arXiv:astro-ph/9702146v1 17 Feb 1997 colo hsc n pc eerh nvriyo Birming of University Research, Space and Physics of School nteOii fGoua lsesi litcladcD and Elliptical in Clusters Globular of Origin the On ioiybttema ealct ftemtlpo nsdoe ones metal–poor the t pare of metallicity with that mean well the find correlates but we minosity GCs metal–rich distributions, the metallicity of subpopulations, GC metallicity metal–poor F bimodal and with systems. metal–rich ies GC of the properties separate new we several forma identify galaxy and for or lution, implications the the explore We populations, GC metallicit galaxies. distinct GC early–type massive multimodal in the seen are tions research (GC) cluster ular niae httemtlrc C r lsl ope othe to coupled closely are GCs metal–rich the that indicates ehp h otntwrh frcn nig nextragalac in findings recent of noteworthy most the Perhaps e rplinLbrtr,40 a rv rv,Psdn C Pasadena Drive, Grove Oak 4800 Laboratory, Propulsion Jet ikOsraoy nvriyo aiona at rz CA Cruz, Santa California, of University Observatory, Lick ikOsraoy nvriyo aiona at rz CA Cruz, Santa California, of University Observatory, Lick lcrncmi:[email protected] mail: Electronic lcrncmi:[email protected] mail: Electronic lcrncmi:[email protected] mail: Electronic T,Uie Kingdom United 2TT, ucnA Forbes A. Duncan alJ Grillmair J. Carl enP Brodie P. Jean ABSTRACT Galaxies and 1 a,Egatn imnhmB15 Birmingham Edgbaston, ham, inadevo- and tion 95064 95064 91109 A tglx lu- galaxy nt gno these of igin not. s aayand galaxy distribu- y rt when irst, ngalax- in emean he i glob- tic This share a common chemical enrichment history with the galaxy field stars. The mean metallicity of the metal–poor population is largely indepen- dent of the galaxy luminosity. Second, the slope of the GC system radial surface density varies considerably in early–type galaxies. However the galaxies with relatively populous GC systems for their luminosity (called high specific frequency, SN ) only have shallow, extended radial distribu- tions. A characteristic of high SN galaxies is that their GCs are pref- erentially located in the outer galaxy regions relative to the underlying starlight. Third, we find that the ratio of metal–rich to metal–poor GCs correlates with SN . In other words, high SN galaxies have proportion- ately more metal–poor GCs, per unit galaxy light, than low SN galaxies. Fourth, we find steeper metallicity gradients in high SN galaxies. This is due to the greater number of metal–poor GCs at large galactocentric radii. We critically review current ideas for the origin of GCs in giant ellip- tical (gE) and cD galaxies and conclude that the gaseous merger model of Ashman & Zepf (1992) is unlikely to account for the GC systems in these galaxies. Tidal stripping of the GCs from nearby galaxies appears to contribute to the GC population in the outer parts of cD galaxies, but we suggest that the vast majority of GCs in gE and cD galaxies have probably formed in situ. We speculate that these indigenous GCs formed in two distinct phases of star formation from gas of differing metallicity, giving rise to the bimodal GC metallicity distributions. The metal–poor GCs are formed at an early stage in the collapse of the protogalactic cloud. The metal–rich GCs formed out of more enriched gas, roughly contemporaneously with the galaxy stars. In this sense the metal–rich GCs in elliptical galaxies are the analog of the metal–poor halo GCs in spirals. The disk GCs in spirals may represent a third phase of this formation process.
2 1. Introduction mation efficiency. When normalized in this way, SN ∼ 5 for most early–type galaxies with Globular clusters have long been recog- magnitudes –19 > MV > –22 (e.g. Kissler– nized as ‘galactic fossils’ and like their coun- Patig 1996). However, for –22 > MV > –23.5 terparts in palaeontology they provide a unique several galaxies have SN ∼ 15 i.e., a factor of tool for studying galaxy formation and evolu- three times as many GCs per starlight as ‘nor- tion at early times. In particular, they may mal’ SN galaxies. Furthermore, these galaxies provide the best probe of chemical enrich- do not simply have higher global SN values ment and star formation in the initial stages but they show evidence for high SN values at of galaxy formation. Observations of globular all galactocentric radii (McLaughlin, Harris & clusters (GCs) may enable us to distinguish Hanes 1994). Finding an explanation for the between variants of the collapse and merger ‘high’ SN GC systems has been described by models for elliptical galaxy formation, a dis- McLaughlin et al. (1994) as “the most out- tinction that has been difficult to make in the standing problem in globular cluster system past (Kormendy 1990). research”. The origin of these GC systems The first step in using GCs to probe galaxy has yet to be clearly identified and indeed the formation is to understand the formation of question of whether they are anomalous or GCs themselves. A variety of models (e.g. merely represent the extreme of some distri- Fall & Rees 1985; Murray & Lin 1992; Ku- bution has been the subject of debate. mai, Basu & Fujimoto 1993a; Harris & Pu- Environmental effects on GC systems have dritz 1994; Vietri & Pesce 1995) have sought been suspected for some time (e.g. cluster to explain the formation mechanism for GCs. ellipticals have on average higher SN values These models focus on gas cloud physics, in- than field ellipticals; Harris 1991). However cluding metallicity, cooling and star forma- it was generally believed that there were no tion rates, shocks etc. Although many of correlations between the properties of a GC the early models have been ruled out, some system and the characteristics of the cluster of the more recent models may eventually be of galaxies in which the parent galaxy resided accepted as a prescription for GC formation. (McLaughlin et al. 1994). Initial evidence In this paper we take a different approach, that such correlations did in fact exist was and attempt to determine where and when presented by West et al. (1995). More re- GCs formed by examining the systemic prop- cently, based on a larger and more homoge- erties of GC systems. This may provide con- neous sample than was previously available, straints for the further refinement of the GC Blakeslee (1996) has shown that SN is corre- and galaxy formation models. lated with cluster velocity dispersion, the lo- The number of GCs in an early–type galaxy cal density of bright galaxies and X–ray tem- scales roughly with the parent galaxy’s lu- perature. His sample, of 5 cDs and 14 gi- minosity. A useful convention is to normal- ant ellipticals (gEs), cover the whole range ize the GC counts per absolute luminosity of from SN ∼4 to ∼10. Therefore, although cD MV = –15 (Harris & van den Bergh 1981). galaxies have undoubtly had more complex The resulting ‘specific frequency’, SN = N × evolutionary histories than gEs, their GC sys- 100.4(MV +15), is probably a measure of GC for- tems may actually have had a similar origin.
