992; Rampazzo and Sulentic, 1992;Schweitzer and Seitzer, 1992; Weil and Hernquist, 1993; Pen, 1994; Reduzzi et al.,1996) that may be of more recent origin. Spectroscopists frequently detect multipledynamical components and counterrotating cores (cf. Capaccioli and Longo, 1994;Bender, 1996) that may be caused by the presence of a triaxial system in projection(Statler, 1991), by mergers (Balcells and Quinn, 1990) or by gas processed at a lateevolutionary stage (Bender and Surma, 1992). Noting that star formation historiesof dwarf ellipticals are also complex and varied (e.g., Sancisi et al., 1987; Priceand Gullixson, 1989; Sandage and Hoffman, 1991; Sandage and Fomalont, 1993;


Skillman, 1996) and that mergers may form both disky and boxy Es (e.g., Bekkiand Shioya, 1997), this has led to the viewpoint that most early-type galaxies mighthave formed at late cosmological times as a result of interaction and merging events(cf. Balcells, 1992; Buson et al., 1992; Franx, 1992; Hau and Thomson, 1994).

The tight relationship between central velocity dispersion, �, effective radius,re, and effective surface brightness, Ie,

re / �� I�e ; (18)

(the fundamental plane; � � 1.4; � � 0.85) limits however the M=L dispersion(within re) to 12% (Renzini and Ciotti, 1993) and underlines a remarkable structuraland dynamic homogeneity that is difficult to reconcile with a large variation of ages.The small dispersion of colors and an analysis of color magnitude sequences inthe Coma and Virgo clusters (Bower et al., 1992) indicate that cluster galaxies areuniformly old and coeval. Being formed within a time interval that is short relativeto their present age, little star formation has occurred at later times. Observationsof more distant clusters show that, at least up to z = 1:2, early-type galaxies areamong the basic units; M=L ratios shift monotonically with redshift (e.g., Benderet al., 1996; Dickinson, 1996; Im et al., 1996; Yee et al., 1996; Kelson et al., 1997).While some cluster ellipticals must have formed at low redshifts, the majority mayhave formed at high z, evolving passively thereafter (but see Kauffmann et al.,1996; de Jong and Davies, 1997). Field ellipticals tend to show a stellar populationcomponent of intermediate age (e.g., Bica and Alloin, 1987; Bower et al., 1990;de Carvalho and Djorgovski, 1992; Allington-Smith et al., 1993; Rose et al., 1994)but are also found at high redshifts (z � 0.7) with intrinsic colors similar to thoseof local ellipticals (cf. Bender, 1996; Koo, 1996).

How does the neutral ISM fit into such a population of early-type galaxiespresumably dominated by old coeval objects? As already mentioned, cooling flows(Section 6.4) are likely not the dominant source for cool dense gas. To give just oneargument: Optical emission line regions in X-ray luminous galaxies are associatedwith substantial dust absorption (e.g., Macchetto and Sparks, 1991; Goudfrooijet al., 1994b) that is difficult to explain in terms of a ‘cooling flow’ scenariowith short dust sputtering time scales. Explaining the existence of a cool gascomponent by mass loss from a stellar population is also problematic, becausethe specific angular momentum of the gas surpasses that of the stars (e.g., Knapp,1987; Schweizer, 1987; Tadhunter et al., 1989). Rings or disks of gas and dust,characterized by a high specific angular momentum and often being warped ormisaligned w.r.t. the stellar body (see Figure 28), may however be explained interms of accretion from nearby galaxies. In such a scenario the cool and warm(Tkin� 104 K) gas will eventually be evaporated by the ambient hot X-ray plasmain an ‘evaporation flow’ (Sparks et al., 1989; de Jong et al., 1990). Additionalsupport for the accretion of gas and dust can be obtained studying MHI/LB andMH2/LB ratios in a statistical way (Knapp et al., 1985b; Wardle and Knapp, 1986;Forbes, 1991; Lees et al., 1991; Wiklind et al., 1995): S0a galaxies are similar to


Figure 29. Schematic sketch of a Cen A-like galaxy. The figure rotation axis is the horizontal linealong the galaxy’s short axis in the plane of the sky with the upper parts receding and the lower partsapproaching the observer. The three tilted straight lines on each side of the major axis indicate theanomalous orbits which simulate a warped disk (from van Albada et al., 1982).

spirals in that the blue luminosity, HI, and H2 content are correlated, with welldefined mean values and a relatively small dispersion.MHI/LB and MH2/LB ratiosfrom elliptical galaxies are instead better described by a power-law differentialdistribution encompassing a larger range of values (see Section 4.4 and Figure 6).The gaseous and stellar contents thus appear to be decoupled, consistent with thehypothesis that the gas is of external origin. While this explains the strong variationsin measured star formation rates with time (Sections 4.3, 5.6, 5.7), strongest supportfor this scenario is provided by detailed studies of kinematical and morphologicalproperties.

In order to become familiar with expected morphologies, it is necessary toconsider the set of simple closed orbits. In the case of S0 galaxies, the accreted gasfrom outside either settles into the equatorial plane or into a plane perpendicular toit giving rise to a polar ring (Steiman-Cameron and Durisen, 1982; Sparke, 1986).The angular momenta of gas and stars become parallel or antiparallel in the first andorthogonal in the second case. In (axisymmetric) oblate galaxies the gas will circlethe short axis and the preferred plane lies along the major axis; in (axisymmetric)prolate or cigar-shaped galaxies stable closed orbits are found along the minor axis(Tohline et al., 1982). Orbits can be nearly circular. Statistically investigating theintrinsic shape of elliptical systems, poor fits are obtained with axisymmetricalpopulations (Fasano and Vio, 1991; Tremblay and Merritt, 1995; Ryden, 1996),


thus providing compelling evidence for the presence of triaxial structures. Wherefigure rotation is negligible, stable almost elliptical gas flow orbits exist in theplanes perpendicular to the long and short axes (cf. Heiligman and Schwarzschild,1979; Schwarzschild, 1979). Including slow (‘tumbling’) figure rotation about oneof the principle axes, an additional family of stable orbits exists that is intrinsicallywarped. The tilt of the gas layer w.r.t. the principle plane perpendicular to the figurerotation axis increases with decreasing distance from the center (Merritt and deZeeuw, 1983, and references therein) until the orbits define a plane containing theaxis of figure rotation (see Figure 29). In a triaxial galaxy, the projected principalaxes do not generally coincide with the major and minor axes of the projectedsurface brightness distribution; there is a significant misalignment between thekinematical axes of the gas and the photometric axes of the optical light distribution.Hence warped or kinematically decoupled gaseous rings and disks with ordered,not necessarily circular motion and high specific angular momentum are a directsignature for accretion of material from other galaxies.

Mainly as a consequence of sensitive spectroscopic surveys at optical wavelengths,an increasing number of such kinematically decoupled interstellar components havebeen found in recent years (e.g., Whitemore et al., 1990; Bertola et al., 1992a; vanDokkum and Franx, 1995). In several cases however not only the warm ionized gasbut also the cool neutral interstellar medium is known to show a peculiar velocityfield. Table VI displays prominent sources with cool gas that can be characterizedby a single parameter, the inclination (i 6= 0�) relative to the main optical axis.While the number of sources is small, a high fraction of those few objects that havebeen mapped in HI or CO belong to the sample (for NGC 3593 and NGC 2685,see Sections 5.1 and 5.3). According to van Driel and van Woerden (1991), len-ticular galaxies have mostly an inner (r < R25) or an outer (r > R25) HI ring(R25: isophotal major axis radius at the blue surface brightness of 25m arc sec�2).Approximately half of the inner rings are slightly tilted relative to the stellar disk,while half of the outer HI rings have inclinations of 25�–55� relative to the plane ofthe stellar disk. Such a behavior is expected for material just settling for equilibrium,since orbital clocks run slower at large radii.

Assuming that the probability for prograde and retrograde tidal encounters haveequal likelihood and that lifetimes for co- and counterrotating structures are thesame, Bertola et al. (1992a) estimate that in a large fraction (� 40%) of S0 galaxiesthe interstellar ionized gas must be of external origin. The fraction would furtherincrease if counterrotating disks have a shorter lifetime than corotating structures(Bettoni and Galletta 1991).

In the few elliptical galaxies studied with high angular resolution, HI featuresare mostly ‘irregular’, also in those cases (not included in Table VI) where theinclination parameter is not a suitable measure of the irregularity (cf. Raimondet al., 1981; Lake et al., 1987; Kim et al., 1988; Appleton et al., 1990). Somesources show structure that is reminiscent of tidal tales (cf. Sancisi et al., 1984;van Gorkom et al., 1986). Two (rare) examples of elliptical galaxies with ‘regular’


Table VIGalaxies with cool interstellar gas inclined relative to the galaxy’s optical axisa

Source Type Decoupled component �P:A:bkin Ref.c

H2 HI Ionized gas degrees

A 0136–0801 S0 + 90 27NGC 660 Sa? + + + 45 12, 17, 24NGC 855 dE + ? 10MCG–5–7–1 S0 + 90 27NGC 1052 E + + 90 4, 8NGC 1316 E + + 90 2, 21IC 2006 E + + 180 9NGC 1947 E + + 90 1, 5, 15, 21NGC 2685 S0 + + + 90 11, 14, 24NGC 2787 SB0 + + 25 11, 14NGC 3593d S0a + + 180 20, 26NGC 3941 S0 + 20 11NGC 3998 S0 + + 90 11, 14NGC 4138 S0 + + 180 28NGC 4203 S0 + 35 11NGC 4262 SB0 + 55 11NGC 4278 E + 45 3UGC 7576 + 90 27A 1230+09 E + ? 6NGC 4546e SB0a + + + 180 23NGC 4650 A S0 + 90 27NGC 4826 Sab + 180 16, 22II Zw 73 + 90 27NGC 5122 S0 + 90 27Cen A E + + + 90 1, 18NGC 5266 E + + + 90 7, 21, 29UGC 9562 + 90 27NGC 7013 S0a + 20 11NGC 7213 S0a + ? 13NGC 7252 E + + 180 19

a See Galletta et al. (1997) for a first polar ring CO survey.b Inclination angle between the gas component and the main stellar rotation axis.c References: (1) Bertola and Galletta (1978); (2) Schweizer (1980); (3) Raimond et al.(1981); (4) van Gorkom et al. (1986); (5) Mollenhoff (1982); (6) Lake et al. (1987); (7)Varnas et al. (1987); (8) Kim (1989); (9) Schweizer et al. (1989); (10) Walsh et al. (1990);(11) van Driel and van Woerden (1991); (12) Whitmore et al. (1990); (13) Wiklind (1991);(14) Bertola et al. (1992a); (15) Bertola et al. (1992b); (16) Braun et al. (1992); (17)Combes et al. (1992); (18) Quillen et al. (1992); (19) Wang et al. (1992b); (20) Wiklindand Henkel (1992a); (21) Sage and Galletta (1993); (22) Braun et al. (1994); (23) Sageand Galletta (1994); (24) Watson et al. (1994); (25) van Driel et al. (1995); (26) Bertolaet al. (1996); (27) Cox and Sparke (1996); (28) Jore et al. (1996); (29) Morganti et al.(1997)d The molecular and the ionized gas disks are counterrotating w.r.t. to the primary butcorotating w.r.t. the secondary disk (Section 5.1).e Because CO was detected only tentatively, this source is included in footnote (f) ofTable I.


HI rings or disks are NGC 807 and NGC 2974 (Dressel, 1987; Kim et al., 1988).The small number of sources in Table VI and applied selection criteria make anyquantitative conclusion preliminary. It is nevertheless interesting that there is alarge percentage of ellipticals with cool gas and dust distributed along the minoraxis. The prototype of this class is Cen A, where the presence of a dust lane iscrucial to disentangle the galaxy’s triaxial shape (e.g., Hui et al., 1995).

What are the processes leading to the formation of kinematically peculiar ringsand disks that are associated with prominent dust lanes? For early-type galaxieswith similar masses of rotating and counterrotating stars, the presence of an ISMmay be explained without invoking galaxy interaction, if stars moving on boxorbits have been scattered into tube orbits with clockwise and counterclockwisestreaming motions. This yields two counterstreaming galactic disks of similar mass,age, and scale length like those seen in NGC 4550 (cf. Evans and Collett, 1994;Sellwood and Merritt, 1994). In most cases, however, galaxy interactions musthave played a role as has already been exemplified in Sections 4.6 and 4.7. Not allinteractions are so severe that the entire system is disrupted (cf. Combes, 1992).Some hyperbolic encounters with large spirals and chance captures of HI-richsatellites or nearby dwarfs will affect accreting systems in a less dramatic way. Aninfall of MHI � 107 � 109 M� would leave the stellar component of the largergalaxy rather unperturbed. HI complexes at the lower end of this mass range withouta significant stellar population may also exist in the intergalactic medium (Brinks,1994). Initially, an S0 galaxy may form a disk from infalling material that is mostlyconverted into stars. If the angular momentum of the continuously infalling gasvaries slowly with time, a pre-existing disk will adjust its orientation to absorb thenew material. A rapid change in the angular momentum, however, will lead to theformation of a gaseous disk that is tilted relative to the main body of the galaxy.Any remaining gas in the original disk that collides with the new material will looseangular momentum and will be dumped into the central region (cf. Merrifield andKuijken, 1994). A recent analysis of core regions of 64 ellipticals suggests that theobserved dust lanes are not always completely relaxed (van Dokkum and Franx,1995). Moreover, stellar and dust lane rotation axes tend to be misaligned. Thisimplies that even in the nuclear regions of elliptical galaxies (r � 250 pc), dust andstars are dynamically decoupled; again an external origin of the ISM is likely.

How stable are the disks and rings and what is their ultimate fate? N-bodycalculations simulating the disruption of a dwarf galaxy and subsequent formationof a ring (Katz and Rix, 1992) suggest that lifetimes in excess of 109 yr are possible.Polar rings massive enough to be self-gravitating (Sparke, 1986; Dubinski andChristodoulou, 1994; Cox and Sparke, 1996) could even survive a Hubble time. Thisreduces the need for a large reservoir of dwarf galaxies. Among spiral galaxies thereare four known cases of retrograde motion (NGC 3626, NGC 4826, NGC 7217,and NGC 7331; Braun et al., 1992, 1994; Casoli and Gerin, 1993; Merrifieldand Kuijken, 1994; Ciri et al., 1995; Prada et al., 1996). The rare occurrenceof kinematically decoupled gaseous features in spirals can be explained by the


interaction with the pre-existing disk containing large amounts of gas and dust.Collisions between clouds may lead to loss of angular momentum, shock heating,and evaporation, eventually including an active phase of star formation.

While interactions between cosmologically old and approximately coeval galax-ies of various sizes explain most of the measured peculiarities (see e.g., Schweizerand Seitzer, 1992, for an analysis of optical data), we should not entirely ignorealternative processes which may not dominate the ISM of early-type galaxies butwhich may nevertheless affect observed gas masses, morphologies, and kinematics(e.g., Bettoni and Galletta, 1991; van Driel and van Woerden, 1991). Such processesare:

(1) Continuous infall of primordial intergalactic matter: the presence of metalenriched material which is required for dust formation and the iron lines in thehot coronal gas of giant ellipticals is inconsistent with an ISM dominated bycontinuous infall of dark primordial matter. Counterrotating features are also notexplained (e.g., Bettoni and Galletta, 1991).