3 Blakeslee concluded that bright central galax- ered or confirmed. They provide important ies with high SN values (these are usually but constraints on any model for GC formation. not always cDs) are not ‘anomalous’ in the We have taken distance moduli and the num- sense that they have too many GCs for their ber of GCs (NGC ) directly from the McMas- luminosity. Rather, these galaxies are ‘under- ter Catalog (Harris 1996). The one exception luminous’ for their clustercentric position and is NGC 1404, for which we take NGC from the number of GCs more accurately reflects the recent HST study of Forbes et al. (1997), the local cluster mass density. The fact that i.e. NGC = 620 ± 130. We believe this to there is never more than one massive high be more accurate than the number quoted by SN galaxy per cluster is consistent with these Harris (1996), i.e. 950 ± 140 (see also sec- findings. tion 3.5). Extinction–corrected V magnitudes In this paper we focus on the origin of GCs come from Faber et al. (1989) and are gener- in massive early–type galaxies (i.e. gE and cD ally within a few tenths of those given in the galaxies). These galaxies contain the most RC3. The one exception here is NGC 3311, in 0 populous GC systems; which are often the which the values are VT = 10.14 (Faber et al. 0 best–studied. Their GCs may have had a sim- 1989) and VT = 11.16 (RC3). In a photomet- ilar origin and the galaxies themselves share ric study of the Hydra cluster, Hamabe (1993) 0 some common properties e.g., the main bod- finds VT = 10.75 which is roughly the average ies of cD galaxies have a similar fundamental– of the above two values. Here we adopt the plane relation to that of elliptical galaxies Hamabe value which gives MV = –22.7. Most (Hoessel, Oegerle & Schneider 1987). First, of the galaxies discussed in this section are lu- we present several recent developments in the minous gE or cD galaxies. In this sense they study of GC systems that provide important are generally boxy, slow rotators with shallow constraints on the origin of GCs. Second, we central cusps (e.g. Faber et al. 1996; Carollo −1 discuss these results within the framework of et al. 1997). We assume H◦ = 75 km s − various models proposed to explain the prop- Mpc 1. erties of GC systems of early–type galaxies. Third, we present our views and supporting 2.1. Multimodal Globular Cluster Metal- data for what we believe are the most vi- licity Distributions able formation mechanisms (i.e. tidal strip- Although the initial evidence for multi- ping and in situ GC formation). Fourth, we modal GC color (metallicity) distributions present a case study of the cD galaxy NGC in early–type galaxies was weak, there are 1399 and the Fornax cluster of galaxies. Fi- now a handful of convincing cases (see Ta- nally, we summarize our preferred scenario for ble 1 for a summary). The existence of the origin of GCs and make several predic- multimodal GC metallicity distributions sug- tions to test these ideas. gests that GCs have formed in distinct star formation episodes from different metallicity 2. Recent Developments in Extragalac- gas; a single star formation episode would tic Globular Cluster Studies give rise to a unimodal distribution. The observed metallicity distributions effectively In this section we highlight several aspects rule out a simple monolithic collapse and pro- of GC systems that have recently been discov-
4 vide a strong constraint for any GC forma- NGC 4472 (Geisler, Lee & Kim 1996) can be tion model. The data for 11 gE/cD galaxies extended to other galaxies, then we can ex- from the literature are summarized in Table plain the long–standing observation that the 1. For NGC 1399 and NGC 4486 (M87) there mean metallicity of GCs is about 0.5 dex have been claims for a trimodal metallicity lower than field stars in gEs (e.g. Harris distribution. In these cases we will refer to 1991). This then simply arises because the the metal–poor, intermediate metallicity and metal–rich GCs have a similar metallicity to metal–rich GC populations. The difference the galaxy stars, and the metal–poor GCs are between the metallicity of the metal–rich and about 1 dex more metal–poor. With a sim- metal–poor peaks has a mean value and sta- ilar number of GCs in each subpopulation, tistical error of ∆[Fe/H] = 1.0 ± 0.1 dex. the mean offset between GCs and field stars The metallicity distributions quoted in Ta- is 0.5 dex, as is generally observed. The low- ble 1 are, strictly speaking, ones of color. est luminosity galaxy in Table 1 has MV = The bimodal colors represent GC populations –21. Less luminous ellipticals with MV > –21 that have different metallicities or ages, or also show a metallicity offset between GC and some combination of both. This is the well– galaxy stars (Forbes et al. 1996). known age–metallicity degeneracy that affects In terms of the number of GCs observed the interpretation of optical colors. If we as- with high quality CCD photometry, the best sume that the two populations have the same studied GC system is that of NGC 4472 metallicity and that the red population is very (M49) in Virgo (Geisler et al. 1996). Al- old (i.e. 15 Gyr), then for a typical color dif- though it is not a high SN galaxy (SN = 5.5 ference of ∆(V–I) ∼ 0.2, the blue population ± 1.7), it is a giant elliptical (MV = –22.7) is younger by ∼ 13 Gyr. On the other hand, that is somewhat more luminous than NGC for contemporaneous populations ∆(V–I) ∼ 4486 (M87). Several interesting results were 0.2 indicates that the red population is more found by Geisler et al. which may also hold metal–rich than the blue one by ∼ 1 dex in true for other ellipticals with bimodal metal- metallicity. The latter appears much more licity distributions. The two metallicity peaks reasonable given, for example, our knowledge in NGC 4472 are located at [Fe/H] = –1.25 of the Milky Way GC system. In the dis- and –0.05, so that ∆[Fe/H] = 1.2. They can cussion below we will assume that GC colors be well represented by two Gaussians, both mostly reflect GC metallicity but in reality with the same dispersion of 0.38 dex, with there will probably be some component due about 2/3 in the metal–poor peak and 1/3 in to age effects (see also Zepf 1996). For exam- the metal–rich peak. (As given in Table 1, ple, Chaboyer et al. (1992) have derived an for a total population of 6300 this indicates age–metallicity relation for Milky Way halo there are about 4200 metal–poor and 2100 GCs in which GCs with ∆[Fe/H] = 1 dex are metal–rich GCs in NGC 4472.) The overall separated in age by ∼ 4 Gyrs. High signal-to- radial metallicity gradient of the GC system noise spectra of a sample of blue and red GCs is almost entirely due to the changing rela- in an elliptical galaxy are needed to further tive mix of the two populations with radius. investigate this issue. There is little or no evidence for a metallic- If the characteristics of the GC system in ity gradient within either GC subpopulation
5 suggesting a largely dissipationless formation better represented by two GC mean metal- process (similar to the halo of our Galaxy). licities corresponding to the metal–rich and The mean metallicity of the metal–rich pop- metal–poor subpopulations. Although obser- ulation closely matches the underlying field vational errors may dominate, there are other star metallicity over a wide range in galac- reasons for expecting some scatter in the Z– tocentric radius. The metal–poor GCs are L relation. In particular, some of the scat- simply offset from the field stars by 1.2 dex. ter will be due to changing proportions of the The metal–rich GCs are more centrally con- metal–poor and metal–rich GC populations centrated than the metal–poor GCs, and have from one galaxy to the next, which is further ′′ ′′ effective radii of 227 and 274 respectively. complicated by observations that cover differ- The galaxy itself has an effective radius of ent galactocentric radii (metal–poor GCs are ′′ about 100 (Bender, Burstein & Faber 1992). more radially extended than the metal–rich ones). In Fig. 1 we plot the mean metallic- 2.2. The Globular Cluster Metallicity ity of the two subpopulations, listed in Table – Galaxy Luminosity Relation 1, against the galaxy luminosity. We have assumed errors of ± 0.1 for Washington color A relationship between GC mean metallic- photometry, ± 0.15 for B–I and ± 0.25 for V–I ity and parent galaxy luminosity (Z–L) was colors which reflect the scatter in the color– first suggested by van den Bergh (1975) and metallicity relations. A weighted fit for the later supported by the spectroscopic metal- metal–rich GC peak gives [Fe/H] = –0.21 (± licity data of Brodie & Huchra (1991). The 0.06) M – 4.6 (±1.4). The scatter of the existence of this relation, at