(2) Accretion from a gaseous halo: the details of the elemental compositionand the infall rate of halo gas are poorly known (see Section 6.4.1). Accretionof halo gas does not explain counterrotating features and the difference betweengas-rich S0 galaxies with radially confined HI rings and gas-rich spirals with moreuniformly filled gas disks.

(3) Stripping of gas through an external medium: gas stripping has long beenthe favorite process to explain the existence of gas poor S0 galaxies in the centerof clusters (e.g., Haynes et al., 1984). The lack of gas in isolated S0 galaxies ishowever not explained. Gas stripping would most efficiently affect the outer partsof a galaxy. While this is consistent with the concentration of molecular gas in thenuclear parts of S0 galaxies (Section 4.3), their outer HI rings as well as the largeradial extent of HI in some giant ellipticals (Figure 30; see also, e.g., Lake et al.,1987) are in conflict with this idea.

(4) Galactic winds: cool gas in the inner region of an early-type galaxy mightbe heated and removed by star formation, supernovae, or a nuclear galactic wind(e.g., White and Chevalier, 1983; Mathews, 1989). To form an outer HI ring withthe expelled cooled down gas requires, however, a significant input of angularmomentum. Warps, inclined and counterrotating rings as well as nuclear CO con-densations are not explained.

(5) Mass loss from giant stars: the ejection of circumstellar envelopes into theISM (cf. Knapp et al., 1992) leads to the formation of filled gaseous disks whichare found in S0a and Sa but not in S0 and E galaxies. The large radial extent ofthe HI gas in some ellipticals (e.g., Lake et al., 1987) and its high specific angularmomentum, surpassing that of the stars, are also not consistent with a cool ISMdominated by circumstellar material.


Figure 30. HI contours in steps of 3� 1019 cm�2 for IC 2006, superposed on a reproduction of twosummed IIa-J plates (from Schweizer et al., 1989).

6.6. ROTATION CURVES AND M=L RATIOS

The combined analysis of radio, near-infrared, optical, and X-ray data allows, inprinciple, to determine rotation curves over a vast range of galactocentric radiiand to trace the radial mass distribution also in regions not accessible to optic-al spectroscopy. The complexities intrinsic to the physics of hot coronae makeshowever mass determinations from the X-ray halos rather model dependent (e.g.,Fabbiano, 1989; Bertin, 1992; Capaccioli and Longo, 1994). Triaxial structures(see Section 6.5) can show a wide variety of stellar orbits and may be affectedby tangential unisotropies in the outer regions. The existence of dynamically cold,


rotationally supported, strongly flattened gaseous disks or rings is thus a greatadvantage (see Figure 30). These rings and disks are those structures commonlyobserved when studying the cool neutral gas phase. Kinematical mass estimatesare risky, potentially referring to gas that has not yet settled into an equilibriumconfiguration or that is on elliptical orbits (e.g., van Dokkum and Franx, 1995; seealso Section 6.5). Nevertheless such data provide important information on totalmass and M=L ratios.

While CO is tracing the innermost region of an (at least partially) opticallyobscured galaxy, HI is elucidating the kinematics of the outer parts (Figure 30).For circular motion, constant inclination, and a rotation curve with a slope lessor equal to that of a solid body, the maximum line of sight velocities w.r.t. thesystemic velocity always occur on the ‘kinematic line of nodes’. At these nodes, thegalactocentric radius is identical to its projected distance on the plane of the sky (cf.Rydbeck et al., 1993). Correcting for velocity dispersion and beam smearing, it isthus possible to derive rotational velocities as a function of galactocentric distancefrom CO and HI data. Using the potential-density pair presented by Hernquist(1990) to approximate the de Vaucouleurs r1=4 luminosity profile, the expectedcircular velocity as a function of radius becomes

Vrot =

pGM r

r + a; (19)

with a = 0:551re (cf. Wiklind et al., 1997).For the central region of Cen A (for the outer parts, including planetary nebulae

and globular clusters, see Capaccioli and Longo, 1994; Hui et al., 1995; Sparke,1996), a comparison between the H� and CO rotation curves (Bland et al., 1987;Nicholson et al., 1992; Rydbeck et al., 1993) shows significant deviations whichcan be attributed to differences in angular resolution and extinction (Figure 31). Themolecular rotation curve, which is free of extinction effects, rises rapidly, passes atr � 800 through the CO ring shown in Figure 19, turns at� 800 into a Keplerian falloff (this angular distance might be an upper limit only, caused by the finite angularresolution of the CO data), and rises again after reaching r � 2000. At r � 6000, therotation curve flattens withVrot � 250�300 km s�1. This flattening is also observedin HI (van Gorkom et al., 1990). Apparently, the CO rotation curve is modified bya mass of order 109 M� inside the central� 100 pc that is possibly associated withthe active nucleus of the radio galaxy. The high velocities might also be caused bystrong non-circular motions induced by the triaxiality in the gravitational potentialof the underlying galaxy or by an oval distortion of the gaseous disk component.In any case, lack of optical line emission from the molecular ring makes CO to themost reliable mass tracer of the central region. Overall, the rotation curve of thegas in Cen A resembles that seen in most disk galaxies (cf. Sofue, 1996). It shows asteep rise within the central� 100 pc, declines further out, rises again, and flattensat radial distances > 1 kpc. This clearly differs from the rotation curve expected


Figure 31. Cen A rotation curves: Maximum velocities of (1) the CO emission after ‘statistical imageanalysis’; (2) the intensity weighted mean velocity of the same profiles; and (3) the H� rotation curve(from Rydbeck et al., 1993).

in the case of a spherical potential with a light distribution following the r1=4 law(Equation (19)).

Modeling the double-horned CO profiles of NGC 759 (see Figure 7) andNGC 7252 (Dupraz et al., 1990), applying Equation (19), the fundamental planecorrelations, and assuming ring-like molecular distributions, the surface gas dens-ities, the inclination of the rings, and the molecular gas fractions can be derived(Wiklind et al., 1997). Molecular gas fractions then become smaller than thosederived applying the virial theorem, an important effect that may be relevant forthe entire class of ultraluminous mergers seen by IRAS.

For the central regions of the edge-on early-type disk galaxies NGC 3593 andNGC 4710, CO rotation curves have also been determined (Wiklind and Henkel,1992a; Wrobel and Kenney, 1992). In NGC 3593, optical and mm-wave data areconsistent with a steep rise of Vrot in the inner (< 1000) part of the galaxy and anapproximately constant Vrot � 125 km s�1 further out. A similarly steep rise to aprojected Vrot �150 km s�1 is observed in more detail towards NGC 4710, wherekinematically distinct features off the major axis are also present.

While CO is a rarely employed but appropriate tool to investigate the rotationalproperties of the gas in the innermost parts of an early-type galaxy, HI traces gas at


large galactocentric distances. If the gas is dynamically relaxed (which is difficultto prove), the kinematical masses can be used to determineM=L ratios. Old stellarpopulations have remarkably consistent mass-to-light ratios of order 1–10M�/L�,with M=Lv / L0:2

v for E galaxies (e.g., Kormendy 1987a, b; Bender et al., 1992,1993). Young stellar populations tend to have lower mass-to-light ratios, with thebulk of the luminosity arising from the most luminous (and massive) stars. ThusM=L ratios well in excess of 10 infer the presence of dark matter.

While dwarf ellipticals are often devoid of extended HI emission (e.g., Skillman1996; Young and Lo, 1997), Schweizer et al. (1989) and Franx et al. (1994) findan Mtot/LB enhancement from 5 to �15 M�/L� between the core and 6.5 re

(this corresponds to r �15 kpc at D = 20 Mpc) for the elliptical galaxy IC 2006(Figure 30). In the elliptical NGC 5666, the HI velocity field is regular enough toprovide a probe of the mass distribution out to almost twenty times the effectiveradius (Lake et al., 1987). The rotation curve remains flat outside r = 3 kpc(D = 35 Mpc). A preliminary analysis suggests that Mtot/LB , including the outerparts of the galaxy, increases by a factor of four relative to the central region. For thelenticular NGC 2685 (see Section 5.3) Schiminovich et al. (1995) find at r = 4R25a mass-to-light ratio of Mtot/LB � 30–35M�/L� that is five times larger than inthe central region.

Mtot/LB ratios from HI at large radii are typically >10M�/L� (Raimondet al., 1981; Knapp et al., 1985b; Krumm et al., 1985; Gottesman and Hawarden,1986; van Gorkom et al., 1986; Shostak, 1987; Kim et al., 1988; van Driel et al.,1988, 1989; van Driel and van Woerden, 1989; Schiminovich et al., 1995), i.e.,larger than the average ratios of 3–10M�/L� from optical studies (e.g., Faberand Gallagher, 1979; Schechter, 1980; van der Marel, 1991; Saglia et al., 1993;Bertin et al., 1994). Analyzing HI data, a statistical comparison between S0/Sa andlater disk galaxies was presented by van Driel and van Woerden (1991). InsideR25, the Mtot/LB ratios are similar to those in spirals (2–15M�/L�). OutsideR25, however, half of the early-type galaxies show drastically increased mass-to-light ratios. Combining optical data (typically extending to re and providingkinematical data with arcsec resolution) with HI data (extending to several re

with arcmin resolution), Bertola et al. (1993) find that the variation of Mtot/LBwith galactocentric radius is similar in elliptical and spiral galaxies; Mtot/LB isalmost constant inside re and increases further out (Figure 32). Differences inradial mass-to-light ratio distributions between E, S0, and spirals thus appear to besmall and matter potentially deposited by cooling flows seems not to have a severeeffect on radial mass distributions (e.g., Mathews, 1988). Dark halos must befairly axisymmetric: Not only the scatter in the Tully-Fisher relation (Franx and deZeeuw, 1992) and the azimuthal structure of stellar disks (Rix and Zaritsky, 1995),but also the expansion of the IC 2006 HI velocity field in harmonics (Franx et al.,1994) suggest that deviations from dark halo axisymmetry must be smaller thanE 0.6. To date, HI velocity fields in elliptical galaxies are mostly fitted assumingplanar circular motion or adopting tilted circular ring models just as it is customary


Figure 32. Cumulative M /LB as a function of radius for spirals (solid line) and data points fromelliptical galaxies. The dashed lines denote the uncertainties (Bertola et al., 1993).

for spirals. Triaxial models with elliptical orbits would provide more realistic fits.When combined with accurate two dimensional surface photometry and stellarabsorption line data, more comprehensive dynamical models could be constructed.A substantial increase in Mtot/LB with r also characterizes at least some of thenearby dE and dSph systems (Carter and Sadler, 1990; Pryor, 1992).

6.7. EARLY-TYPE GALAXIES AND MASSIVE STARS

The apparent absence of type II and Ib/c supernovae (SNe) in E/S0 galaxies and thedetection of only a handful of such objects in S0a/Sa galaxies (e.g., Barbon et al.,1989; van den Bergh and Tammann, 1991) was historically the main indication that


the stellar progenitors must be massive stars. Nowadays, with the well established20M� type II progenitor of SN 1987 a, this argument can be inverted: Because noSN II and Ib/c events are seen in E’s, we might conclude that massive star formationmust have ceased in this type of galaxy a long time ago.

Stellar evolution theory predicts that single stars withM <� 5–8M� end as whitedwarfs. According to the standard model, SN Ia events occur in binary systemscontaining such stars when the primary object is driven beyond the Chandrasekharlimit by mass transfer from the secondary (e.g., Branch, 1992). The type Ia rate perblue luminosity, � 0.15 ‘SNu’ units (1 SNu = 1 SN [LB /1010 L�] � [t/100 yr]),is almost independent of Hubble type. In spirals, SN(Ib+II)/SN(Ia) � 2 (van denBergh and Tammann, 1991). In ellipticals this ratio is a factor of at least 20 smaller(e.g., Cappellaro, 1996a, b). Among the almost 1000 SNe that were discovereduntil 1994, 45 confirmed SNe have been observed in elliptical galaxies; those 10with a secure classification are all of type Ia; another 11 have been classified astype I.

Analysing the global properties of the cool gas and dust in early-type galaxies(Sections 4.3 and 4.4), we found that MH2/LB ratios can be more than an order ofmagnitude smaller than in spirals, while the SFE (see Section 3.1) is approximatelythe same. Without having to invoke an IMF biased in favor of low mass stars, thiseffect reduces star formation and thus SN (Ib+II)/SN (Ia) ratios from 2 (as measuredin spirals) to <�0.2. Ellipticals without detected CO will often be characterized byparticularly small MH2 /LB values (cf. Lees et al., 1991; Wiklind et al., 1995) thatyield SN (Ib+II)/SN (Ia) < 0.1. Another effect potentially reducing the SN (Ib+II)/ SN (Ia) ratio should be investigated in more detail: We noted already, that COemission from early-type galaxies is confined to the nuclear region. While there isa lack of CO maps with arcsec resolution, H2 surface densities tend to be large.For NGC 759 (cf. Section 4.4) we obtain several 100M� pc�2, which is well inexcess of the �(HI + H2)� 10M� pc�2 (Prantzos and Aubert, 1995) near the solarsystem. The high-resolution map towards NGC 4710 (Wrobel and Kenney, 1992)yields �H2 = 800M� yr�1 at the peak of the CO distribution. Although some H II

regions are visible (e.g., Pogge and Eskridge, 1993), type II and Ib/c supernovae inearly-type galaxies might thus be more often obscured by dust than in the disks ofspiral galaxies.

Obviously, the present supernova rates are not providing significant limits toexclude the formation of massive stars. If we define all objects as ‘massive’ thathave been OB ZAMS stars in an early stage of their evolution, such stars arealready observed since more than four decades (e.g., Baade, 1951; Hodge 1963,1973; Wilcots et al., 1990; Lee et al., 1993): NGC 185 and NGC 205 formed asmall number of blue massive stars as recently as 1� 4� 107 yr ago, implying thatstars as massive as 7–15M� have not yet reached the red giant branch (for stellarlifetimes, see Schaller et al., 1992). More distant dwarf ellipticals with opticalspectra resembling those of galactic HII regions (cf. Section 6.3) may be similar.


Even if we ignore ‘proto-elliptical’ mergers with luminous starbursts, there arefurther arguments favoring the presence of massive stars. In Section 4.5, we reportedthe recent discovery of a cool very dense gas component (n (H2)> 104 cm�3) inseveral early-type galaxies. The quasithermal HCN and HCO+ emission shows noobvious peculiarity. Relative to the less dense interstellar medium traced by CO,the very dense component out of that stars may form is at least as massive as in thegalactic disk. The optically obscured dense nuclear gas component of Cen A (seeSection 5.4) shows all the properties that are elsewhere interpreted as the molecularsignature of a nuclear starburst, inferring the formation of massive stars. While inmost cases it is difficult to discriminate between UV radiation from horizontalbranch and from young massive stars (e.g., Binette et al., 1994; Lee, 1994), thedetection of blue condensations, UV imaging, and UV and optical spectroscopysupport the view that at least in NGC 1275, Hydra A, NGC 3199, NGC 5102,NGC 5173, and NGC 6166 (Wirth et al., 1983; Bertola et al., 1986; Shields andFilippenko, 1990; Ferguson et al., 1991; Vader and Vigroux, 1991; Smith et al.,1992; Nørgaard-Nielsen et al., 1993; Hansen et al., 1995; Deharveng et al., 1996;van Dyke Dixon et al., 1996) massive stars were recently formed.

7. Outlook

As we have seen, the latest generation of mm-wave telescopes and the successfulIRAS launch have greatly improved our understanding of the cool interstellarmedium of early-type galaxies. The absence of a correlation between the presenceof spiral arms or density waves and star formation efficiencies or gas consumptiontimes, the gradual variation of the cool gas properties from S0a to E galaxies, andthe detection of kinematically decoupled rings and disks have provided fascinatinginsights. First maps of molecular emission from early-type galaxies have also beenobtained.

Nevertheless, there remain more questions than answers: What is causing vari-ations in star formation activity over time intervals of �108 yr in lenticular anddwarf elliptical galaxies? Are these variations always triggered by galaxy inter-actions? Are the physical processes regulating star formation the same in spiral,lenticular, elliptical, and dwarf irregular galaxies? From a molecular perspective,compact ellipticals are totally unexplored. The situation is not much better for cDand dE galaxies: So far a single CO map has been published for each type. Is thepresence of a notable molecular component typical for dwarf ellipticals at the upperend of their luminosity range? Or are the two M 31 companions, NGC 185 andNGC 205, exceptional? Can the postulated smooth dust component (Goudfrooijand de Jong, 1995) be directly confirmed with the help of an appropriate gas phasetracer? And where is the associated gas? Are VLBA or VLBI measurements of theH2O megamasers in NGC 1052 and TXFS 2226–184 providing new insights into


the nuclear engine of early-type galaxies? And are there more H2O masers to befound?

The study of decoupled molecular gas in early-type galaxies is still in its infancy.What is the relative number of structures dominated by HI, H2, ionized gas andstars as a function of age? How many of the cool gas rings and disks are remnantsof a galaxy interaction, how long do they survive, and what is their ultimate fate?While many sources have been studied at optical wavelengths, high resolution COand HI data would be crucial to obtain information on the conversion of HI intoH2 gas and stars and to estimate the relative timescales for such processes. Whatis its detailed kinematical state, its excitation, chemistry, and star formation rate?Is it necessary to modify Mtot/LB ratios determined from HI accounting for thepresence of triaxial potentials and effects of figure rotation? Can HI and CO databe used to model the underlying axisymmetric or triaxial potential in a significantnumber of galaxies? Interesting insights into astrophysical phenomena might beprovided by the interface between the decoupled and the co-rotating ISM. Are largekinematical differences between the interstellar and stellar component inhibitingmassive star formation? Where is the cool gas predicted by the standard coolingflow model? Will vibrationally excited H2 provide new important insights? Andare ‘bimodal’ X-ray clusters, with distinct X-ray merging components causingenhanced gas densities and cooling rates, suitable targets for searches for cooldense gas?

Finding answers to these and other questions will require substantial efforts. Acrucial step ahead will be the construction of a sensitive array of mm-wave tele-scopes, preferrably at southern latitudes, where many of the prominent ellipticals(most notably Cen A) can be observed. This would allow to determine NH2=ICO

conversion factors in various early-type environments. Such an experiment is fun-damental since it is the basis for any molecular mass estimate, requiring detailedCO, C18O, and mm-wave continuum maps with arcsec resolution (cf. Mauersbergeret al., 1996a, b). The connection between the cool neutral gas, the warm (� 104 K)interstellar medium, and the diffuse X-ray emitting gas could be studied in muchmore detail if the spatial distribution of the cool gas were known for a larger sampleof sources. Questions related to the hypothesized ‘evaporation flows’ could alsobe attacked. Of highest interest are the nuclear regions: A thorough study of thecentral portions of sources like NGC 759 and Cen A would be important to elucid-ate nuclear properties and to compare the results with those obtained for spirals.Are there many ellipticals with a compact nuclear component of cool dense gasas is suggested by the large number of nuclear dust lane detections with the HST?Detailed observations of molecular gas in the central regions of giant ellipticalsmay be the best way to determine the mass distribution and, possibly, to study thekinematics of material approaching a hypothesized supermassive black hole.

Presently, the evolution of early-type galaxies is studied by observing richclusters at redshifts up to z � 1.2 (e.g., Baum et al., 1995a; Bender et al., 1996;Caldwell, 1996; Dickinson, 1996; Djorgovski, 1996; Im et al., 1996; Kauffmann


et al., 1996; Schade et al., 1996; van Dokkum and Franx, 1996; Kelson et al., 1997).Were elliptical galaxies not only bluer but also dustier in the distant past? Did theycontain significant amounts of molecular gas well after their initial star formingphase? And is there more cool neutral gas in the outer ‘halo ellipticals’ than in thecentrally located early-type objects (cf. Guzman et al., 1992)?

Searching for the cool dust component in the outer regions of large cooling flowswould be another important and technically challenging experiment, requiringsubmillimeter observations either avoiding the chopping technique or using widebeam throws. Details of massive star formation could be analysed by a combinationof molecular cloud observations and HST images and, less directly, by an accurateclassification of more than merely a handful of supernova events. The large extentof the HI emission in some early-type galaxies has led to the determination of highMtot/LB ratios at radii of several re, but data from more sources are desirable tofurther constrain models requiring the presence of dark matter.

To date, progress is most obviously shaped by developments in two separatefields: (1) Sensitive multichannel bolometers will soon provide information on thetotal flux and the spatial distribution of the (sub)millineter continuum emission, thusproviding direct estimates of the relative amounts of thermal dust emission and non-thermal radiation from relativistic electrons; column densities of the thermal dustcomponent will also be determined. (2) The successfully launched ISO Satelliteprovides data from the cool dust component (T < 30 K) that was not observedby IRAS. In addition, a wide range of spectroscopic measurements is being made,up to wavelengths of � = 200 �m. This, combined with the refurbished HST,the fourthcoming Parkes HI multibeam survey (e.g., Sadler, 1997) and the stilloperating ASCA and ROSAT satellites, comprises an exciting combination ofobservational opportunities that will further elucidate the physical and chemicalproperties of lenticulars, cD’s, elliptical, and dwarf elliptical galaxies.


Appendix

A.1. CO BEAM FILLING FACTORS

Since the largest filled-aperture mm-wave antennas have beam sizes not smallerthan 1000–6000, the CO beam filling factor is an important quantity. With opticallythick 12CO lines and cool molecular gas (an excitation temperatureTex= 10 K maythen be a good approximation), the beam filling factor f (12CO) is defined by

f(12CO) = [J�;Tex � J�;Tbg ]�1ZTmb;12CO 1�0 dv: (20)

J�;T is the Planck function, Tbg = 2:7 K, and Tmb is the main beam brightnesstemperature. Assuming that 13CO is optically thin (which is likely the case) andwith a rough estimate of the excitation temperature Tex (in the case that the gasis cool, Tex= 10 K is again a good approximation), we can compare the beamaveraged column density N1 with the column densityN2 along those lines-of-sightfrom where the bulk of the emission arises. The 13CO beam filling factor thenbecomes with �max;13CO 1�0 � –ln (1 – Tmb;13CO 1�0=Tmb;12CO 1�0)

f(13CO) =N1(

13CO)N2(13CO)

= [J�;Tex � J�;Tbg ]�1RTmb;13CO 1�0 dvR�13CO 1�0 dv

: (21)

A.2. STAR FORMATION RATES

Adopting a Salpeter IMF between 0.1 and 100M� and assuming that the molecularclouds out of which stars form are disrupted within 2 106 yrs, Thronson and Telesco(1986) deduce star formation rates of

SFR (FIR) � 6:5� 10�10 (LFIR=L�) M� yr�1 (22)

from the far-infrared luminosities. The corresponding equation between the SFRand the H� luminosity is

SFR(H�) � 2:5� 10�8 ��1 (LH�=L�) M� yr�1 (23)

(Gallagher and Hunter, 1986; � is related to the loss of ionizing photons caused bydust absorption and escape).

For determinations of the long-term SFR, the blue luminosity is a more appro-priate tracer. Gallagher et al. (1984) propose

SFR(blue) � 2:9� 10�11 (LB=L�) M� yr�1: (24)

A.3. ACRONYMS

ASCA: Advanced Satellite for Cosmology and Astrophysics (Tanaka et al., 1994)FIR: Far infrared


FR I: Fanarov and Riley (1974) class I radio galaxyFWHP: Full width to half powerHST: Hubble Space TelescopeHV: High velocity system in NGC 1275IMF: Initial mass functionIRAM: Institut de Radio Astronomie MillimetriqueIRAS: Infrared Astronomical Satellite (e.g., IRAS, 1989)ISM: Interstellar mediumISO: Infrared Space ObservatoryNRAO: National Radio Astronomy ObservatoryLV: Low velocity system in NGC 1275ROSAT: Rontgensatellit (Trumper, 1983; Pfeffermann et al., 1986)RC3: Galaxy catalog (De Vaucouleurs et al., 1991)RSA: Galaxy catalog (Sandage and Tammann, 1987)SFE: Star formation efficiencySFR: Star formation rateVLBA: Very Long Baseline Array (Kellermann and Thompson, 1985)VLBI: Very long baseline interferometryZAMS: Zero age Main Sequence

Acknowledgements

It is a pleasure to thank L. J. Sage, J. Braine, E. Churchwell, S. Dobereiner,P. Goudfrooij, P. Kalberla, J. J. Salzer, and W. Wild for helpful discussions orcomments in the course of preparing the manuscript.

References

Aaronson, M., Huchra, J., Mould, J., Schechter, P. L., and Tully, R. B.: 1982, Astrophys. J. 258, 64.Aguilar, L. and White, and S. D. M.: 1986, Astrophys. J. 307, 97.Allen, S. W., Edge, A. C., Fabian, A. C., Bohringer, H., Crawford, C. S., Ebeling, H., Johnstone, R.

M., Naylor, T., and Schwarz, R. A.: 1992, Monthly Notices Roy. Astron. Soc. 259, 67.Allen, S. W. and Fabian, A. C.: 1997, Monthly Notices Roy. Astron. Soc. 286, 583.Allen, S. W., Fabian, A. C., Edge, A. C., Bautz, M. W., Furuzawa, A., and Tawara, Y.: 1996,

Monthly Notices Roy. Astron. Soc. 283, 263.Allen, S. W., Fabian, A. C., Edge, A. C., Bohringer, H., and White, D. A.: 1995, Monthly Notices

Roy. Astron. Soc. 275, 741.Allen, S. W., Fabian, A. C., Johnstone, R. M., White, D. A., Daines, S. J., Edge, A. C., and Steward,

G. C.: 1993, Monthly Notices Roy. Astron. Soc. 262, 901.Allington-Smith, J. R., Ellis, R. S., Zirbel, E. L., and Oemler, A.: 1993, Astrophys. J. 404, 521.Annis, J. and Jewitt, D.: 1993, Monthly Notices Roy. Astron. Soc. 264, 593.Antonucci, R. and Barvainis, R.: 1994, Astron. J. 107, 448.Appleton, P. N., Pedlar, A., and Wilkinson, A.: 1990, Astrophys. J. 357, 426.Arimoto, N.: 1996a, in R. Bender and R. L. Davies (eds.), IAU Symp. 171, New Light on Galaxy

Evolution, p. 139.


Arimoto, N.: 1996b, in A. Buzzoni, A. Renzini, and A. Serrano (eds.), Fresh Views of EllipticalGalaxies, ASP Conf. Ser. 86, BookCrafters Inc., San Francisco, p. 239.

Arimoto, N., Sofue, Y., and Tsujimoto, T.: 1996, Publ. Astron. Soc. Japan 48, 275.Arnaud, K. A.: 1988, in A. C. Fabian (ed.), Cooling Flows in Clusters and Galaxies, Kluwer Academic

Publishers, Dordrecht, p. 31.Athanassoula, E. and Bosma, A.: 1985, Ann. Rev. Astron. Astrophys. 23, 147.Awaki, H., Mushotzky, R., Tsuru, T. et al.: 1994, Publ. Astron. Soc. Japan 46, L65.Baade, W.: 1951, Publ. Univ. Mich. Obs. 10, 7.Balbus, S. A.: 1991, Astrophys. J. 372, 25.Balbus, S. A. and Soker, N.: 1990, Astrophys. J. 357, 353.Balcells, M.: 1992, in G. Longo, M. Capaccioli, and G. Busarello (eds.), Morphological and Physical

Classification of Galaxies, Kluwer Academic Publishers, Dordrecht, p. 221.Balcells, M. and Quinn, P. J.: 1990, Astrophys. J. 361, 381.Balcells, M. and Stanford, S. A.: 1990, Astrophys. J. 362, 443.Balick, B., Faber, S. M., and Gallagher, J. S.: 1976, Astrophys. J. 209, 710.Balkowski, C., Alloin, D., and Le Denmat, G.: 1986, Astron. Astrophys. 167, 223.Bally, J. and Thronson, H. A.: 1989, Astron. J. 97, 69.Barbon,R., Capellaro, E., and Turatto, M.: 1989, Astron. Astrophys. Suppl. 81, 421.Barnes, J. E.: 1996, in R. Bender and R. L. Davies (eds.), IAU Symp. 171, New Light on Galaxy

Evolution, p. 191.Barthel, P. D., Schilizzi, R. T., Miley, G. K., Jagers, W. J., and Strom, R. G.: 1985, Astron. Astrophys.

148, 243.Barvainis, R. and Antonucci, R.: 1994, Astron. J. 107, 1291.Baum, S. A.: 1992, Publ. Astron. Soc. Pacific 104, 848.Baum, S. A., de Koff, S., Sparks, W., Miley, G., Biretta, J., Golombek, D., Macchetto, D., and

McCarthy, P.: 1995a, Science with the Hubble Telescope, Abstract booklet.Baum, S. A., Zirbel, E. L., and O’Dea, C. P.: 1995b, Astrophys. J. 451, 88.Becker, R., Henkel, C., Bomans, D. J., and Wilson, T. L.: 1995, Astron. Astrophys. 295, 302.Becker, R., Schilke, P., and Henkel, C.: 1989, Astron. Astrophys. 211, L19.Begelman, M. C. and Fabian, A. C.: 1990, Monthly Notices Roy. Astron. Soc. 244, 26.Bekki, K. and Shioya, Y.: 1997, Astrophys. J. 478, L17.Bell, M. B. and Seaquist, E. R.: 1988, Astrophys. J. 329, L17.Bender, R.: 1992, in G. Longo, M. Capaccioli, and G. Busarello (eds.), Morphological and Physical

Classification of Galaxies, Kluwer Academic Publishers, Dordrecht, p. 357.Bender, R.: 1996, in R. Bender and R. L. Davies (eds.), IAU Symp. 171, New Light on Galaxy

Evolution, p. 181.Bender, R. and Mollenhoff, C.: 1987, Astron. Astrophys. 177, 71.Bender, R. and Surma, P.: 1992, Astron. Astrophys. 258, 250.Bender, R., Burstein, D., and Faber, S. M.: 1992, Astrophys. J. 399, 462.Bender, R., Burstein, D., and Faber, S. M.: 1993, Astrophys. J. 411, 153.Bender, R., Capaccioli, M., Macchetto, F. et al.: 1992, in I. J. Danziger, W. W. Zeilinger, and K.

Kjar (eds.), ESO Conf. and Workshop Proc. 45, Structure, Dynamics and Chemical Evolution ofElliptical Galaxies, ESO, Garching, p. 1.

Bender, R., Dobereiner, S., and Mollenhoff, C.: 1988, Astron. Astrophys. Suppl. 74, 385.Bender, R., Surma, P., Dobereiner, S., Mollenhoff, C., and Madjesky, R.: 1989, Astron. Astrophys.