least for low lu- V metal–rich data points about the relation is minosity galaxies, was disputed by Ashman less than the estimated error bars, suggest- & Bird (1993). They claimed that the GC ing that our assumed metallicity errors are metallicity was independent of galaxy lumi- overestimated. In the lower panel, for the nosity. Larger galaxy samples (Durrell et al. metal–poor GCs, we have simply offset the 1996; Forbes et al. 1996) have now confirmed metal–rich relation downward by ∆[Fe/H] = that such a relationship of the form Z ∝ L0.4, 1.0 dex (i.e. the average metallicity difference does indeed exist, over 10 magnitudes. The between the two peaks). There is consider- slope of the relation is the same as the galaxy ably more scatter in the Z–L relation for the Z–L relation (Brodie & Huchra 1991). Taken metal–poor data points. This is unlikely to together these results indicate that GCs have be caused by errors or calibration problems had a similar chemical enrichment history to and so the difference in scatter must be a real their parent galaxy. The Z–L relation holds effect. Either there is virtually no Z–L rela- important clues about GC and galaxy forma- tion for metal–poor GCs or there is a ‘second tion processes, but the large scatter, which parameter effect’ causing the significant scat- may be mostly observational error, makes it ter. Figure 1 clearly indicates that [Fe/H] for difficult to reach any definitive conclusions. the metal–rich GCs is more closely coupled In the previous studies of the GC Z–L re- to the luminosity of the galaxy than for the lation, the GC metallicity was taken to be metal–poor GCs. This result will be discussed the mean of the entire GC system. We now further in section 3.4. know that in several cases, the GC system is
6 2.3. The Spatial Distribution of Glob- at large radii in the main body (r1/4 part) of ular Clusters the galaxy. In NGC 4486 (M87) the GC radial profile shows evidence for a change in slope at It has been known for sometime that the the radius of the cD envelope (e.g. McLaugh- radial distribution of GCs is often more ex- lin et al. 1993). This suggests that the outer tended than the starlight of the parent galaxy GCs and the envelope of cD galaxies have a (see Harris 1991 for a review). The GC radial similar origin. As cD envelopes are generally surface density profile is usually fit by: thought to be a product of their cluster envi- ronment (e.g. Kormendy & Djorgovski 1989) α ρ = ρor this implies that GCs in the envelope(1) are also associated with some aspect of the cluster. where α is the radial slope. Kissler–Patig (1996) has collected measurements of α from 3. Models for the Origin of Globular the literature. After excluding the S0 galax- Clusters in gE and cD Galaxies ies, we show in Fig. 2 this slope versus the SN value as given by Kissler–Patig. This figure Fall & Rees (1988) characterized GC for- shows that low SN galaxies have a wide range mation models into three types – primary, of GC radial slopes, whereas high SN galaxies secondary and tertiary. In the primary for- only have relatively flat slopes. For SN > 7 mation models, GCs may form well before the mean slope is α = –1.43 ± 0.12 (statisti- the galaxy. In secondary formation models, cal error) and for SN ≤ 7 it is α = –1.76 ± GCs form at essentially the same time as the 0.09. There are no cases known of high SN galaxy itself, and for tertiary formation mod- galaxies with steep GC radial slopes. This els, GCs form after most of the parent galaxy ‘zone of avoidance’ is unlikely to be due to has formed. Given the numerous correlations selection effects as high SN galaxies have nu- between GCs and their parent galaxy (see merous GCs and so are relatively easy to find Harris 1991) and the absence of dark matter and study. Thus another property of high SN in GCs, formation well before the galaxy itself galaxies is that their GC systems have a very (e.g. Peebles & Dicke 1968) and GC forma- extended spatial distribution (see also Kaisler tion that is unrelated to galaxies (e.g. West et al. 1996). 1993) seems to be ruled out. Theories for the A more useful quantity to examine would formation of galaxies have traditionally been be the difference between the GC radial slope divided into the ‘collapse’ or the ‘merger’ pic- and the underlying starlight slope. Unfor- ture. The present majority view is that both tunately, the starlight slopes have only been collapse and the merging of subunits (hier- measured over similar galactocentric radii for archical clustering) play a role in the forma- a handful of galaxies. In the inner (i.e. r1/4) tion and subsequent evolution of early–type regions of cD galaxies, the GC system is gen- galaxies. The presence of more than one pop- erally flatter than the galaxy starlight (Harris ulation of GCs in massive early–type galaxies 1991; McLaughlin, Harris & Hanes 1993). In effectively rules out a simple monolithic col- lapse. Either the collapse was more complex other words, a characteristic of high SN galax- ies is that their GCs are preferentially located (e.g. episodic) or some tertiary GC process has occurred (e.g. the accretion of existing
7 GCs or the creation of new GCs in a gaseous velocity gas. Their models predict a mass– merger). metallicity relation for individual GCs which Below we discuss specific models for the is not supported by the observational data origin of GCs in giant elliptical and cD galax- (see e.g. Forbes, Brodie & Huchra 1996). Fur- ies: 1) a gaseous merger; 2) cooling flows; thermore, this model can not apply to the GC 3) intracluster globular clusters; 4) in situ systems of dwarf galaxies which would prob- GC formation and 5) tidal stripping of GCs ably be disrupted by high velocity collisions. from nearby galaxies. The merger model of Another merger model that has been the Ashman & Zepf (1992) and Zepf & Ashman focus of much attention is that of Ashman & (1993), makes detailed predictions for GC sys- Zepf (1992) and Zepf & Ashman (1993). Be- tem formation on a macroscopic level which low we discuss this model in detail in the con- are tested below. Most other models have not text of GC formation in gE and cD galaxies. made such clear predictions and can therefore In their model the bimodal GC metallicity only be discussed on a more qualitative basis. distributions are due to the presence of metal– poor GCs from the progenitor spirals and 3.1. The Gaseous Merger Model metal–rich ones created from the gas of the merged galaxies. The metal–rich GCs are ob- Van den Bergh’s (1984) argued against the served to be more centrally concentrated than merger hypothesis on the grounds that ellip- the metal–poor ones, as expected in their ticals have many more GCs per unit light model. One of the most noteworthy successes than spirals. An additional problem is that of this model is the discovery of protoglobu- a purely stellar merger of two ellipticals (in lar clusters in currently merging galaxies (e.g. which no new GCs are created) cannot ex- Whitmore et al. 1993; Schweizer et al. 1996). plain the ratios of metal–poor to metal–rich The sizes, magnitudes, colors and mass distri- GCs. In particular, the GC metallicity – bution of the identified objects generally ap- galaxy luminosity relation suggests that the pear to be consistent with the properties of metal–poor GCs would come from the lower protoglobular clusters. Although some issues luminosity progenitor and would therefore also remain (such as destruction of the very low be fewer in number. Column 9 of Table 1 mass objects) it seems plausible that these ob- shows that this is generally not the case, i.e. jects will evolve into bona fide GCs. Thus the the number of metal–poor GCs exceeds the Ashman & Zepf (1992) model appears consis- number of metal–rich GCs. The GC metallic- tent with the observations of GCs in currently ity – galaxy luminosity relation also suggests merging galaxies. that metal–rich (luminous) ellipticals are not made by simply merging several metal–poor Here we are interested in the origin of GCs ellipticals. These issues and others have been in the most massive early–type galaxies. In discussed by van den Bergh (1990). the framework of the Ashman & Zepf (1992) hypothesis the gEs have been formed by mas- Schweizer (1987) has suggested that new sive gas–rich galaxies or several ∼ L∗ galax- GCs are created from the gas of the merger ies. For cD galaxies they advocate multiple progenitors. For a gaseous merger, the mod- mergers which in turn leads to several “differ- els of Kumai, Basu & Fujimoto (1993a,b) pro- ent epochs of [globular] cluster formation and pose that GCs form in the collision of high