217, 35.Bender, R., Ziegler, B., and Bruzual, G.: 1996, Astrophys. J. 463, L51.Bertin, G.: 1992, in I. J. Danziger, W. W. Zeilinger, and K. Kjarin (eds.), ESO Conf. and Workshop

Proc. 45, Structure, Dynamics and Chemical Evolution of Elliptical Galaxies, ESO, Garching,p. 243.

Bertin, G., Bertola, F., Buson, L. M., Danziger, I. J., Dejonghe, H., Sadler, E. M., Saglia, R. P., deZeeuw, P. T., and Zeilinger, W. W.: 1994, Astron. Astrophys. 292, 381.

Bertola, F.: 1992, in G. Longo, M. Capaccioli, and G. Busarello (eds.), Morphological and PhysicalClassification of Galaxies, Kluwer Academic Publishers, Dordrecht, p. 115.

Bertola, F. and Galletta, G.: 1978, Astrophys. J. 226, L115.


Bertola, F., Buson, L. M., and Zeilinger, W. W.: 1992a, Astrophys. J. 401, L79.Bertola, F., Galletta, G., and Zeilinger, W. W.: 1992b, Astron. Astrophys. 254, 89.Bertola, F., Cinzano, P., Corsini, E. M., Pozzela, A., Persic, M., and Salucci, P.: 1996, Astrophys. J.

458, L67.Bertola, F., Gregg, M. D., Gunn, J. E., and Oemler, A.: 1986, Astrophys. J. 303, 624.Bertola, F., Pizzella, A., Persic, M., and Salucci, P. : 1993, Astrophys. J. 416, L45.Bettoni, D. and Galletta, G.: 1991, in B. Sundelius (ed.), Dynamics of Disk Galaxies, Goteborg,

p. 317.Bi, H. G., Arp, H., and Zimmermann, H. U.: 1994, Astron. Astrophys. 282, 386.Bica, E. and Alloin, D.: 1987, Astron. Astrophys. 181, 270.Bieging, J. H. and Biermann, P.: 1977, Astron. Astrophys. 60, 361.Binette, L., Magris, C. G., Stasinska, G., and Bruzual, A. G.: 1994, Astron. Astrophys. 292, 13.Binggeli, B., Sandage A., and Tammann, G. A.: 1988, Ann. Rev. Astron. Astrophys. 26, 509.Binney, I.: 1988, A. C. Fabian (ed.), Cooling Flows in Clusters and Galaxies, Kluwer Academic

Publishers, Dordrecht, p. 225.Birkinshaw, M. and Davies, R. L.: 1985, Astrophys. J. 291, 32.Bland, J., Taylor, K., and Atherton, P. D.: 1987, Monthly Notices Roy. Astron. Soc. 228, 5950Bohringer, H. and Morfill, G. R.: 1988, Astrophys. J. 330, 609.Borne, K. D. and Richstone, D.: 1991, Astrophys. J. 369, 111.Bothun, G. D. and Mould, J. R.: 1988, Astrophys. J. 324, 123.Bower, R. G., Ellis, R. S., Rose, J. A., and Sharples, R. M.: 1980, Astron. J. 99, 530.Bower, R. G., Lucey, J. R., and Ellis, R. S.: 1992, Monthly Notices Roy. Astron. Soc. 254, 601.Boyle, B. J., Fong, R., and Shanks, T.: 1988, Monthly Notices Roy. Astron. Soc. 231, 897.Braatz, J. A., Wilson, A. S., and Henkel, C.: 1994, Astrophys. J. 437, L99.Braatz, J. A., Claussen, M., Diamond, P., Wilson, A. S., and Henkel, C.: 1996a, Bull. Am. Astron. Soc.

28, 1392.Braatz, J. A., Wilson, A. S., and Henkel, C.: 1996b, Astrophys. J. Suppl. 106, 51.Braine, J. and Dupraz, C.: 1994, Astron. Astrophys. 283, 407.Braine, J. and Wiklind, T.: 1993, Astron. Astrophys. 267, L47.Braine, J., Henkel, C., and Wiklind, T.: 1997, Astron. Astrophys. 321, 765.Braine, J., Wyrowski, F, Radford, S. J. E., Henkel, C., and Lesch, H.: 1995, Astron. Astrophys. 293,

315.Branch, D.: 1992, Astrophys. J. 392, 35.Braun, R., Walterbos, R. A. M., and Kennicutt, R. C.: 1992, Nature 360, 442.Braun, R., Walterbos, R. A. M., Kennicutt, R. C., and Tacconi, L. J.: 1994, Astrophys. J. 420, 558.Bregman, J. N. and David, L. P.: 1988, in A. C. Fabian (ed.), Cooling Flows in Clusters and Galaxies,

Kluwer Academic Publishers, Dordrecht, p. 41.Bregman, J. N. and Hogg, D. E.: 1988, Astron. J. 96, 455.Bregman, J. N., Hogg, D. E., and Roberts, M. S.: 1992, Astrophys. J. 387, 484.Bregman, J. N., Hogg, D. E., and Roberts, M. S.: 1995, Astrophys. J. 441, 561.Bregman, J. N., McNamara, B. R., and O’Connell, R. W.: 1990, Astrophys. J. 351, 406.Bregman, J. N., Roberts, M. S., and Giovanelli, R.: 1988, Astrophys. J. 330, L93.Bridges, T. J., Carter, D., Harris, W. E., and Pritchet, C. J.: 1996, Monthly Notices Roy. Astron. Soc.

281, 1290.Brighenti, F. and Mathews, W. G.: 1996, Astrophys. J. 470, 747.Brinks, E.: 1994, in T. Montmerle, C. J. Lada, I. F. Mirabel, and J. Tran Thanh Van (eds.), The Cold

Universe, Editions Frontieres, Gif-sur-Yvette, p. 303.Bronfman, L., Cohen, R. S., Alvarez, H., May, J., and Thaddeus, P.: 1988, Astrophys. J. 324, 248.Brosch, N. and Loinger, F.: 1991, Astron. Astrophys. 249, 327.Brosch, N., Greenberg, J. M., and Grosbøl, P. J.: 1985, Astron. Astrophys. 143, 399.Burkert, A.: 1994, Monthly Notices Roy. Astron. Soc. 266, 877.Burstein, D., Bender, R., and Faber, S. M.: 1992, in I. J. Danziger, W. W. Zeilinger, and K. Kjar (eds.),

ESO Conf. and Workshop Proc. 45, Structure, Dynamics and Chemical Evolution of EllipticalGalaxies, ESO, Garching, p. 31.

Burstein, D., Bertola, F., Buson, L. M., Faber, S. M., and Lauer, T. R.: 1988, Astrophys. J. 328, 440.


Burstein, D., Davies, R. L., Dressler, A., Faber, S. M., Stone, R. P. S., Lynden-Bell, D., Terlevich,R. J., and Wegner, G.: 1987, Astrophys. J. Suppl. 64, 601.

Buson, L. M., Bertola, F., Burstein, D., and Renzini, A.: 1992, in I. J. Danziger, W. W. Zeilinger, andK. Kjar (eds.), ESO Conf. and Workshop Proc. 45, Structure, Dynamics and Chemical Evolutionof Elliptical Galaxies, ESO, Garching, p. 493.

Buson, L. M., Sadler, E. M., Zeilinger, W. W., Bertin, G., Bertola, F., Danziger, J., Dejonghe, H.,Saglia, R. P., and de Zeeuw, P. T.: 1993, Astron. Astrophys. 280, 409.

Byram, E. T., Chubb, T. A., and Friedman, H.: 1966, Science 152, 66.Byun, Y.-I., Grillmair, C. J., Faber, S. M., Ajhar, E. A., Dressler, A., Kormendy, J., Lauer, T. R.,

Richstone, D., and Tremaine, S.: 1996, Astron. J. 111, 1889.Caldwell, N.: 1984, Publ. Astron. Soc. Pacific 96, 287.Caldwell, N.: 1996, in A. Buzzoni, A. Renzini, and A. Serrano (eds.), ASP Conf. Ser. 86, Fresh Views

of Elliptical Galaxies, BookCrafters Inc., San Francisco, p. 93.Canizares, C. R., Fabbiano, G., and Trinchieri, G.: 1987, Astrophys. J. 312, 503.Canizares, C. R., Markert, T. H., and Donahue, M. E.: 1988, in A. C. Fabian (ed.), Cooling Flows in

Clusters and Galaxies, Kluwer Academic Publishers, Dordrecht, p. 63.Capaccioli, M.: 1987, in T. de Zeeuw (ed.), ‘Structure and Dynamics of Elliptical Galaxies’, IAU

Symp. 127, 47.Capaccioli, M. and Longo, G.: 1994, Astron. Astrophys. Rev. 5, 293.Cappellaro, E.: 1996a, in R. Bender and R. L. Davies (eds.), ‘New Light on Galaxy Evolution’, IAU

Symp. 171, 81.Cappellaro, E.: 1996b, in A. Buzzoni, A. Renzini, and A. Serrano (eds.), ASP Conf. Ser. 86, Fresh

Views of Elliptical Galaxies, BookCrafters Inc., San Francisco, p. 253.Cardiel, N., Gorgas, J., and Aragon-Salamanca, A.: 1995, Monthly Notices Roy. Astron. Soc. 277,

502.Carignan, C., Demers, S., and Cote, S.: 1991, Astrophys. J. 381, L13.Carilli, C. L. and Barthel, P. D.: 1996, Astron. Astrophys. Rev. 7, 1.Carlberg, R. G.: 1984, Astrophys. J. 286, 405.Carlstrom, J. E.: 1988, in R. E. Pudritz and M. Fich (eds.), Galactic and Extragalactic Star Formation,

Kluwer Academic Publishers, Dordrecht, p. 571.Carollo, C. M., Franx, M., Illingworth, G. D., and Forbes, D. A.: 1997, Astrophys. J. 481, 710.Carter, D. and Sadler, E. M.: 1990, Monthly Notices Roy. Astron. Soc. 245, 12.Casoli, F. and Gerin, M.: 1993, Astron. Astrophys. 279, L41.Chamaraux, P., Balkowski, C., and Fontanelli, P.: 1987, Astron. Astrophys. Suppl. 69, 263.Chin, Y.-N., Henkel, C., Millar, T. J., and Whiteoak, J. B.: 1996, Astron. Astrophys.312, L33.Chin, Y.-N., Henkel, C., Whiteoak, J. B., Millar, T. J., Hunt, M. R., and Lemme, C.: 1997, Astron.

Astrophys. 317, 548.Ciotti, L., D’Ercole, A., Pellegrini, S., and Renzini, A.: 1991, Astrophys. J. 376, 380.Ciri, R., Bettoni, D., and Galletta, G.: 1995, Nature 375, 661.Claussen, M. J. and Lo, K.-Y.: 1986, Astrophys. J. 308, 592.Colina, L. and de Juan, L.: 1995, Astrophys. J. 448, 548.Colina, L., Sparks, W. B., and Macchetto. F.: 1991, Astrophys. J. 370, 102.Combes, F.: 1992, in G. Longo, M. Capaccioli, and G. Busarello (eds.), Morphological and Physical

Classification of Galaxies, Kluwer Academic Publishers, Dordrecht, p. 265.Combes, F., Braine, J., Casoli, F., Gerin, M., and van Driel, W.: 1992, Astron. Astrophys. 259, L65.Combes, F. and Elmegreen, B. G.: 1993, Astron. Astrophys. 271, 391.Condon, J. J., Helou, G., Sanders, D. B., and Soifer, B. T.: 1990, Astrophys. J. Suppl. 73, 359.Conway, J. E.: 1997, in A. Pedlar (ed.), High Sensitivity Radio Astronomy, in press.Conway, J. E. and Blanco, P. R.: 1995, Astrophys. J. 449, L131.Cowie, L. L., Hu, E. M., Jenkins, E. B., and York, D. G.: 1983, Astrophys. J. 272, 29.Cox, A. L., Sparke, L. S.: 1996, in E. D. Skillman (ed.), <it ASP Conf. Ser. 106, The Minnesota

Lectures on Extragalactic Neutral Hydrogen, Astron. Soc. Pac., San Francisco, p. 169.Crane, P. C., van der Hulst, J. M., and Haschick, A. D.: 1982, in D. S. Heeschen and C. M. Wade

(eds.), ‘Extragalactic Radio Sources’, IAU Symp. 97, 307.Crawford, C. S. and Fabian, A. C.: 1992, Monthly Notices Roy. Astron. Soc. 259, 265.


Crawford, C. S., Arnaud, K. A., Fabian, A. C., and Johnstone, R. M.: 1989, Monthly Notices Roy.Astron. Soc. 236, 277.

Curtis, H. D.: 1918, Publ. Lick Obs. 13, 11.Da Costa, G. S.: 1992, in B. Barbuy and A. Renzini (eds.), ‘The Stellar Populations of Galaxies’, IAU

Symp. 149, 191.Dahmen, G., Huttemeister, S., Wilson, T. L., Mauersberger, R., and Linhart, A.: 1996, in R. Gredel

(ed.), ASP Conf. Ser. 102, The Galactic Center, BookCrafters Inc., San Francisco, p. 54.Daines, S. J., Fabian, A. C., and Thomas, P. A.: 1994, Monthly Notices Roy. Astron. Soc. 268, 1060.Davies, R. D. and Lewis, B. M.: 1973, Monthly Notices Roy. Astron. Soc. 165, 231.Davies, R. L., Efstathiou, G., Fall, S. M., Illingworth, G., and Schechter, P. L.: 1983, Astrophys. J.

266, 41.de Carvalho, R. R. and Djorgovski, S.: 1992, Astrophys. J. 389, L49.Deharveng, J.-M., Jedrzejewski, R., and Rocca-Volmerange, B.: 1996, in P. Benvenutti, F. D. Mac-

chetto, and E. D. Schreier (eds.), Science with the Hubble Space Telescope II, Space Tel. Sci.Inst., Baltimore, p. 230.

de Jong, R. S. and Davies, R. L.: 1997, Monthly Notices Roy. Astron. Soc. 285, L1.de Jong, T., Nørgaard-Nielsen, H. U., Jørgensen, H. E., and Hansen, L.: 1990, Astron. Astrophys. 232,

317.de Juan, L., Colina, L., and Golombek, D.: 1996, Astron. Astrophys. 305, 776.D’Ercole, A., Renzini, A., Ciotti, L., and Pellegrini, S.: 1989, Astrophys. J. 341, L9.de Vaucouleurs, G., de Vaucouleurs, A., Corwin, H. G., Buta, R. J., Paturel, G., and Fouque, P.: 1991,

Third Reference Catalogue of Bright Galaxies (RC3), Springer-Verlag, New York.de Young, D. S., Roberts, M. S., and Saslaw, W. C.: 1973, Astrophys. J. 185, 809.de Zeeuw, T. and Franx, M.: 1991, Ann. Rev. Astron. Astrophys. 29, 239.Dickinson, M.: 1996, in A. Buzzoni, A. Renzini, and A. Serrano (eds.), ASP Conf. Ser. 86, Fresh

Views of Elliptical Galaxies, BookCrafters Inc., San Francisco, p. 283.Dickman, R. L.: 1975, Astrophys. J. 202, 50.Djorgovski, S. and Santiago, B. X.: 1992, Astrophys. J. 391, L85Djorgovski, S. G., Pahre, M. A., and de Carvalho, R. R.: 1996, in A. Buzzoni, A. Renzini, and

A. Serrano (eds.), ASP Conf. Ser. 86, Fresh Views of Elliptical Galaxies, BookCrafters Inc., SanFrancisco, p. 129.

Dobereiner, S., Junkes, N., Wagner, S. J., Zinnecker, H., Fosbury, R., Fabbiano, G., and Schreier,E. J.: 1996, Astrophys. J. 470, L15.

D’Odorico, S., di Serego Alighieri, S., Pettini, M., Magain, P., Nissen, P. E., and Panagia, N.: 1989,Astron. Astrophys. 215, 21.

Donahue, M. and Stocke, J. T.: 1994, Astrophys. J. 422, 459.Donahue, M. and Stocke, J. T.: 1995, Astrophys. J. 449, 554.Donahue, M., Stocke, J. T., and Gioia, I. M.: 1992, Astrophys. J. 385, 49.Donnelli, R. H., Faber, S. M., and O’Connell, R. M.: 1990, Astrophys. J. 354, 52.Dos Santos, P. M. and Lepine, J. R. D.: 1979, Nature 278, 34.Draine, B. T. and Salpeter, E. E.: 1979, Astrophys. J. 231, 77.Dressel, L. L.: 1987, in T. de Zeeuw (ed.), ‘Structure and Dynamics of Elliptical Galaxies’, IAU

Symp. 127, 423.Dubinski, J. and Christodoulou, D. M.: 1994, Astrophys. J. 424, 615.Dupraz, C., Casoli, F., Combes, F., and Kazes, I.: 1990, Astron. Astrophys. 228, L5.Duvert, G., Cernicharo, J., and Baudry, A.: 1986, Astron. Astrophys. 164, 349.Dwarakanath, K. S., van Gorkom, J. H., and Owen, F. N.: 1994, Astrophys. J. 432, 469.Dwek, E., Rephaeli, Y., and Mather, J. C.: 1990, Astrophys. J. 350, 104.Ebneter, K. and Balick, B.: 1985, Astron. J. 90, 183.Ebneter, K., Djorgovski, S., and Davies, M.: 1988, Astron. J. 95, 422.Eckart, A., Cameron, M., Genzel, R., Jackson, J. M., Rothermel, H., and Stutzki, J.: 1990a,

Astrophys. J. 365, 522.Eckart, A., Cameron, M., Rothermel, H., Wild, W., Zinnecker, H., Rydbeck, G., Olberg, M., and

Wiklind, T.: 1990b, Astrophys. J. 363, 451.


Edge, A. C., Fabian, A. C., Allen, S. W., Crawford, C. S., White, D. A., Bohringer, H., and Voges,W.: 1994, Monthly Notices Roy. Astron. Soc. 270, L1.

Elmegreen, B. G.: 1985, in D. C. Black and M. S. Matthews (eds.), Protostars and Planets II,University of Arizona Press, Tucson, p. 33.

Elmegreen, D. M.: 1980, Astrophys. J. Suppl. 43, 37.Elston, R. and Maloney, P.: 1994, in I. S. McLean (ed.), Infrared Astronomy with Arrays, ApandSS

Library 190, Kluwer Academic Publishers, Dordrecht, p. 169.Eskridge, P. B. and Pogge, R. W.: 1991, Astron. J. 101, 2056.Eskridge, P. B., Fabbiano, G., and Kim, D.-W.: 1995, Astrophys. J. Suppl. 97, 141.Evans, N. W. and Collett, J. L.: 1994, Astrophys. J. 420, L67.Fabbiano, G.: 1989, Ann. Rev. Astron. Astrophys. 27, 87.Fabbiano, G.: 1992, in I. J. Danziger, W. W. Zeilinger, and K. Kjar (eds.), ESO Conf. and Workshop

Proc. 45, Structure, Dynamics and Chemical Evolution of Elliptical Galaxies, ESO, Garching,p. 617.

Fabbiano, G.: 1996a, in R. Bender and R. L. Davies (eds.), ‘New Light on Galaxy Evolution’, IAUSymp. 171, 123.

Fabbiano, G.: 1996b, in A. Buzzoni, A. Renzini, and A. Serrano (eds.), ASP Conf. Ser. 86, FreshViews of Elliptical Galaxies, BookCrafters Inc., San Francisco, p. 103.

Fabbiano, G., Gioa, I. M., and Trinchieri, G.: 1989, Astrophys. J. 347, 127.Fabbiano, G., Kim, D.-W., and Trichieri, G.: 1992, Astrophys. J. Suppl. 80, 531.Faber, S. M. and Gallagher, J. S.: 1976, Astrophys. J. 204, 365.Faber, S. M. and Gallagher, J. S.: 1979, Ann. Rev. Astron. Astrophys. 17, 135.Fabian, A. C.: 1994, Ann. Rev. Astron. Astrophys. 32, 277.Fabian, A. C. and Crawford, C. S.: 1995, Monthly Notices Roy. Astron. Soc. 274, L63.Fabian, A. C., Arnaud, K. A., Bautz, M. W., and Tawara, Y.: 1994a, Astrophys. J. 436, L63.Fabian, A. C., Canizares, C. R., and Bohringer, H.: 1994b, Astrophys. J. 425, 40.Fabian, A. C., Johnstone, R. M., and Daines, S. J.: 1994c, Monthly Notices Roy. Astron. Soc. 271,

737.Fabian, A. C., Nulsen, P. E. J., and Canizares, C. R.: 1991, Astron. Astrophys. Rev. 2, 191.Falcke, H., Gopal-Krishna, and Biermann, P.: 1995, Astron. Astrophys. 298, 395.Falcke, H., Rieke, M. J., Rieke, G. H., Simpson, C., and Wilson, A. S.: 1997, Astrophys. J., in press.Fanaroff, B. L. and Riley, J. M.: 1974, Monthly Notices Roy. Astron. Soc. 167, 31.Farquhar, P. R. A., Millar, T. J., and Herbst, E.: 1994, Monthly Notices Roy. Astron. Soc. 269, 641.Fasano, G. and Vio, R.: 1991, Monthly Notices Roy. Astron. Soc. 249, 629.Ferguson, H. C.: 1993, Monthly Notices Roy. Astron. Soc. 263, 343.Ferguson, H. C. and Binggeli, B.: 1994, Astron. Astrophys. Rev. 6, 67.Ferguson, H. C., Davidson, A. F., Kriss, G. A. et al.: 1991, Astrophys. J. 382, L69.Ferland, G. J., Fabian, A. C., and Johnstone, R. M.: 1994, Monthly Notices Roy. Astron. Soc. 266,

399.Ferrarese, L., Ford, H., and Jaffe, W.: 1995, Science with the Hubble Telescope, Abstract booklet.Ferruit, P. and Pecontal, E.: 1994, Astron. Astrophys. 288, 65.Fich, M. and Hodge, P.: 1993, Astrophys. J. 415, 75.Fischer, J., Geballe, T. R., Smith, H. A., Simon, M., and Storey, J. W. V.: 1987, Astrophys. J. 320,

667.Forbes, D. A.: 1991, Monthly Notices Roy. Astron. Soc. 249, 779.Forbes, D. A. and Thomson, R. C.: 1992, Monthly Notices Roy. Astron. Soc. 254, 723.Forbes, D. A., Franx, M., and Illingworth, G. D.: 1995, Astron. J. 109, 1988.Forman, W., Jones, C., and Tucker, W.: 1985, Astrophys. J. 293, 102.Franx, M.: 1992, in G. Longo, M. Capaccioli, and G. Busarello (eds.), Morphological and Physical

Classification of Galaxies, Kluwer Academic Publishers, Dordrecht, p. 23.Franx, M. and de Zeeuw, T.: 1992, Astrophys. J. 392, L47.Franx, M., van Gorkom, J. H., and de Zeeuw, T.: 1994, Astrophys. J. 436, 642.Freedman, W. L.: 1992, in B. Barbury and A. Renzini (eds.), ‘The Stellar Populations of Galaxies’,

IAU Symp. 149, 169.Frerking, M. A., Langer, W. D., and Wilson, R. W.: 1982, Astrophys. J. 262, 590.


Fritze-v. Alvensleben, U., and Gerhard, O. E.: 1994, Astron. Astrophys. 285, 775.Gallagher, J. S. and Hunter, D. A.: 1986, in C. J. Lonsdale Persson (ed.), Star Formation in Galaxies,

NASA CP-2466, p. 167.Gallagher, J. S., Hunter, D. A., and Tutukov, A. V.: 1984, Astrophys. J. 284, 544.Gallagher, J. S. and Wyse, R. F. G.: 1994, Publ. Astron. Soc. Austr. 106, 1225.Galletta, G., Sage, L. J., and Sparke, L. S.: 1997, Monthly Notices Roy. Astron. Soc. 284, 773.Garcıa-Barretto, J. A., Franco, J., Guichard, J., and Carrillo, R.: 1995, Astrophys. J. 451, 156.Gardner, F. F. and Whiteoak. J. B.: 1976a, Monthly Notices Roy. Astron. Soc. 175, 9.Gardner, F. F. and Whiteoak, J. B.: 1976b, Proc. ASA, 3, 63.Gardner, F. F. and Whiteoak, J. B.: 1979, Monthly Notices Roy. Astron. Soc. 189, p. 51.Garwood, R. W., Helou, G., and Dickey, J. M.: 1987, Astrophys. J. 322, 88.Genzel, R.: 1992, in W. B. Burton, B. G. Elmegreen, and R. Genzel (eds.), The Galactic Interstellar

Medium, Saas-Fee Advanced Course 21, Springer-Verlag, Berlin, p. 275.Gerola, H. and Glassgold, A. E.: 1978, Astrophys. J. Suppl. 37, 1.Goldsmith, P. F. and Langer, W. D.: 1978, Astrophys. J. 222, 881.Gordon, M. A.: 1990, Astrophys. J. 350, L29.Gordon, M. A.: 1991, Astrophys. J. 371, 563.Gottesman, S. T. and Hawarden, T. G.: 1986, Monthly Notices Roy. Astron. Soc. 219, 759.Goudfrooij, P. and de Jong, T.: 1995, Astron. Astrophys. 298, 784.Goudfrooij, P., de Jong, T., Hansen, L., and Nørgaard-Nielsen, H. U.: 1994a, Monthly Notices Roy.

Astron. Soc. 271, 833.Goudfrooij, P., Hansen, L., Jørgensen, H. E., and Nørgaard-Nielsen, H. U.: 1994b, Astron. Astrophys.

Suppl. 105, 341.Goudfrooij, P., Hansen, L., Jørgensen, H. E., Nørgaard-Nielsen, H. U., and van den Hoeck, L. B.:

1994c, Astron. Astrophys. Suppl. 104, 179.Grabelski, D. A. and Ulmer, M. P.: 1990, Astrophys. J. 355, 401.Graham, A., Lauer, T. R., Colless, M., and Postman, M.: 1996, Astrophys. J. 465, 534.Greenhill, L. J., Herrnstein, J. R., Moran, J. M., Menten, K. M., and Velusamy, T.: 1997, Astrophys. J.,

in press.Greggio, L. and Renzini, A.: 1990, Astrophys. J. 364, 35.Gregorini, L., Messina, A., and Vettolani, G.: 1989, Astron. Astrophys. Suppl. 80, 239.Grewing, M. and Mebold, U.: 1975, Astron. Astrophys. 42, 119.Gusten, R. and Henkel, C.: 1983, Astron. Astrophys. 125, 136.Guzman, R., Lucey, J.R., Carter, D., and Terlevich, R.J.: 1992, Monthly Notices Roy. Astron. Soc.

257, 187.Hagiwara, Y., Kohno, K., Kawabe, R., and Nakai, N.: 1997, Publ. Astron. Soc. Japan 49.Hansen, L., Jørgensen, H. E., and Nørgaard-Nielsen, H. U.: 1995, Astron. Astrophys. 297, 13.Haschick, A. D., Baan, W. A., Schneps, M. H., Reid, M. J., Moran, J. M. and Gusten, R.: 1990,

Astrophys. J. 356, 149.Hattori, M. and Habe, A.: 1990, Monthly Notices Roy. Astron. Soc. 242, 399.Hau, G. K. T. and Thomson, R. C.: 1994, Monthly Notices Roy. Astron. Soc. 270, L23.Hawarden, T. G., Elson, R. A. W., Longmore, A. J., Tritton, S. B., and Corwin, H. G.: 1981,

Monthly Notices Roy. Astron. Soc. 196, 747.Hawarden, T. G., Sandell, G., Matthews, H. E., Friberg, P., Watt, G. D., and Smith, P. A.: 1993,

Monthly Notices Roy. Astron. Soc. 260, 844.Haynes, M.: 1981, Astron. J. 86, 1126.Haynes, M. P., Giovanelli, R., and Chincarini, G. L.: 1984, Ann. Rev. Astron. Astrophys. 22, 445.Heckman, T. M.: 1981, Astrophys. J. 250, L59.Heckman, T. M., Baum, S. A., van Breugel, W. J. M., and McCarthy, P.: 1989a, Astrophys. J. 338, 48.Heckman, T. M., Blitz, L., Wilson, A. S., and Armus, L.: 1989b, Astrophys. J. 342, 735.Heikkila, A., Johansson, L. E. B., and Olofsson, H.: 1997, Astron. Astrophys. 319, L21.Heiles, C., Goodman, A. A., McKee, C. F., and Zweibel, E. G.: 1993, in E. H. Levy and J. I. Lunine

(eds.), Protostars and Planets III, The University of Arizona Press, Tucson, p. 279.Heiligman, G. and Schwarzschild, M.: 1979, Astrophys. J. 233, 872.Held, E. V. and Mould, J. R.: 1994, Astron. J. 107, 1307.


Helou, G.: 1986, Astrophys. J. 311, L33.Henkel, C., Baan, W. A., and Mauersberger, R.: 1991, Astron. Astrophys. Rev. 3, 47.Henkel, C. and Mauersberger, R.: 1991, in P. D. Singh (ed.), ‘Astrochemistry of Cosmic Phenomena’,

IAU Symp. 150, 111.Henkel, C. and Mauersberger, R.: 1993, Astron. Astrophys. 274, 730.Henkel, C., Wiklind, T., Wyrowski, F., and Combes, F.: 1996, in P. Shaver (ed.), Science with Large

Millimetre Arrays, ESO-IRAM-NFRA-Onsala Workshop, Springer-Verlag, Berlin, p. 144.Hernquist, L.: 1990, Astrophys. J. 356, 359.Hernquist, L. and Barnes, J. E.: 1991, Nature 354, 210.Hernquist, L. and Spergel, D. N.: 1992, Astrophys. J. 399, L117.Hibbard, J. E., Guhathakurta, P., van Gorkom, J. H., and Schweizer, F.: 1994, Astron. J. 107, 67.Hibbard, J. E. and Mihos, J. C.: 1995, Astron. J. 110, 140.Hjorth, J. and Madsen, J.: 1991, Monthly Notices Roy. Astron. Soc. 253, 703.Hodge, P. W.: 1963, Astron. J. 68, 691.Hodge, P. W.: 1973, Astrophys. J. 182, 671.Hodge, P. W.: 1989, Ann. Rev. Astron. Astrophys. 27, 139.Hogg, D. E., Roberts M. S., and Sandage, A.: 1993, Astron. J. 106, 907.Hollenbach, D. and McKee, C. F.: 1979, Astrophys. J. Suppl. 41, 555.Holtzman, J. A., Faber, S. M., Shaya, E. J. et al.: 1992, Astron. J. 103, 691.Horellou, C., Casoli, F., Combes, F., and Dupraz, C.: 1995, Astron. Astrophys. 298, 743.Hu, E. M.: 1988, in A. C. Fabian (ed.), Cooling Flows in Clusters and Galaxies, Kluwer Academic

Publishers, Dordrecht, p. 73.Hu, E. M.: 1992, Astrophys. J. 391, 608.Hu, E. M., Cowie, L. L., and Wang, Z.: 1985, Astrophys. J. Suppl. 59, 447.Hubble, E.: 1936, The Realm of the Nebula, Dover Publ. Inc., New York.Huchtmeier, W. K.: 1997, Astron. Astrophys. 319, 401.Huchtmeier, W. K. and Bohnenstengel, H.-D.: 1975, Astron. Astrophys. 44, 479.Huchtmeier, W. K. and Richter, O.-G.: 1989, A General Catalog of HI Observations of Galaxies,

Springer-Verlag, New York.Huchtmeier, W. K. and Tammann, G. A.: 1992, Astron. Astrophys. 257, 455.Huchtmeier, W. K., Bregman, J. N., Hogg, D. E., and Roberts, M. S.: 1988, Astron. Astrophys. 198,

L17.Huchtmeier, W. K., Bregman, J. N., Hogg, D. E., and Roberts, M. S.: 1994, Astron. Astrophys. 281,

327.Huchtmeier, W. K., Sage, L. J., and Henkel, C.: 1995, Astron. Astrophys. 300, 675.Hwang, U., Mushotzky, R. F., Loewenstein, M., Markert, T. H., Fukazawa, Y., and Matsumoto, H.:

1997, Astrophys. J. 476, 560.Hui, X., Ford, H. C., Freeman, K. C., and Dopita, M. A.: 1995, Astrophys. J. 449, 592.Im, M., Griffiths, R. E., Ratnatunga, K. U., and Sarajedini, V. L.: 1996, Astrophys. J. 461, L79.Impey, C. D., Wynn-Williams, C. G., and Becklin, E. E.: 1990, Astrophys. J. 356, 62.Inoue, M. Y., Kameno, S., Kawabe, R., Inoue, M., Hasegawa, T., and Tanaka, M.: 1996, Astron. J.