8 thus a spread in cluster ages”. Can a gaseous GCs are formed. Instead we find that it de- merger and subsequent GC formation, as de- creases, as shown in Fig. 3. Thus a property scribed by Ashman & Zepf (1992), account of high SN galaxies is that they have relatively for the properties of GC systems in gE and more metal–poor GCs, in direct contradiction cD galaxies ? to that expected in the Ashman & Zepf (1992)
Previous studies of high SN galaxies (McLaugh- model. For the cD galaxies in Table 1 the lin et al. 1994) and brightest cluster galaxies number of metal–poor GCs significantly ex- (Blakeslee 1996) have mentioned some prob- ceed the metal–rich ones. lems in trying to account for the populous GC systems with the Ashman & Zepf (1992) • The radial metallicity slope increases merger model. The case of NGC 3311 pro- with SN . Zepf & Ashman (1993) predict that vides a particularly strong constraint for any the “colour [metallicity] gradients of ellipti- GC formation model as its GCs are relatively cals with high SN values will be shallower metal–rich with peaks at [Fe/H] = –0.5 and than those with normal SN values”. This +0.15 (Secker et al. 1995). It is hard to see would be so because in a high SN galaxy, how the GCs in a spiral galaxy (which typi- large numbers of metal–rich GCs should be cally have mean [Fe/H] = –1.5) can account formed near the center and the metal–poor for the ‘metal–poor’ peak of [Fe/H] ∼ –0.5. GCs would be ‘pufffed up’ to large galacto- Here we detail some of the observational data centric radii (see Zepf & Ashman 1993). We which are contrary to the expectations of the have collected metallicity gradients and SN Ashman & Zepf (1992) merger model, in de- values from the literature and listed them in creasing order of significance. Table 2. In Fig. 4 we plot the [Fe/H] metal- licity gradients versus SN . Although there is • The metal–poor GCs are more nu- considerable scatter, the trend is for high SN merous than the metal–rich population. ellipticals to have steeper metallicity slopes In order to explain high SN galaxies by the than normal SN ellipticals. This trend is in merger of several normal SN galaxies, the ma- direct contradiction to the Zepf & Ashman jority of GCs must be created in the merger (1993) prediction and appears to be a gen- event. This implies that the metal–rich GCs eral problem for the merger model. Geisler should be more numerous (by a factor of et al. (1996) found that the merger model ∼ 3×) than the metal–poor GCs. This is prediction failed to match the observed GC not the case. In Table 1 we give the ratio metallicity in both absolute value and radial of the number of metal–rich to metal–poor variation for NGC 4472. GCs (this is equivalent to ∆S = Nnew/Nold as We have made a crude prediction for the defined by Ashman & Zepf 1992). This ra- variation of metallicity slope with increasing tio is expected to vary from ∼ 1 for normal SN based on the data for NGC 4472 (Geisler SN ellipticals to 4 for high SN galaxies. Table et al. 1996). As appears to be the case for 1 shows that this ratio (column 9) is gener- NGC 4472 (M49), we will assume that the ally less than 1. Furthermore, in the merger observed GC metallicity gradients in ellipti- model this ratio should increase with increas- cals are due solely to the changing relative ing SN as larger numbers of new (metal–rich) mix of metal–rich and metal–poor GCs. Fur-
9 thermore, we will assume that the only differ- ent GC mean metallicities and the gas (from ence between a galaxy of SN = 5.5 (i.e. NGC which new GCs form) should be enriched by 4472) and one with SN = 11 is an increase varying amounts. Therefore, unless all of the in the number of metal–poor GCs (i.e. the merging galaxies are similar, we would ex- number of metal–rich GCs remains constant). pect a broad, top–hat type distribution cov- Thus we can predict the metallicity slope of a ering a range of metallicity, but instead dis- galaxy with SN = 11. If these additional GCs tinct metallicity peaks are seen. In the case are distributed with the same surface density of NGC 4472 (although not a cD it has a sim- distribution as is currently seen, then we ex- ilar luminosity to a cD galaxy) the metal–rich pect a metallicity gradient of ∆[Fe/H]/∆logR and metal–poor peaks have a similar disper- ′′ = –0.58 dex/ . The dashed line shown in sion when fitted with a Gaussian. Fig. 4 simply connects the data point for NGC 4472 (i.e SN = 5.5 and ∆[Fe/H]/∆logR • Galaxies with the same luminosity = –0.41) to that expected for an SN = 11 should have similar SN values. If galax- galaxy with a metallicity gradient of –0.58. ies are made from multiple mergers, why do It reproduces the general trend surprisingly some apparently create GCs much more effi- well. Note that the Ashman & Zepf (1992) ciently than others ? For example, NGC 4486 predicted relation would have a gradient of (M87) and NGC 4472 (M49) both lie in the the opposite sign to this line. Virgo cluster and have essentially the same luminosity, and yet they have quite different • High SN galaxies do not have peaky SN values (13 and 5.5 respectively). In the central GC densities. The gas in a merger merger hypothesis, this would indicate that quickly dissipates to the galaxy center, where the GCs in NGC 4486 formed ∼ 3× more ef- large numbers of new GCs should form. The ficiently than those in NGC 4472. expected higher GC surface densities in the centers of high SN galaxies are not seen. In • The GC surface density profiles of fact the surface density flattens off in log normal SN galaxies are not always flat- space giving a GC system ‘core’ (e.g. Grill- ter than the underlying starlight. Zepf mair, Pritchet & van den Bergh 1986; Lauer & Ashman (1993) “expect the globular clus- & Kormendy 1986; Grillmair et al. 1994a). ter system to be noticeably flatter than the The size of this core is larger in more lumi- galaxy” for normal (i.e. not high SN ) galaxies. nous galaxies and this does not appear to be As shown in Fig. 2, galaxies with SN ≤ 7 have due to destruction effects over time (Forbes a wide range in GC surface density slopes sug- et al. 1996). gesting a wide range in the difference between GC and starlight slopes. Kissler–Patig et al. • Multiple mergers will not produce dis- (1996) have studied several galaxies in the tinct metallicity peaks. Based solely in lu- Fornax cluster with SN < 7. All have GC minosity terms, roughly 10 L∗ galaxies would slopes that are consistent with the starlight be required to form a cD galaxy or the more slope – none are noticeably flatter contrary massive ellipticals, from multiple mergers. In to the prediction of Zepf & Ashman (1993). general these galaxies will have slightly differ-
10 • The main body and envelope of cD = 4.6 ± 0.7. Whitmore et al. (1997) show galaxies appear to be distinct. For high that in the best–studied merging galaxies the SN galaxies, Zepf & Ashman (1993) predict in number of new GCs created is equal to or less general that “the properties of the galaxy and than the number of original GCs from the the globular cluster system should be simi- progenitor spirals. This means that the SN lar”. Both the main body and the cD en- values, after 15 Gyrs, are ∼ 2–3. Globular velope should be formed by the same merg- clusters need to form with much higher effi- ers. In practice, the envelopes of cD galax- ciency in a gaseous merger if the SN values of ies are generally found to be structurally dis- merger remnants are to match those of ellip- tinct from the rest of the galaxy (e.g. Mackie tical galaxies. So the argument that there are 1992). Furthermore, in both NGC 1399 and too many GCs in ellipticals to be explained NGC 4486 (M87) the velocity dispersion of by mergers (van den Bergh 1984) may still the GCs is well in excess of the stellar veloc- be valid. If the GC systems of low luminos- ity dispersion at large radii (Grillmair et al. ity ellipticals are not the result of a merger, 1994b; Huchra & Brodie 1987; Mould et al. then currently merging galaxies must form a 1990); for NGC 1399 the GC velocity disper- galaxy with characteristics different from to- sion in the cD envelope is close to that for the day’s ellipticals (see also van den Bergh 1995). galaxies in the cluster. 3.2. Cooling Flows To summarize, the gaseous merger model It has been suggested that GCs could form as proposed by Ashman & Zepf (1992) for in the material that cools and condenses from GC formation in massive early–type galaxies cluster–wide hot gas (Fabian et al. 1984). faces some serious difficulties. We note that Richer et al. (1993) proposed that the pro- some of these difficulties have been identified toglobular clusters seen in NGC 1275 were the −1 by Ashman & Zepf (1997) and possible solu- result of the M˙ ∼ 300M⊙yr cooling flow. tions suggested. However it is our view that Subsequently, protoglobulars have been iden- the disagreements between the observational tified in several galaxies without cooling flows data and the expectations from the Ashman (e.g. Whitmore et al. 1993, 1997; Schweizer & Zepf merger model are still substantial. In et al. 1996). Bridges et al. (1996a) have particular, the model cannot explain the large discussed cooling flow GC formation and con- numbers of excess GCs present in high SN cD cluded that “it seems most unlikely that cool- galaxies and appears to be in conflict with GC ing flows are responsible for the high SN phe- trends in normal SN giant ellipticals. nomenon”. Similar views were put forth by An open question is whether the Ashman Grillmair et al. (1994b) and Harris, Pritchet & Zepf (1992) model can explain the origin & McClure (1995). We refer the reader to of GCs in low luminosity ellipticals. It is these works for further details. currently not clear whether such galaxies are smaller versions of gEs or if they form a dis- 3.3. Intracluster Globular Clusters tinct population (e.g. Carollo et al. 1997). West et al. (1995) proposed that there Using the list of Kissler–Patig (1996), we find exists a population of GCs that are associ- ellipticals with –19.5 > M > –21.0 have S V N ated with the potential well of the cluster