111, 1852.IRAS: 1989, Catalogued Galaxies and Quasars Observed in the IRAS Survey, Version 2, produced

by L. Fullmer and C. Lonsdale, JPL D-1932.Israel, F. P.: 1992, Astron. Astrophys. 265, 487.Israel, F. P., van Dishoeck, E. F., Baas, F., de Graauw, T., and Phillips, T. G.: 1991, Astron. Astrophys.

245, L13.Israel, F. P., van Dishoeck, E. F., Baas, F., Koornneef, J., Black, J. H. and de Graauw, T.: 1990,

Astron. Astrophys. 227, 342.Jackson, J. M., Snell, R. L., Ho, P. T. P., and Barrett, A. H.: 1989, Astrophys. J. 337, 680.Jafelice, L. C.: 1992, Astron. J. 104, 1279.Jafelice, L. C. and Friaca, A. C. S. : 1996, Monthly Notices Roy. Astron. Soc. 280, 438.Jaffe, W.: 1983, Monthly Notices Roy. Astron. Soc. 202, 995.Jaffe, W.: 1987, Astron. Astrophys. 171, 378.Jaffe, W.: 1990, Astron. Astrophys. 240, 254.


Jaffe, W.: 1992, in A. C. Fabian (ed.), Clusters and Superclusters of Galaxies, Kluwer AcademicPublishers, Dordrecht, p. 109.

Jaffe, W. and Bremer, M. N.: 1997, Monthly Notices Roy. Astron. Soc. 284, L1.Jaffe, W., de Bruyn, A. G., and Sijbreng, D.: 1988, in A. C. Fabian (ed.), Cooling Flows in Clusters

and Galaxies, Kluwer Academic Publishers, Dordrecht, p. 145.Jaffe, W., Ford, H. C., Ferrarese, L., van den Bosch, F., and O’Connell, R. W.: 1993, Nature 364, 213.Jaffe, W., Ford, H. C., Ferrarese, L., van den Bosch, F., and O’Connell, R. W.: 1996, Astrophys. J.

460, 214.Jaffe, W., Ford, H. C., O’Connell, R. W., van den Bosch, F. C., and Ferrarese, L.: 1994, Astron. J.

108, 1567.Jaffe, W. and McNamara, B. R.: 1994, Astrophys. J. 434, 110.Jenkins, C. R.: 1982, Monthly Notices Roy. Astron. Soc. 200, 705.Johansson, L. E. B., Olafsson, H., Hjalmarson, A., Gredel, R., and Black, J. H.: 1994, Astron.

Astrophys. 291, 89.Johnson, D. W. and Gottesman, S. T.: 1979, in D. S. Evans (ed.), Photometry, Kinematics and

Dynamics of Galaxies, Proceedings of a Conference held at the University of Texas, Austin,p. 57.

Johnson, D. W. and Gottesman, S. T.: 1983, Astrophys. J. 275, 549.Johnstone, R. M., Fabian, A. C., Edge, A. C., and Thomas, P. A.: 1992, Monthly Notices Roy. Astron.

Soc. 255, 431.Johnstone, R. M., Fabian, A. C., and Nulsen, P. E. J.: 1987, Monthly Notices Roy. Astron. Soc. 224,

75.Jones, C. and Forman, W.: 1984, Astrophys. J. 276, 38.Jore, K. P., Broeils, A. H., and Haynes, M. P.: 1996, Astron. J. 112, 438.Jørgensen, I. and Franx, M.: 1994, Astrophys. J. 433, 533.Jura, M., Kim, D.-W., Knapp, G. R., and Guhathakurta, P.: 1987, Astrophys. J. 312, L11.Kaisler, D., Harris, W. E., Crabtree, D. R., and Richer, H. B.: 1996, Astron. J. 111, 2224.Kauffmann, G., Charlot, S., and White, S. D. M.: 1996, Monthly Notices Roy. Astron. Soc. 283, L117.Katz, N. and Rix, H. W.: 1992, Astrophys. J. 389, L55.Kawara, K. and Taniguchi, Y.: 1993, Astrophys. J. 410, L19.Kazes, I., Troland, T. H., and Crutcher, R. M.: 1991, Astron. Astrophys. 245, L17.Kellermann, K. I. and Thompson, A. R.: 1985, Science 229, 123.Kellermann, K. I., Zensus, J. A., and Cohen, M. H.: 1997, Astrophys. J. 475, L93.Kelson, D. D., van Dokkum, P. G., Franx, M., Illingworth, G. D., and Fabricant, D.: 1997, Astrophys. J.

478, L13.Kennicutt, R. C.: 1989, Astrophys. J. 344, 685.Kim, D.-W.: 1989, Astrophys. J. 346, 653.Kim, D.-W. and Fabbiano, G.: 1995, Astrophys. J. 441, 182.Kim, D.-W., Guhathakurta, P., van Gorkom, J. H., Jura, M., and Knapp, G. R.: 1988, Astrophys. J.

330, 684.Kinzer, R. L., Johnson, W. N., Derner, C. D. et al.: 1995, Astrophys. J. 449, 105.Kissler-Patig, M.: 1997, Astron. Astrophys. 319, 83.Kley, W. and Mathews, W. G.: 1995, Astrophys. J. 438, 100.Knapen, J. H., Hes, R., Beckman, J. E., and Peletier, R. F.: 1991, Astron. Astrophys. 241, 42.Knapp, G. R.: 1987, in T. de Zeeuw (ed.), ‘Structure and Dynamics of Elliptical Galaxies’, IAU Symp.

127, 145.Knapp, G. R. and Patten, B. M.: 1991, Astron. J. 101, 1609.Knapp, G. R. and Rupen, M. P.: 1996, Astrophys. J. 460, 271.Knapp, G. R., Bies, W. E., and van Gorkom, J. H.: 1990, Astron. J. 99, 476.Knapp, G. R., Guhathakurta, P., Kim, D.-W., and Jura, M.: 1989, Astrophys. J. Suppl. 70, 329.Knapp, G. R., Gunn, J. E., and Wynn-Williams, C. G.: 1992, Astrophys. J. 399, 76.Knapp, G. R., Rupen, M. P., Fich, M., Harper, D. A., and Wynn-Williams, C. G.: 1997,

Astron. Astrophys. 315, L75.Knapp, G. R., Turner, E. L., and Cunniffe, P. E.: 1985a, Astron. J. 90, 454.Knapp, G. R., van Driel, W., and van Woerden, H.: 1985b, Astron. Astrophys. 142, 1.


Koekemoer, A. M., Henkel, C., Greenhill, L. J., Dey, A., van Breugel, W., Codella, C., and Antonucci,R.: 1995, Nature 378, 697.

Koo, D. C.: 1996, in R. Bender and R. L. Davies (eds.), ‘New Light on Galaxy Evolution’, IAU Symp.171, 217.

Kormendy, J.: 1977, Astrophys. J. 218, 333.Kormendy, J.: 1987a, in S. M. Faber (ed.), Nearly Normal Galaxies: From the Planck Time to the

Present, Springer-Verlag, New York, p. 163.Kormendy, J.: 1987b, in T. de Zeeuw (ed.), ‘Structure and Dynamics of Elliptical Galaxies’, IAU

Symp. 127, 17.Kormendy, J.: 1989, Astrophys. J. 342, L63.Kormendy, J. and Bender, R.: 1994, in G. Meylan and P. Prugniel (eds.), ESO Conf. and Workshop

Proc. 49, Dwarf Galaxies, ESO, Garching, p. 161.Kormendy, J. and Bender, R.: 1996, Astrophys. J. 464, L119.Kormendy, J. and Djorgovski, S.: 1989, Ann. Rev. Astron. Astrophys. 27, 235.Kormendy, J. and Sanders, D. B.: 1992, Astrophys. J. 390, L53.Kormendy, J., Byun, Y.-I., Ajhar, E. A. et al.: 1996, in R. Bender and R. L. Davies (eds.), ‘New Light

on Galaxy Evolution’, IAU Symp. 171, 105.Kotanyi, C. G. and Ekers, R. D.: 1979, Astron. Astrophys. 73, L1.Kritsuk, A. G.: 1992, Astron. Astrophys. 261, 78.Kritsuk, A. G.: 1997, Monthly Notices Roy. Astron. Soc. 284, 327.Krumm, N., van Driel, W., and van Woerden, H.: 1985, Astron. Astrophys. 144, 202.Kulkarni, S. R. and Heiles, C.: 1988, in G. L. Verschuur and K. I. Kellerman (eds.), Galactic and

Extragalactic Radio Astronomy, Springer-Verlag, New York, p. 132.Kumar, C. K. and Thonnard, N.: 1983, Astron. J. 88, 260.Laing, R. A., Riley, J. M., and Longair, M.: 1983, Monthly Notices Roy. Astron. Soc. 204, 151.Lake, G., Schommer, R. A., and van Gorkom, J. H.: 1987, Astrophys. J. 314, 57.Larson, R. B.: 1981, Monthly Notices Roy. Astron. Soc. 194, 809.Lauberts, A.: 1982, The ESO/Uppsala Survey of the ESO(B) Atlas, ESO, Garching.Lauer, T. R., Ajhar, E. A., Byun, Y.-I., Dressler, A., Faber, S. M., Grillmair, C., Kormendy, J.,

Richstone, D., and Tremaine, S.: 1995, Astron. J. 110, 2622.Lazareff, B., Castets, A., Kim, D.-W., and Jura, M.: 1989, Astrophys. J. 336, L13.Lee, M. G., Freedman, W. L., and Madore, B. F.: 1993, Astron. J. 106, 964.Lee, Y.-W.: 1994, Astrophys. J. 430, L113.Lees, J. F., Knapp, G. R., Rupen, M. P., and Phillips, T. G.: 1991, Astrophys. J. 379, 177.Lester, D. F., Zink, E. C., Doppmann, G. W., Gaffney, N. I., Harvey, P. M., and Smith, B. J.: 1995,

Astrophys. J. 439, 185.Levinson, A., Laor, A., and Vermeulen, R. C.: 1995, Astrophys. J. 448, 589.Lewis, B. M.: 1970, Observatory 90, 264.Li, J. G., Seaquist, E. R., and Sage, L. J.: 1993, Astrophys. J. 411, L71.Lo, K. Y., Sargent, W. L. W., and Young, K.: 1993, Astron. J. 102, 507.Loose, H.-H. and Thuan, T. X.: 1996, Astrophys. J. 309, 59.Loewenstein, M.: 1990, Astrophys. J. 349, 471.Loewenstein, M., Zweibel, E. G., and Begelman, M. C.: 1991, Astrophys. J. 377, 392.Maccagni, D., Garilli, B., Gioia, I. M., Maccacaro, T., Vettolani, G., and Wolter, A.: 1988, Astrophys. J.

334, L1.Macchetto, F., Pastoriza, M., Caon, N., Sparks, W. B., Giavalisco, M., Bender, R., and Capaccioli,

M.: 1996, Astron. Astrophys. Suppl. 120, 463.Macchetto, F. and Sparks, W. B.: 1992, in G. Longo, M. Capaccioli, and G. Busarello (eds.), Mor-

phological and Physical Classification of Galaxies, Kluwer Academic Publishers, Dordrecht,p. 191.

Mahabal, A., Kembhavi, A., Singh, K. P., Bhat, P. N., and Prabhu, T. P.: 1996, Astrophys. J. 457, 598.Maloney, P. R., Begelman, M. C., and Rees M. J.: 1994, Astrophys. J. 432, 606.Mamon, G. A.: 1992, Astrophys. J. 401, L3.Mathews, W. G.: 1988, in A. C. Fabian (ed.), Cooling Flows in Clusters and Galaxies, Kluwer

Academic Publishers, Dordrecht, p. 279.


Mathews, W. G.: 1989, Astron. J. 97, 42.Mathews, W. G.: 1990, Astrophys. J. 354, 468.Mathews, W. G., Loewenstein, M.: 1986, Astrophys. J. 306, L7.Matsushita, K., Makishima, K., Awaki, H. et al.: 1994, Astrophys. J. 436, L41.Mauersberger, R. and Henkel, C.: 1993, Rev. Mod. Astron. 6, 69.Mauersberger, R., Henkel, C., Whiteoak, J. B., Chin, Y.-N., and Tieftrunk, A. R.: 1996a,

Astron. Astrophys. 309, 705.Mauersberger, R., Henkel, C., Wielebinski, R., Wiklind, T., and Reuter, H.-P.: 1996b,

Astron. Astrophys. 305, 421.Mauersberger, R., Henkel, C., Wilson, T. L., and Walmsley, C. M.: 1986, Astron. Astrophys. 162, 199.Mayall, N. U.: 1936, Publ. Astron. Soc. Pacific 48, 14.Mayall, N. U.: 1939, Publ. Astron. Soc. Pacific 51, 282.McNamara, B. R., Bergman, J. N., and O’Connell, R. W.: 1990, Astrophys. J. 360, 20.McNamara, B. R. and Jaffe, W.: 1994, Astron. Astrophys. 281, 673.McNamara, B. R. and O’Connell, R. W.: 1992, Astrophys. J. 393, 579.McNamara, B. R., O’Connell, R. W., and Sarazin, C. L.: 1996, Astron. J. 112, 91.Meiksin, A.: 1988, in A. C. Fabian (ed.), Cooling Flows in Clusters and Galaxies, Kluwer Academic

Publishers, Dordrecht, p. 47.Meiksin, A.: 1990, Astrophys. J. 352, 466.Merrifield, M. R. and Kuijken, K.: 1994, Astrophys. J. 432, 575.Merritt, D. and de Zeeuw, T.: 1983. Astrophys. J. 267, L19.Michard, R.: 1994, Astron. Astrophys. 288, 401.Mihos, J. C., Bothun, G. D., and Richstone, D. O.: 1993, Astrophys. J. 418, 82.Miller, L.: 1986, Monthly Notices Roy. Astron. Soc. 220, 713.Miller, W. W. and Cox, D. P.: 1993, Astrophys. J. 417, 579.Minkowski, R.: 1957, in H. C. van der Hulst (ed.), ‘Radio Astronomy’, IAU Symp. 4, 107.Mirabel, I. F. and Wilson, A. S.: 1984, Astrophys. J. 277, 92.Mirabel, I. F., Sanders, D. B., and Kazes, I.: 1989, Astrophys. J. 340, L9.Mitchell, G. F. and Deveau, T. J.: 1983, Astrophys. J. 266, 646.Miyoshi, M., Moran, J. M., Herrnstein, J., Greenhill, L. J., Nakai, N., Diamond, P., and Inoue, M.:

1995, Nature 373, 127.Molinari, E., Buzzoni, A., Chincarini, G., and Pedrana, M. D.: 1994, Astron. Astrophys. 292, 54.Mollenhoff, C.: 1982, Astron. Astrophys. 108, 130.Mollenhoff, C. and Bender, R.: 1987, Astron. Astrophys. 174, 63.Mollenhoff, C. Hummel, E., and Bender, R.: 1992, Astron. Astrophys. 255, 35.Morganti, R.: 1992, in G. Longo, M. Capaccioli, and G. Busarello (eds.), Morphological and Physical

Classification of Galaxies, Kluwer Academic Publishers, Dordrecht, p. 415.Morganti, R., Sadler, E. M., Oosterloo, T., Pizzella, A., and Bertola, F.: 1997, Astron. J. 113, 937.Murray, S. D. and Balbus, S. A.: 1992, Astrophys. J. 395, 99.Mushotzky, R. F., Loewenstein, M., Awaki, H., Makishima, K., and Matsumoto, H.: 1994,

Astrophys. J. 436, L79.Mushotzky, R. F. and Szymkowiak, A. E.: 1988, in A. C. Fabian (ed.), Cooling Flows in Clusters and

Galaxies, Kluwer Academic Publishers, Dordrecht, p. 53.Nesci, R., Perola, G. C., and Wolter, A.: 1995, Astron. Astrophys. 299, 34.Neufeld, D. A., Maloney, P. R., and Conger, S.: 1994, Astrophys. J. 436, L127.Nicholson, R. A., Bland-Hawthorne, J., and Taylor, K.: 1992, Astrophys. J. 387, 503.Nieto, J.-L. and Lagoute, Ch.: 1992, in G. Longo, M. Capaccioli, G. Busarello (eds), Morphological

and Physical Classification of Galaxies, Kluwer Academic Publishers, Dordrecht, p. 387.Nieto, J.-L., Bender, R., Poulain, P., and Surma, P.: 1992, Astron. Astrophys. 257, 97.Nilson, P.: 1973, Uppsala General Catalogue of Galaxies, Uppsala.Nørgaard-Nielsen, H. U., Goudfrooij, P., Jørgensen, H. E., and Hansen, L.: 1993, Astron. Astrophys.

279, 61.Norman, C. and Meiksin, A.: 1996, Astrophys. J. 468, 97.Nulsen, P. E. J.: 1986, Monthly Notices Roy. Astron. Soc. 221, 377.O’Dea, C. P., Baum, S. A., and Gallimore, J. F.: 1994a, Astrophys. J. 436, 669.


O’Dea, C. P., Baum, S. A., Maloney, P. R., Tacconi, L. J., and Sparks, W. B.: 1994b, Astrophys. J.422, 467.

Ohta, K., Tomita, A., Saito, M., Sasaki, M., and Nakai, N.: 1993, Publ. Astron. Soc. Japan 45, L21.Peletier, R. F. and Christodoulou, D. M.: 1993, Astron. J. 105, 1378.Peletier, R. F., Valentijn, E. A., and Jameson, R. F.: 1987, in T. de Zeeuw (ed.), ‘Structure and

Dynamics of Elliptical Galaxies’, IAU Symp. 127, 443.Pen, U.-L.: 1994, Astrophys. J. 431, 104.Peterson, R. G. and Caldwell, N.: 1993, Astron. J. 105, 1411.Pfeffermann, E. et al.: 1986, Proc. SPIE 733, 519.Phillips, M. M., Jenkins, C. R., Dopita, M. A., Sadler, E. M., and Binette, L.: 1986, Astron. J. 91,

1062.Phillips, T. G., Ellison, B. N., Keene, J. B., Leighton, R. B., Howard, R. J., Masson, C. R., Sanders,

D. B., Veidt, B., and Young, K.: 1987, Astrophys. J. 322, L73.Pinkney, J., Holtzman, J., Garasi, C. et al.: 1996, Astrophys. J. 468, L13.Pistinner, S., Levinson, A., and Eichler, D.: 1996, Astrophys. J. 467, 162.Pistinner, S. and Shaviv, G.: 1995, Astrophys. J. 446, L11.Pistinner, S. and Shaviv, G.: 1996, Astrophys. J. 459, L147.Pogge, R. W. and Eskridge, P. B.: 1993, Astron. J. 106, 1405.Prada, F., Gutierrez, C. M., Peletier, R. F., and McKeith, C. D.: 1996, Astrophys. J. 463, L9.Prantzos, N. and Aubert, O.: 1995, Astron. Astrophys. 302, 69.Price, J. S. and Gullixson, C. A.: 1989, Astrophys. J. 337, 658.Pringle, J. E.: 1989, Monthly Notices Roy. Astron. Soc. 239, 479.Prugniel, P. and Simien, F.: 1997, Astron. Astrophys. 321, 111.Prugniel, P., Nieto, J.-L., and Simien, F.: 1987, Astron. Astrophys. 173, 49.Pryor, C.: 1992, in G. Longo, M. Capaccioli, and G. Busarello (eds.), Morphological and Physical

Classification of Galaxies, Kluwer Academic Publishers, Dordrecht, p. 163.Quillen, A. C., de Zeeuw, P. T., Phinney, E. S., and Phillips, T. G.: 1992, Astrophys. J. 391, 121.Quillen, A. C., Graham, J. R., and Frogel, J. A.: 1993, Astrophys. J. 412, 550.Quinn, P. J. and Hernquist, L.: 1987, in T. de Zeeuw (ed.), ‘Structure and Dynamics of Elliptical

Galaxies’, IAU Symp. 127, 249.Ragaranjan, F. V. N., Fabian, A. C., Forman, W. R., and Jones, C.: 1995, Monthly Notices Roy. Astron.

Soc. 272, 665.Raimond, E., Faber, S. M., Gallagher, J. S., and Knapp. G. R.: 1981, Astrophys. J. 246, 708.Rampazzo, R. and Sulentic, J. W.: 1992, Astron. Astrophys. 259, 43.Reduzzi, L., Langhetti, M., and Rampazzo, R.: 1996, Monthly Notices Roy. Astron. Soc. 282, 149.Reid, N., Boisson, C., and Sanson, A. E.: 1994, Monthly Notices Roy. Astron. Soc. 269, 713.Reif, K., Mebold, U., Goss, W. M., van Woerden, H., and Siegman, B.: 1982, Astron. Astrophys. 50,

451.Renzini, A.: 1996, in R. Bender and R. L. Davies (eds.), ‘New Light on Galaxy Evolution’, IAU

Symp. 171, 131.Renzini, A. and Ciotti, L.: 1993, Astrophys. J. 416, L49.Reuter, H.-P., Pohl, M., Lesch, H., and Sievers, A. W.: 1993, Astron. Astrophys. 277, 21.Richer, H., Crabtree, D. R., Fabian, A. C., and Lin, D. N. C.: 1993, Astron. J. 105, 877.Richter, O.-G., Sackett, P. D., and Sparke, L. S.: 1994, Astron. J. 107, 99.Rix, H.-W. and White, S. D. M.: 1990, Astrophys. J. 362, 52.Rix, H.-W. and Zaritsky, D.: 1995, Astrophys. J. 447, 82.Roberts, M. S.: 1970, Astrophys. J. 161, L9.Roberts, M. S., Hogg, D. E., Bregman, J. N., Forman, W. R., and Jones, C.: 1991, Astrophys. J. Suppl.

75, 751.Robinson, A., Binette, L., Fosbury, R. A. E., and Tadhunter, C. N.: 1987, Monthly Notices Roy. Astron.

Soc. 227, 97.Romani, R. W. and Moaz, D.: 1992, Astrophys. J. 386, 36.Rose, J. A., Bower, R. G., Caldwel, N., Ellis, R. S., Sharples, R. M., and Teague, P.: 1994, Astron. J.

108, 2054.Rosner, R. and Tucker, W. H.: 1989, Astrophys. J. 338, 761.


Rubio, M., Lequeux, J., and Boulanger, F.: 1993, Astron. Astrophys. 271, 9.Rydbeck, G., Wiklind, T., Cameron, M., Wild, W., Eckart, A., Genzel, R., and Rothermel, H.: 1993,

Astron. Astrophys. 270, L13.Ryden, B. S.: 1996, Astrophys. J. 461, 146.Sadler, E. M.: 1987, in T. de Zeeuw (ed.), ‘Structure and Dynamics of Elliptical Galaxies’, IAU Symp.

127, 125.Sadler, E. M.: 1997, Publ. Astron. Soc. Austr. 14, 45.Sadler, E. M. and Gerhard, O. E.: 1985, Monthly Notices Roy. Astron. Soc. 214, 177.Sage, L. J.: 1989, Astrophys. J. 344, 200.Sage, L. J.: 1990, Astron. Astrophys. 239, 125.Sage, L. J.: 1993, Astron. Astrophys. 272, 123.Sage, L. J. and Galletta, G.: 1993, Astrophys. J. 419, 544.Sage, L. J. and Galletta, G.: 1994, Astron. J. 108, 1633.Sage, L. J. and Solomon, P. M.: 1989, Astrophys. J. 342, L15.Sage, L. J. and Wrobel, J. M.: 1989, Astrophys. J. 344, 204.Sage, L. J., Salzer, J. J., Loose, H.-H., and Henkel, C.: 1992, Astron. Astrophys. 265, 19.Saglia, R. P., Bertin, G., Bertola, F., Danziger, J., Dejonghe, H., Sadler, E. M., Stiavelli, M., de Zeeuw,

P. T., and Zeilinger, W. W.: 1993, Astrophys. J. 403, 567.Sakamoto, S.: 1996, Astrophys. J. 462, 215.Sancisi, R., Thonnard, N., and Ekers, R. D.: 1987, Astrophys. J. 315, L39.Sancisi, R., van Woerden, H., Davies, R. D., and Hart, L.: 1984, Monthly Notices Roy. Astron. Soc.

210, 497.Sandage, A. and Binggeli, B.: 1984, Astron. J. 89, 919.Sandage, A. and Fomalont, E.: 1993, Astrophys. J. 407, 14.Sandage, A. and Hoffman, G. L.: 1991, Astrophys. J. 379, L45.Sandage, A. and Tammann, G. A.: 1987, A Revised Shapley Ames Catalog of Bright Galaxies (RSA),

second edition, Carnegie Institution, Washington D.C.Sanders, D. B.: 1991, in F. Combes and F. Casoli (eds.), ‘Dynamics of Galaxies and Their Molecular

Cloud Distributions’, IAU Symp. 146, 417.Sanders, D. B. and Mirabel, I. F.: 1985, Astrophys. J. 298, L31.Sansom, A. E., Danziger, I. J., Ekers, R. D., Fosbury, R. A. E., Goss, W. M., Monk, A. S., Shaver,

P. A., Sparks, W. B., and Wall, J. V.: 1987, Monthly Notices Roy. Astron. Soc. 229, 15.Sarazin, C. L.: 1990, in H. A. Thronson and J. M. Shull (eds.), The Interstellar Medium in Galaxies,

Second Wyoming Conference, Kluwer Academic Publishers, Dordrecht, p. 201.Sarazin, C. L., O’Connell, R. W., and McNamara, B. R.: 1992, Astrophys. J. 397, L31.Schade, D., Carlberg, R. G., Yee, H. K. C., Lopez-Cruz, O., and Ellingsen, E.: 1996, Astrophys. J.

464, L63.Schaller, G., Schaerer, D., Meynet, G., and Maeder, A.: 1992, Astron. Astrophys. Suppl. 96, 269.Schechter, P. L.: 1980, Astron. J. 85, 801.Schilke, P. and Walmsley, C. M., Pineau de Forets, G., Roueff, E., Flower, D. R., and Guilloteau, S.:

1992, Astron. Astrophys. 256, 595.Schiminovich, D., van Gorkom, J. H., van der Hulst, J. M., and Kasow, S.: 1994, Astrophys. J. 423,

L101.Schiminovich, D., van Gorkom, J. H., van der Hulst, J. M., and Malin, D. F.: 1995, Astrophys. J. 444,

L77.Schombert, J. H., Barsony, M., and Hanlon, P. C.: 1993, Astrophys. J. 416, L61.Schombert, J. H.: 1992, in G. Longo, M. Capaccioli, and G. Busarello (eds.), Morphological and

Physical Classification of Galaxies, Kluwer Academic Publishers, Dordrecht, p. 53.Schreier, E. J., Capetti, A., Maccheto, F., Sparks, W. B., and Ford, H. J.: 1996, Astrophys. J. 459, 535.Schwarzschild, M.: 1979, Astrophys. J. 232, 236.Schweizer, F.: 1980, Astrophys. J. 237, 303.Schweizer, F.: 1982, Astrophys. J. 252, 455.Schweizer, F.: 1987, in T. de Zeeuw (ed.), ‘Structure and Dynamics of Elliptical Galaxies’, IAU Symp.

127, 109.Schweizer, F. and Seitzer, P.: 1992, Astron. J. 104, 1039.


Schweizer, F., van Gorkom, J. H., and Seitzer, P.: 1989, Astrophys. J. 338, 770.Scorza, C. and Bender, R.: 1995, Astron. Astrophys. 293, 20.Scorza, C. and Bender, R.: 1996, in R. Bender and R. L. Davies (eds.), ‘New Light on Galaxy

Evolution’, IAU Symp. 171, 55.Seab, C. G.: 1987, in D. J. Hollenbach and H. A. Thronson (eds.), Interstellar Processes, D. Reidel

Publ. Co., Dordrecht, p. 491.Seaquist, E. R. and Bell, M. B.: 1986, Astrophys. J. 303, L67.Seaquist, E. R. and Bell, M. B.: 1990, Astrophys. J. 364, 94.Sellwood, J. A. and Merritt, D.: 1994, Astrophys. J. 425, 530.Shane, W. W.: 1980, Astron. Astrophys. 82, 314.Shields, J. C.: 1991, Astron. J. 102, 1314.Shields, J. C. and Filippenko, A. V.: 1990, Astrophys. J. 353, L7.Shostak, G. S.: 1987, Astron. Astrophys. 175, 4.Skillman, E. D.: 1996, in E. D. Skillman (ed.), ASP Conf. Ser. 106, The Minnesota Lectures on

Extragalactic Neutral Hydrogen, Astron. Soc. Pac., San Francisco, p. 208.Smith, E. P., O’Connell, R. W., Bohlin, R. C. et al.: 1992, Astrophys. J. 395, L49.Smith, S., Kennel, C. F., and Coroniti, F. V.: 1993, Astrophys. J. 412, 82.Sodroski, T. J., Odegard, N., Dwek, E. et al.: 1995, Astrophys. J. 452, 262.Sofue, Y.: 1996, Astrophys. J. 458, 120.Sofue, Y. and Wakamatsu, K.: 1993, Publ. Astron. Soc. Japan 45, 529.Solomon, P. M. and Barrett, J. W.: 1991, in F. Combes and F. Casoli (eds.), ‘Dynamics of Galaxies

and Their Molecular Cloud Distributions’, IAU Symp. 146, 235.Solomon, P. M., Downes, D., and Radford, S. J. E.: 1992a, Astrophys. J. 387, L55.Solomon, P. M., Downes, D., and Radford, S. J. E.: 1992b, Astrophys. J. 398, L29.Solomon, P. M., Downes, D., Radford, S. J. E., and Barrett, J. W.: 1997, Astrophys. J. 478, 144.Sparke, L. S.: 1986, Monthly Notices Roy. Astron. Soc. 219, 657.Sparke, L. S.: 1996, Astrophys. J. 473, 810.Sparks, W. B.: 1992a, Astrophys. J. 399, 66.Sparks, W. B.: 1992b, in I. J. Danziger, W. W. Zeilinger, and K. Kjar (eds.), ESO Conf. and Workshop

Proc. 45, Structure, Dynamics and Chemical Evolution of Elliptical Galaxies, ESO Garching,p. 567.