11 of galaxies rather than with any individual merging galaxy NGC 3921, Schweizer et al. galaxy. A trend for the number of GCs to (1996) concluded that GMCs are associated correlate with X–ray temperature (once cor- with the formation of protoglobular clusters. rected for the distance of the galaxy from the The largest GMCs may have similar masses dynamical center of the cluster) was taken as and densities to the primordial fragments pro- evidence in favor of their intracluster globu- posed by Searle & Zinn (1978) as the building lar cluster (IGC) hypothesis. They suggested blocks of galaxies (see also Harris & Pudritz three possible sources for these IGCs – cooling 1994). In this case a dwarf elliptical may rep- flows, local density enhancements and tidal resent a few Searle–Zinn fragments whereas stripping. Cooling flow GC formation has large galaxies are made up of numerous such been discussed above. Even if a mechanism fragments. could create GCs in local density enhance- For cD galaxies, an argument in favor of ments outside of galaxies, it would not ac- in situ formation is that with a GC formation count for the known correlations of GCs with efficiency of ≤1% (Ashman & Zepf 1992; Lar- 11 their parent galaxy (e.g. the GC metallicity– son 1993) about 10 M⊙ of gas would be re- galaxy luminosity relation). Tidally stripped quired to form the ‘excess’ ∼10,000 GCs seen 5 GCs may be more viable, especially in light (assuming each GC has a mass of 10 M⊙). of the recent discoveries of intergalactic plan- Thus apparently most of the GCs in cD galax- etary nebulae (Theuns & Warren 1997). This ies were formed at an early epoch when the is however unlikely to be the source of the en- galaxy itself was in a mostly gaseous phase. tire ∼ 10,000 excess GCs required in high SN In early–type galaxies, the observations that galaxies and metallicity is still an important GCs are more metal–poor in the mean than constraint. We further discuss tidal stripping the underlying starlight at a given radius and below. that the GC systems tend to be more spatially extended than the starlight suggest that the 3.4. In Situ Formation GCs formed before the bulk of the galaxy field Harris and co–workers (e.g. Harris 1991; stars (Harris 1991). The various known corre- McLaughlin et al. 1994; Harris, Pritchet & lations between the GC system and the par- McClure 1995; Durrell et al. 1996) advo- ent galaxy suggest that the epoch of this ‘pre– cate in situ formation for most of the GCs galaxy’ GC formation occurred just before the in gE and cD galaxies, with possibly a small main collapse phase. But we now know, at contribution from other sources. Durrell et least for the GC metallicity – galaxy lumi- 5 al. (1996) found that GCs with M ≥ 10 M⊙ nosity relation (section 2.2), that the metal– have a power–law mass distribution of the poor GCs are poorly correlated, if at all, with same form, which is remarkably independent galaxy luminosity. This relaxes the constraint of galaxy type, luminosity and environment. on the epoch of formation for these GCs. A This suggests a similar GC formation mecha- further modification to the standard collapse nism for all galaxies, including dwarfs, giants picture is required to explain the presence of and cD galaxies. The power–law is identical two distinct GCs populations (as indicated by to that for giant molecular clouds (GMCs) the now well–established bimodal color distri- in our Galaxy. In a study of the currently butions) which rules out a simple monolithic
12 collapse. Below we discuss our view for GC GCs may depend on local density. In general formation in a multiphase collapse. the star formation time scale varies as ∝ ρ−1, At early times, low–metallicity pre–galaxy whereas the dynamical collapse time varies as − GCs and some field stars will form, enriching ∝ ρ 1/2. So if the local ambient density is the protogalactic gas. If only a small fraction high, a larger fraction of the initial material of the total gas is used in this initial phase, turns into GCs leaving less gas for the second then the remaining metal–enriched gas should episode of GC formation. High SN galaxies undergo further collapse. A second episode could simply be those with more locally dense of star formation is required to produce the and/or clumpy gas. Another possibility, is metal–rich GCs. In this ‘galaxy’ phase the that the number of pre–galaxy GCs depends vast majority of field stars also form. Such on the total mass of the protogalactic cloud. episodic formation in different metallicity gas However we do not find a strong trend for the could be expected to give rise to bimodal GC number of metal–poor GCs (see Table 1) to metallicity distributions. We would expect scale with galaxy luminosity, suggesting that the pre–galaxy GCs to have a nearly spher- the mass of the protocloud is not a major fac- ical distribution, show little or no rotation tor. and have a large velocity dispersion. There Perhaps the main problem with this col- is some evidence to support this prediction. lapse picture is the lack of a well–understood Grillmair et al. (1994b) found that a sam- mechanism for initiating GC (and star) for- ple of GCs in the outer regions of NGC 1399 mation in the pre–galaxy phase, turning it off (which would be dominated by the metal– and turning it on again at some later stage poor ones) had a much higher velocity dis- in the collapse. Recently, Burkert & Ruiz– persion than the stars in the main body of Lapuente (1997) have proposed a model in NGC 1399 and with no measurable rotation. which, after the first star formation phase, In NGC 4486, GCs also have higher velocity subsequent star formation may be stalled by dispersions than the galaxy stars in the outer several Gyrs. They propose that the hot gas regions (Huchra & Brodie 1987; Mould et al. initially heated by SNe type II in the first 1990). The metal–rich GCs formed in the sec- phase, will only partially cool before being re– ond (galaxy) phase may show some rotation heated by SNe type Ia. This re–heating delays depending on the degree of dissipation during the second star formation phase. Their model the collapse and the presence of torques. was applied to small galaxies; it would be very We also expect the pre–galaxy GCs to have interesting if it could be extended to include a high SN value associated with them as very large galaxies. It is interesting that the metal- few field stars form in the initial phase. Thus licity difference between the two GC popula- high SN galaxies may simply be those with tions is about 1 dex in [Fe/H] for most large a higher fraction of pre–galaxy (metal–poor) galaxies. If confirmed by larger samples, this GCs. Some evidence to support this idea is would suggest some underlying uniformity in shown in Fig. 3 which shows that high SN the formation process and the length of the galaxies have relatively more metal–poor GCs dormant phase. If we crudely assume that all than metal–rich ones. The physical mecha- large galaxies have metal–poor GCs with a nism for regulating the number of pre–galaxy mean of [Fe/H] = –1.2 and metal–rich ones
13 with a a mean of [Fe/H] = –0.2, then the GCs is increased, the SN value is unaffected age–metallicity relation for Milky Way halo (van den Bergh 1984; Harris 1991). How- GCs (Chaboyer et al. 1992) indicates that ever there is an important point to be made. the metal–rich GCs formed ∼ 4 Gyrs after The GCs in at least some galaxies tend to be the metal–poor ones. Some evidence for early more extended radially than the underlying star formation followed by an extended dor- starlight so that the local SN value increases mant phase comes from a recent study of the with galactocentric radius (e.g. Forbes et al. Local Group Carina galaxy by Smecker–Hane 1996). In the case of NGC 4472, the global et al. (1996). In this case the first epoch of SN is 5.5 whereas the local SN exceeds 30 star formation occurred about 12 Gyrs ago at 90 kpc (McLaughlin et al. 1994). Thus and the second phase ∼ 5 Gyrs later. GCs tidally stripped from the outer parts of In the Milky Way, and by inference spiral a galaxy may have a high SN value, similar galaxies in general, there are two distinct pop- to that of the cD envelope of M87 (SN ∼ 25). ulations of GCs associated with the disk and Tidal effects would, in their weakest form, re- halo components. With reference to the col- move only the outermost GCs and if strong lapse picture above, the halo GCs have kine- enough, the whole galaxy could be accreted matics and metallicities similar to those of by its more massive neighbor. Fully–accreted halo stars (Carney 1993). This indicates that galaxies will have their orbital energy con- halo stars and GCs were formed in the halo verted into internal heat, which acts to ‘puff– at much the same time, and may therefore up’ the outer envelope of the primary galaxy be the spiral galaxy analog to the metal–rich (Hausman & Ostriker 1978). GC population in ellipticals. We speculate In Table 3 we list three nearby galaxies that the disk GCs in spirals represent a third that may have undergone a tidal interaction stage of the collapse process (which is absent with their more massive neighbors. They are: in gE galaxies). NGC 1404 (E1) located near NGC 1399 in Fornax; NGC 4486B (cE0) near NGC 4486 in 3.5. Tidal Stripping Virgo and NGC 5846A (cE2) near NGC 5846 in a compact group. Faber (1973) suggested The transfer of GCs from small galaxies that the two compact ellipticals were actually to a high S galaxy by tidal stripping has N tightly bound cores of tidally stripped galax- been suggested by several authors and dis- ies. This idea is supported by the simula- cussed by van den Bergh (1990). For exam- tions of Aguliar & White (1986) who show ple, Forte et al. (1982) proposed that NGC that tidal interaction results in a more com- 4486 (M87) had captured a significant num- pact secondary galaxy. As the central veloc- ber of GCs from other Virgo cluster galax- ity dispersion is largely unaffected by tidal ies. The exchange of GCs between galaxies stripping we can estimate the original lumi- has been modeled in a simple way by Muzzio nosity of the companion galaxy from the L– and collaborators (e.g. Muzzio et al. 1984; σ scaling relation of normal ellipticals (e.g. Muzzio 1987). The main argument put for- Faber et al. 1996). Furthermore using the GC ward against this hypothesis is that GCs and metallicity–luminosity relation from Forbes et starlight will be removed in the same propor- al. (1996) we can estimate the mean metal- tions so that, although the total number of
14 licity of the tidally stripped GCs. The pre- the HST data offers a more reliable estimate dicted original luminosity of the galaxy and of the number of GCs and we will adopt this GC metallicity, along with estimates of the value (as given in Table 1) which is close to uncertainty, are given in Table 3. The pre- the weighted average of the two ground–based dicted luminosity, suggests that NGC 4486B studies. If we assume that NGC 1404 once has lost about ∼95% of its stars and NGC had a ‘normal’ GC population with SN = 5 5846A ∼70%. On the other hand, NGC 1404 ± 1, then we can estimate the number of miss- has an observed luminosity that is consistent ing GCs. Thus: with its velocity dispersion, suggesting that −0.4(MV +15) if tidal stripping has occurred it has removed Nmissing =5×10 −Nnow (2) very little galaxy starlight (but it may still have stripped off outer GCs). For MV = –21.08, we estimate that NGC 1404 If some fraction of the GC system of the had an original GC population of 1350 ± 270. companion galaxies has been removed via So NGC 1404 may have lost ∼ 730 GCs with tidal stripping, then perhaps these GCs can a mean metallicity of [Fe/H] = –0.7 ± 0.3. be identified by their metallicity in the more On the otherhand, NGC 1399 reveals an in- massive galaxy. Figure 5 shows the metallic- termediate metallicity peak at [Fe/H] ∼ –0.8 ity distribution of GCs in NGC 1399, 4486 (Ostrov et al. 1993), as shown in Fig. 5. The and 5846. We also show the predicted metal- number of GCs with this intermediate metal- licities of the acquired GCs (as listed in Table licity is estimated to be 40% of the total num- 3) from the companion galaxies, i.e. NGC ber of GCs (see Table 1), i.e. about 2100 ± 1404, NGC 4486B and NGC 5846A respec- 700. It is plausible that a large fraction of tively. the intermediate metallicity population GCs For NGC 1404 the predicted GC mean in NGC 1399 have been tidally stripped from metallicity is [Fe/H] = –0.7 ± 0.3. Estimates NGC 1404. This possibility is explored fur- of the number of GCs in NGC 1404 have been ther in section 4 below. made by Richtler et al. (1992) and Hanes The GC metallicity distribution in NGC & Harris (1986) from ground–based studies. 4486, from Washington photometry, is claimed For a distance modulus of 31.0, they found to have three peaks at [Fe/H] = –1.35,–0.8 NGC = 880 ± 120 (SN = 3.2 ± 0.4) and NGC and –0.1 (Lee & Geisler 1993). However, we = 190 ± 80 (SN = 0.7 ± 0.3) respectively. note that central (Whitmore et al. 1995) and Most recently Forbes et al. (1997), using HST slightly off–center (Elson & Santiago 1996) data, calculated the number to be NGC = 620 HST images reveal only the metal–poor and ± 130 (SN = 2.3 ± 0.4). The Forbes et al. metal–rich peaks from V–I colors. Geisler et (1997) data had the advantages over ground– al. (1996) discussed this possible contradic- based studies of very low contamination rates tion and suggested that it could be due to the from foreground stars and background galax- different radial samplings of the three stud- ies, a background field and a pointing ∼ 10 ies. Whatever the explanation, the presence arcmin from NGC 1399 on the opposite side of an intermediate metallicity population is to NGC 1404 (allowing subtraction of GCs not crucial to our analysis, only that there associated with NGC 1399). We believe that exist some fraction of GCs with [Fe/H] ∼ –
15 0.8, which in turn could be those stripped 4. A Case Study: The Fornax Cluster from NGC 4486B. The predicted mean metal- licity of the stripped GCs from NGC 4486B The Fornax cluster of galaxies is a good is [Fe/H] = –0.7 ± 0.3. The total number of region in which to search for the effects of GCs in the NGC 4486 system is 13,000 ± 500 tidal stripping. The cluster itself is relatively (Harris 1996). We estimate from the data of nearby at 16 Mpc and well–studied. Further- Lee & Geisler (1993) that the fraction of GCs more it is more compact and perhaps dy- at this intermediate metallicity is about 40% namically older than the Virgo cluster. A sketch of the inner regions of the Fornax clus- or 5200 ± 200. For a normal SN value of 5 ± 1 ter is shown in Fig. 6. The cD galaxy, NGC and MV = –21.0 we would expect NGC 4486B 1399, is centered in the cluster X–ray emis- to have an original population of 1250 ± 250 ′ GCs. Thus NGC 4486B has not contributed sion which extends out to at least 38 or 175 significantly to the overall GC system in NGC kpc (Ikebe et al. 1996). The nearby ellipti- 4486 but may represent some fraction of the cal NGC 1404 lies within the X–ray envelope intermediate metallicity GCs. Other galaxies of NGC 1399 and appears to have an abnor- may have also contributed GCs to NGC 4486. mally low SN value of 2.3 ± 0.4. In section 4.5 we noted that NGC 1404 may have been The GC metallicity distribution in NGC tidally stripped of GCs, which have since been 5846 has two broad peaks, the metal–poor acquired by NGC 1399 which has S = 13 ± peak ranging from –1.3 < [Fe/H] < –0.9 N 4. (Forbes, Brodie & Huchra 1997). The metal- licity of NGC 5846A GCs, predicted from its Following the arguments of Faber (1973) original luminosity, is [Fe/H] = –0.8 ± 0.3. we can make a crude estimate of the tidal The total number of GCs in NGC 5846 is es- radius in NGC 1404, i.e. the radius beyond timated to be 3120 ± 1850 (Harris 1996). We which material has been tidally stripped by note that Forbes, Brodie & Huchra (1997) es- NGC 1399. The tidal radius R in kpc is given timate a total of 4670 ± 1270 from HST imag- by: ing. Using the Harris value, there are about −1/3 780 ± 463 metal–poor GCs. This compares R = D(3.5M1/M2) (3) with 800 ± 150 GCs as the estimate of the original GC population in NGC 5846A. Thus where D is the distance of closest approach, tidally stripped GCs from NGC 5846A may and the masses of NGC 1399 and NGC 1404 have contributed to the metal–poor peak in are denoted by M1 and M2 respectively. We NGC 5846. will assume that D is equal to the current ′ Blakeslee (1996) found a correlation of SN projected separation of 9.7 or 45 kpc (note with local galaxy density for cluster galax- that the orbital pericenter may be somewhat ies, and noted that this can be understood in smaller, and that the current physical separa- the context of tidal stripping. As the cross- tion is quite probably larger than 45 kpc). For ing time in high velocity dispersion clusters is MV = –21.45 and –21.08, LV = 3.2 and 2.3 × 10 shorter, individual galaxies can pass close to 10 L⊙ for NGC 1399 and NGC 1404 respec- the cluster center many times, increasing the tively. If both galaxies have the same M/L occurrence of interactions accordingly. ratio then the tidal radius R = 15 kpc. There