Sparks, W. B., Macchetto, F., and Golombek, D.: 1989, Astrophys. J. 345, 153.Sparks, W. B., Wall, J. V., Thorne, D. J., Jorden, P. R., van Breda, I. G., Rudd, P. J., and Jørgensen,

H. E.: 1985, Monthly Notices Roy. Astron. Soc. 217, 87.Spergel, D. N. and Blitz, L.: 1992, Nature 357, 665.Spinrad, H., Djorgovski, S., Marr, J., and Aguilar, L.: 1985, Publ. Astron. Soc. Pacific 97, 932.Stanford, S. A. and Bushouse, H. A.: 1991, Astrophys. J. 371, 92.Stark, A. A., Knapp, G. R., Bally, J., Wilson, R. W., Penzias, A. A., and Rowe, H. E.: 1986,

Astrophys. J. 310, 660.Statler, T. S.: 1991, Astron. J. 102, 882.Steiman-Cameron, T. Y. and Durisen, R. H.: 1982, Astrophys. J. 263, L51.Strong, A. W., Bloemen, J. B. G. M., Dame, T. M., Grenier, I. A., Hermsen, W., Lebrun, F., Nyman, L.-

A, Pollock, A. M. T., and Thaddeus, P.: 1988, Astron. Astrophys. 207, 1.Sumi, D.: 1988, in A. C. Fabian (ed.), Cooling Flows in Clusters and Galaxies, Kluwer Academic

Publishers, Dordrecht, p. 381.Tabor, G. and Binney, J.: 1993, Monthly Notices Roy. Astron. Soc. 263, 323.Tacconi, L. J., Tacconi-Garman, L. E., Thornley, M., and van Woerden, H.: 1991, Astron. Astrophys.

252, 541.Tadhunter, C. N., Fosbury, R. A. E., and Quinn, P. J.: 1989, Monthly Notices Roy. Astron. Soc. 240,

225.Tanaka, Y., Inoue, H., and Holt, S. S.: 1994, Publ. Astron. Soc. Japan 46, L37.Taniguchi, Y., Murayama, T., Nakai, N., Suzuki, M., and Kameya, O.: 1994, Astron. J. 108, 468.Taniguchi, Y., Sofue, Y., Wakamatsu, K.-I., and Nakai, N.: 1990a, Astron. J. 100, 1086.Taniguchi, Y., Takato, N., and Aoki, T.: 1990b, Publ. Astron. Soc. Japan 42, 819.Thomasson, M. and Donner, K. J.: 1993, Astron. Astrophys. 272, 153.


Thronson, H. A. and Telesco, C. M.: 1986, Astrophys. J. 311, 98.Thronson, H. A., Tacconi, L., Kenney, J., Greenhouse, M. A., Margulis. M., Tacconi-Garman, L., and

Young, J. S.: 1989, Astrophys. J. 344, 747.Thronson, H. A., Wilton, C., and Ksir, A.: 1991, Monthly Notices Roy. Astron. Soc. 252, 543.Thuan, T. X. and Puschell, J. J.: 1989, Astrophys. J. 346, 34.Tohline, J. E., Simonson, G. F., and Caldwell, N.: 1982, Astrophys. J. 252, 92.Trembley, B. and Merritt, D.: 1995, Astron. J. 110, 1039.Trinchieri, G. and Di Serego Alighieri, S.: 1991, Astron. J. 101, 1647.Trinchieri, G., Kim, D.-W., Fabbiano, G., and Canizares, C. R. C.: 1994, Astrophys. J. 428, 555.Trumper, J.: 1983, Adv. Space Res. 2 (4), 241.Tsai, J. C. and Mathews, W. G.: 1995, Astrophys. J. 448, 84.Vader, J. P. and Vigroux, L.: 1991, Astron. Astrophys. 246, 32.Valentijn, E. A.: 1988, in A. C. Fabian (ed.), Cooling Flows in Clusters and Galaxies, Kluwer

Academic Publishers, Dordrecht, p. 189.Valentijn, E. A. and Bijleveld, W.: 1983, Astron. Astrophys. 125, 223.van Albada, T. S.: 1982, Monthly Notices Roy. Astron. Soc. 201, 939.van Albada, T. S., Kotanyi, C. G., and Schwarzschild, M.: 1982, Monthly Notices Roy. Astron. Soc.

198, 303.van den Bergh, S. and Tammann, A.: 1991, Ann. Rev. Astron. Astrophys. 29, 363.van den Bosch, F. C., Ferrarese, L., Jaffe, W., Ford, H. C., and O’Connell, R. W.: 1994, Astron. J.

108, 1579.van der Hulst, J. M.: 1979, Astron. Astrophys. 75, 97.van der Hulst, J. M., Golisch, W. F., and Haschick, A. D.: 1983, Astrophys. J. 264, L37.van der Marel, R. P.: 1991, Monthly Notices Roy. Astron. Soc. 253, 710.van Dyke Dixon, W., Davidson, A. F., and Ferguson, H. C.: 1996, Astron. J. 111, 130.van Dokkum, P. G. and Franx, M.: 1995, Astron. J. 110, 2027.van Dokkum, P. G. and Franx, M.: 1996, Monthly Notices Roy. Astron. Soc. 281, 985.van Driel, W., Balkowski, C., and van Woerden, W.: 1989, Astron. Astrophys. 218, 49.van Driel, W., Combes, F., Casoli, F. et al.: 1995, Astron. J. 109, 942.van Driel, W., Rots, A. H., and van Woerden, H.: 1988, Astron. Astrophys. 204, 39.van Driel, W. and van Woerden, H.: 1989, Astron. Astrophys. 225, 317.van Driel, W. and van Woerden, H.: 1991, Astron. Astrophys. 243, 71.van Gorkom, J. H.: 1992, in G. Longo, M. Capaccioli, and G. Busarello (eds.), Morphological and

Physical Classification of Galaxies, Kluwer Academic Publishers, Dordrecht, p. 233.van Gorkom, J. H., Knapp, G. R., Ekers, R. D., Ekers, D. D., Lang, R. A. and Polk, K. S.: 1989,

Astron. J. 97, 708.van Gorkom, J. H., Knapp, G. R., Raimond, E., Faber, S. M., and Gallagher, J. S.: 1986, Astron. J.

91, 791.van Gorkom, J. H., Schechter, P. L., and Kristian, J.: 1987, Astrophys. J. 314, 457.van Gorkom, J. H., van der Hulst, J. M., Haschick, A. D., and Tubbs, A. D.: 1990, Astron. J. 99, 1781.van Langefelde, H. J., van Dishoeck, E. F., Sevenster, M. N., and Israel, F. P.: 1995, Astrophys. J.

448, L123.Varnas, S. R., Bertola, F., Galletta, G., Freeman, K. C., and Carter, D.: 1987, Astrophys. J. 313, 69.Vermeulen, R. C., Readhead, A. C. S., and Backer, D. C.: 1994, Astrophys. J. 430, L41.Veron-Cetty, M.-P. and Veron, P.: 1988, Astron. Astrophys. 204, 28.Verter, F.: 1985, Astrophys. J. Suppl. 57, 261.Verter, F.: 1987, Astrophys. J. Suppl. 65, 555.Vigroux, L., Lachieze Rey, M., Souviron, J., and Vader, J. P.: 1985, in O.-G. Richter and B. Binggeli

(eds.), ESO Conf. and Workshop Proc. 20, The Virgo Cluster of Galaxies, ESO, Garching, p. 297.Voit, G. M. and Donahue, M.: 1995, Astrophys. J. 452, 164.Voshchinnikov, N. V. and Khersonskij, V. K.: 1984, Astrophys. Space Sci. 103, 301.Walsh, D. E. P., van Gorkom, J. H., Bies, W. E., Katz, N., Knapp, G. R., and Wallington, S.: 1990,

Astrophys. J. 352, 532.Wang, Q. and Stocke, J. T.: 1993, Astrophys. J. 408, 71.Wang, Z., Kenney, J. D. P., and Ishizuki, S.: 1992a, Astron. J. 104, 2097.


Wang, Z., Schweizer, F., and Scoville, N. Z.: 1992b, Astrophys. J. 396, 510.Walker, R. C., Romney, J. D., and Benson, J. M.: 1994, Astrophys. J. 430, L45.Walmsley, C. M., Gusten, R., Angerhofer, P., Churchwell, E., and Mundy, L.: 1996, Astron. Astrophys.

155, 129.Wardle, M. and Knapp, G. R.: 1986, Astron. J. 91, 23.Warren-Smith, R. F. and Berry, D. S.: 1983, Monthly Notices Roy. Astron. Soc. 205, 889.Watson, D. M., Guptil, M. T., and Buchholz, L. M.: 1994, Astrophys. J. 420, L21.Weil, M. L. and Hernquist, L.: 1993, Astrophys. J. 405, 142.Welch, G. A. and Mitchell, G. F.: 1994, Astrophys. J. 427, 770.Welch, G. A., Mitchell, G. F., and Yi, S.: 1996, Astrophys. J. 470, 781.Weliachew, L. and Gottesman, S. T.: 1973, Astron. Astrophys. 24, 59.White, D. A., Fabian, A. C., Johnstone, R. M., Mushotzky, R. F., and Arnaud, K. A.: 1991,

Monthly Notices Roy. Astron. Soc. 252, 72.White, R. E. and Chevalier, R. A.: 1983, Astrophys. J. 275, 69.White, R. E. and Sarazin, C. L.: 1991, Astrophys. J. 367, 476.Whiteoak, J. B., Gardner, F. F., and Hoglund, B.: 1980, Monthly Notices Roy. Astron. Soc. 190, 17.Whitmore, B. C., Lucas, R. A., McElroy, D. B., Steiman-Cameron, T. Y., Sackett, P. D., and Olling,

R. P.: 1990, Astron. J. 100, 1489.Wiklind, T.: 1991, in F. Combes and F. Casoli (eds.), ‘Dynamics of Galaxies and Their Molecular

Cloud Distributions’, IAU Symp. 146, 33.Wiklind, T. and Combes, F.: 1997, Astron. Astrophys., in press.Wiklind, T. and Henkel, C.: 1989, Astron. Astrophys. 225, 1.Wiklind, T. and Henkel, C.: 1990, Astron. Astrophys. 227, 394.Wiklind, T. and Henkel, C.: 1992a, Astron. Astrophys. 257, 437.Wiklind, T. and Henkel, C.: 1992b, in I. J. Danziger, W. W. Zeilinger, and K. Kjar (eds.), ESO Conf.

and Workshop Proc. 45, Structure, Dynamics and Chemical Evolution of Elliptical Galaxies,ESO, Garching, p. 599.

Wiklind, T. and Henkel, C.: 1995, Astron. Astrophys. 297, L71.Wiklind, T. and Henkel, C.: 1997, Astron. Astrophys., in preparation.Wiklind, T. and Rydbeck, G.: 1986, Astron. Astrophys. 164, L22.Wiklind, T., Henkel, C., and Sage, L. J.: 1993, Astron. Astrophys. 271, 71.Wiklind, T., Combes, F., and Henkel, C.: 1995, Astron. Astrophys. 297, 643.Wiklind, T., Combes, F., Henkel, C., and Wyrowski, F.: 1997, Astron. Astrophys., in press.Wilcots, E. M., Hodge, P. W., Eskridge, P. B., Bertola, F., and Buson, L.: 1990, Astrophys. J. 364, 87.Wild, W., Eckart, A., Cameron, M., Genzel, R., Rothermel, H., Rydbeck, G., and Wiklind, T.: 1996,

‘25 Years CO’, IAU Symp., in press.Wilkinson, P. N., Polatides, A. G., Readhead, A. C. S., Xu, W., and Pearson, T. J.: 1994, Astrophys. J.

432, L87.Wilson, C. D.: 1995, Astrophys. J. 448, L97.Wirth, A. and Gallagher, J. S.: 1984, Astrophys. J. 282, 85.Wirth, A., Kenyon, S. J., and Hunter, D. A.: 1983, Astrophys. J. 269, 102.Wise, M. W. and Silva, D. R.: 1996, Astrophys. J. 461, 155.Witt, A. N., Thronson, H. A., and Capuano, J. M.: 1992, Astrophys. J. 393, 611.Wouterloot, J. G. A., and Walmsley, C. M.: 1986, Astron. Astrophys. 168, 237.Wrobel, J. M. and Kenney, J. D. P.: 1992, Astrophys. J. 399, 94.Wunderlich, E., Klein, U., and Wielebinski, R.: 1987, Astron. Astrophys. Suppl. 69, 487.Yee, H. K. C., Sawicki, M. J., Ellingson, E., and Carlberg, R. G.: 1996, in A. Buzzoni, A. Renzini,

and A. Serrano (eds.), ASP Conf. Ser. 86, Fresh Views of Elliptical Galaxies, BookCrafters Inc.,San Francisco, p. 301.

Yoshida, T., Hattori, M., and Habe, A.: 1991, Monthly Notices Roy. Astron. Soc. 248, 630.Young, J. S. and Knezek, P. M.: 1989, Astrophys. J. 347, L55.Young, J. S. and Lo, K. Y.: 1996, Astrophys. J. 464, L59.Young, J. S. and Scoville, N. Z.: 1991, Ann. Rev. Astron. Astrophys. 29, 581.Young, J. S., Scoville, N. Z., and Brady, E.: 1985, Astrophys. J. 288, 487.Young, J. S., Xie, S., Kenney, J. D. P., and Rice, W. L.: 1989, Astrophys. J. Suppl. 70, 699.


Young, J. S., Xie, S., Tacconi, L. et al.: 1995, Astrophys. J. Suppl. 98, 219.Young, L. M. and Lo, K. Y.: 1996, Astrophys. J. 464, L59.Young, L. M. and Lo, K. Y.: 1997, Astrophys. J. 476, 127.Young, P. J.: 1976, Astron. J. 81, 807.Zepf, S. E., Carter, D., Sharpless, R. M., and Ashman, K. M.: 1995, Astrophys. J. 445, L19.Zirbel, E. L.: 1996, Astrophys. J. 473, 713.Zirbel, E. L.: 1997, Astrophys. J. 476, 489.Zoabi, E., Soker, N., and Regev, O.: 1996, Astrophys. J. 460, 244.

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