16 is evidence from X–ray observations and GC tem of NGC 1399 ? The GC systems of sev- kinematics that M/L ∼ 50–100 M⊙/L⊙ for eral Fornax cluster galaxies have been studied NGC 1399 (Grillmair et al. 1994b). If NGC recently by Kissler–Patig et al. (1996). Us- 1399 has a M/L some 10–20 times that of ing (m–M) = 31.0, V band magnitudes from NGC 1404, then the tidal radius becomes R Faber et al. (1989) and NGC from Kissler– ∼ 10 kpc. Richtler et al. (1992) and Forbes Patig et al. we list their SN values, lumi- et al. (1997) estimate that the GC system nosities, predicted mean GC metallicities and of NGC 1404 terminates at a galactocentric projected distances from NGC 1399 in Ta- distance of about 18 kpc. This radius is simi- ble 4. Table 4 seems to indicate that SN lar to that estimated above and suggests that increases with distance from the cluster cen- tidal stripping of the outer GCs is a plausi- ter, although the formal errors are consistent ble mechanism to explain the absence of GCs with a constant SN value. Nevertheless the beyond 18 kpc. trend with projected distance may indicate Globular cluster velocity dispersions also that several Fornax galaxies have undergone appear to match the expectations from the some tidal stripping. tidal stripping hypothesis. Grillmair et al. Further evidence for tidal stripping comes (1994b) found that a sample of 50 GCs in from a comparison of the GC surface density NGC 1399 (with a mean galactocentric ra- slope and that of the underlying starlight. Ini- dius ∼ 40 kpc) had a velocity dispersion of tially a GC system may be more extended nearly 400 km s−1. This value is much higher than the starlight with a flatter slope, tidal than the stellar velocity dispersion measured stripping will tend to remove the outermost in NGC 1399 at somewhat smaller radii, but GCs so that the GC system becomes trun- is similar to the velocity dispersion measured cated and the GC surface density slope may for galaxies of the Fornax cluster. The sim- approach that of the stellar profile. Fig. 2 ilarities in velocity dispersions of GCs and shows that low to normal SN galaxies have a Fornax galaxies is naturally consistent with a wide range of GC radial slopes and presum- tidal stripping scenario within the dark mat- ably a wide range in the difference between ter halo of NGC 1399, since the material GC and starlight radial profiles. Kissler– stripped from a passing or infalling galaxy Patig et al. (1996) have measured the ra- will retain most or all of the specific angu- dial slopes for NGC 1399 and four other For- lar momentum (with respect to NGC 1399) nax galaxies. The latter galaxies (NGC 1427, of its parent galaxy. Arnaboldi et al. (1994) 1374, 1379 and 1387) have GC slopes that are find a remarkably similar velocity dispersion the same as the underlying starlight within among the planetary nebulae at distances ∼ the errors, whereas NGC 1399, has a GC slope 25 kpc, increasing confidence in the GC re- that is noticeably flatter than the starlight. sult and lending support to the idea that tidal In the context of the above discussion, stripping contributes to both the GC system there are many similarities between the For- and stellar envelope of cD galaxies. nax cluster and the Virgo cluster. In Virgo, Have any other galaxies, besides NGC 1404, the central cD (NGC 4486) is also embedded in the Fornax cluster been tidally stripped of within an extensive X–ray envelope. The GCs GCs which may also contribute to the GC sys- in the outer regions have a higher mean veloc-
17 ity dispersion than the galaxy stars (Huchra tional data for GCs in the massive early–type & Brodie 1987; Mould et al. 1990). However, galaxies. We find that the bimodal GC metal- there are also some noteable differences be- licity distributions do not match the merger tween NGC 4486 (M87) and NGC 1399. NGC model expectations when examined in detail. 4486 has far fewer near neighboring galaxies In particular, the GC systems of high SN and a much smaller cD envelope than NGC galaxies do not have more metal–rich GCs as 1399. This combined with the higher velocity would be expected in the merger model but in- dispersion and dynamical youth of the Virgo stead have more metal–poor ones. This means cluster will lead to differences in the number that GC metallicity gradients are steeper in of accreted galaxies (and their GCs) between high SN galaxies, contrary to the predictions the two cD galaxies. of Ashman & Zepf (1992). We favor an early collapse picture for galaxy 5. Summary, Speculations and Predic- formation (see for example Eggen, Lynden– tions Bell & Sandage 1962; Larson 1975; Gott 1977; Sandage 1990) that probably involves After reviewing the current ideas on GCs some chaotic merging of many small sub- in massive early–type galaxies, we have come units (Searle & Zinn 1978; Katz 1992). We to the conclusion that most GCs were formed speculate that there are three main phases in situ in two star formation episodes. in the collapse process (which we call pre– In cD galaxies there may be an additional galaxy, galaxy and disk phases) each of which contribution from GCs that have been tidally may have associated GC formation. First, in stripped from neighboring galaxies. This ac- the early stages of the collapse, there occurs creted population should have a high local rapid star formation. Only a small fraction SN value and may have played an important of the available gas is converted into stars, role in the development of the cD envelope. and most of these stars are located in GCs In the cases of NGC 1399 and M87 we have which formed from the chaotic merging of identified an intermediate metallicity popula- dense, gravitationally bound clumps. These tion of GCs that may have been acquired via GCs have a high velocity dispersion, are dis- tidal stripping. To a much lesser extent all tributed throughout the cloud volume and are galaxies are accreting small galaxies and their metal–poor. The ratio of stars in GCs to field GCs. We discuss the Fornax cluster in some stars is large (i.e. the specific frequency is detail and suggest that several galaxies near high). There will be little or no radial abun- the central cD (NGC 1399) may have been dance gradient in the GC system from this tidally stripped of GCs. Evidence for this in- phase. Both the field stars and GCs provide cludes currently low SN values, steep GC ra- enrichment for the remaining gas. In the sec- dial distributions that match the starlight and ond phase, after the gas cloud has undergone the high velocity dispersion in the outer NGC further collapse, field star formation is pre- 1399 GCs. ferred over star formation in GCs (possibly The gaseous merger model for GC forma- due to more efficient cooling at higher metal- tion (Ashman & Zepf 1992), although gener- licities). Both field stars and GCs will have ally consistent with currently merging galax- essentially the same metallicity (i.e. relatively ies, does not adequately explain the observa-
18 metal–rich). They will have some rotation de- a galaxy cluster, a cD galaxy is formed. pending on the degree of dissipation in the We conclude that episodic in situ GC for- collapse. Unconsumed gas will eventually set- mation and a contribution from accretion/tidal tle into a disk–like structure near the galaxy stripping in the outer parts of cD galaxies may center. Thus the small disks (e.g. Scorza & offer the ‘best bet’ for solving the “most out- Bender 1996) and peaky stellar distributions standing problem in globular cluster system (e.g. Forbes, Franx & Illingworth 1995) seen research” (McLaughlin et al. 1994). This ori- in low luminosity ellipticals may indicate a gin for GCs suggests several testable predic- prolonged dissipative collapse in the less mas- tions. They include: sive systems. The variation of galaxy proper- ties from large to small ellipticals may reflect • We expect the metal–poor GCs in early– a sequence of decreasing gas-to-star conver- type galaxies to form non–rotating, hot sys- sion efficiency (see also Bender, Burstein & tems (high velocity dispersion) if their origin Faber 1992). We note that the main diffi- is either from the pre–galaxy phase or from culty with this scenario is the lack of detailed tidal stripping. The outer (metal–poor) GCs mechanism for creating two distinct phases of in both NGC 1399 and NGC 4486 (M87) have GC formation from a single halo collapse. high velocity dispersions. Future large sam- Spiral galaxies may represent the extreme ples, will hopefully be able to determine GC of this process with very inefficient conver- kinematics as a function of metallicity. The sion of gas into stars. Spirals will have virtu- metal–poor GCs will be near spherically dis- ally no pre–galaxy star formation, so that the tributed. enrichment levels of the second phase of col- lapse are reduced leading to GCs that are in • We expect the metal–rich GCs in early–type ∼ general more metal–poor (i.e. [Fe/H] –1.5) galaxies to be slowly rotating systems, with than the second phase GCs in ellipticals. Al- more rotation seen in low luminosity ellipti- though they have quite different metallicities, cals than the giants. Some tentative evidence halo GCs in spirals appear to be the analog for GC rotation comes from Hui et al. (1995) to the metal–rich GCs in ellipticals. Spirals who noted that the metal–rich GCs in NGC then go on to form a prominent disk (and as- 5128 appeared to share the rotation proper- sociated GCs) in the third phase of the col- ties of the system of planetary nebulae. The lapse. The type of the final galaxy produced metal–rich GCs will closely match the ellip- by the above process will depend largely on ticity of the underlying galaxy. the initial conditions of the the protogalac- tic gas cloud. For example, protoellipticals • We have found that the metallicity of the may be more dense, more clumpy and more metal–rich GC population is more closely cou- chaotic than protospirals. As the ambient pled to parent galaxy luminosity than the density and clumpiness of the protoelliptical metal–poor population. This suggests that increases, more metal–poor GCs with a high the GC mean metallicity – galaxy luminosity S value are produced. At the extreme end N relation is driven by the metal–rich popula- of this process, and when the protoelliptical tion, and galaxies with a significant metal– is located close to the gravitational center of poor population (i.e. the high SN galaxies)
19 will lie below the mean relation, to lower metallicity, for a given luminosity. Acknowledgments We thank R. Elson, W. Harris, M. Kissler– • As more bimodal GC systems are discov- Patig, B. Whitmore, A. Zabludoff and S. Zepf ered in ellipticals, we predict that they will for useful discussions and comments. We also follow the current trends of proportionately thank the referee, S. van den Bergh, for his more metal–poor GCs and steepening radial careful reading and suggestions for improv- metallicity gradients with increasing specific ing the paper. This research was funded by frequency. the HST grant GO-05920.01-94A and GO- 05990.01-94A
• We expect low luminosity, low SN ellipti- cal galaxies to have relatively more metal– rich than metal–poor GCs. Given the smaller number of GCs present (and an even smaller population of metal–poor ones, if present at all) it may be very difficult to detect a bi- modal metallicity distribution. We predict that many ellipticals, without obvious bi- modality, will have metallicity (color) dis- tributions that are intrinsically broad, i.e. broader than the photometric or spectroscopic measurement errors.
• Tidally stripped galaxies may have steeper GC radial surface density slopes than non– stripped galaxies. Some initial support for this prediction comes from Fleming et al. (1995) who finds that galaxies in clusters have steeper GC radial slopes, at a given mag- nitude, than those in low density environ- ments. The study of Fornax cluster galaxies by Kissler–Patig et al. (1996) also supports this prediction.
• Tidally stripped galaxies will have lower than average SN values, i.e. ≤ 5. In the case of NGC 4486B and NGC 5846A we suggest that the vast majority of GCs and field stars have been removed by tidal interaction with their more massive neighbor, and we expect SN ∼ 1 for these two compact ellipticals.
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23 .2
0
-.2
-.4
-.6 Metal--rich [Fe/H] peak
-.8
-20.5 -21 -21.5 -22 -22.5 -23
-.5
-1
Metal--poor [Fe/H] peak Metal--poor [Fe/H] -1.5
-20.5 -21 -21.5 -22 -22.5 -23
Parent galaxy MV
Fig. 1.— Globular cluster mean metallicity – galaxy luminosity relation. The upper panel shows the mean metallicity of the metal–rich population from Table 1. The solid line is a weighted fit to the data. A good correlation between the metallicity of the metal–rich globular clusters and the parent galaxy luminosity is present. The lower panel shows the mean metallicity of the metal–poor population. The solid line is the metal–rich relation offset lower by 1 dex in [Fe/H] (i.e the average offset between the metal–rich and metal–poor populations). There is little or 24 no correlation between the metallicity of the metal–poor globular clusters and the parent galaxy luminosity. -1
-1.5
surface density slope surface -2
-2.5
0 5 10 15 20 25
SN
Fig. 2.— Globular cluster surface density slope versus specific frequency for elliptical galaxies. The power–law slope of the globular cluster surface density profile and the specific frequency (SN ) are taken from the list of Kissler–Patig (1996). For low SN galaxies there is a wide range of density slopes, but high SN galaxies only have flat, extended globular cluster distributions.
25 18
16
14
12
10 N S
8
6
4 N5846
2
0 0 .2 .4 .6 .8 1 1.2
NR/NP
Fig. 3.— Specific frequency versus the ratio of metal–rich to metal–poor globular clusters. The data point for NGC 5846 (NR/NP = 3.0, SN = 2.8 ± 1.7) is not shown for display purposes. A typical error bar for NR/NP is shown in the lower left. The data show that as the ratio of metal–rich to metal–poor globular clusters increases the specific frequency (SN ) value decreases. High SN galaxies have more metal–poor globular clusters than low SN ones. The merger model of Ashman & Zepf (1992) predicts an opposite trend to that seen i.e., an increasing NR/NP ratio with higher SN values. 26 0
-.2
-.4 logR ∆ [Fe/H]/ ∆ -.6
-.8
-1 0 5 10 15 20
SN
Fig. 4.— Globular cluster radial metallicity gradient versus specific frequency. The data suggest that high SN galaxies (from Table 2) have steeper radial metallicity slopes than low SN ones. The dashed line is the expected trend if an increase in SN corresponds to only larger numbers of metal–poor globular clusters which are placed with the same radial distribution as seen in NGC 4472 (see text for details). The zeropoint is taken to be the current values for NGC 4472 (i.e. SN = 5.5 and ∆[Fe/H]/∆logR = –0.41). The merger model of Ashman & Zepf (1992) predicts an opposite trend to that seen i.e., shallower metallicity27 gradients in high SN galaxies. lblrcutr a aebe curdb ia stripping. tidal by of acquired some case been each have conve In NG may been distribution. from clusters each have clusters globular below globular 5846 marked are stripped NGC 5846A tidally for NGC of (1996) metallicity Huchra mean & expected Brodie Forbes, and i.5—Goua lse ealct itiuini he eta domina central three Ostrov in distribution from metallicity data cluster Globular tometric 5.— Fig.
50 50
N1399 N4486 N5846 80
40 40 tal. et
60 19)frNC19,Giesler 1399, NGC for (1993) 30 30 28
Number 40 20 20
20 10 10 tal. et h nemdaemetallicity intermediate the tdit ealct.The metallicity. into rted 44 G 46 and 4486B NGC 1404, C 19)frNC4486 NGC for (1996) tglxe.Tepho- The galaxies. nt
0 0 0 -2 -1 0 1 -2 -1 0 1 -2 -1 0 1 [Fe/H] [Fe/H] [Fe/H] The Fornax Cluster
NGC 1380
NGC 1428 NGC 1373
NGC 1382 NGC 1374 NGC 1375
NGC 1381
NGC 1427
NGC 1379 NGC 1399 NGC 1387
NGC 1404 NGC 1427A
NGC 1389
10 ' 47 kpc NGC 1386
Fig. 6.— Schematic sketch of the galaxies in the central regions of the Fornax cluster. The X–ray ′ emission extends from the cD galaxy (NGC 1399) out to at least 38 .
29
Table Multimo dal Metallicity Distributions
Galaxy Type m M M N S Fe H Percent N N Source Ref
V GC N R P
N E B I
N E cD Wash
N E B I
N E cD Wash
N E Wash
N E Wash
N E cD Wash
N E V I
N E p Wash
N E V I
IC E V I
Notes to Table
Galaxy name Galaxy type Distance mo dulus from Harris Absolute V magnitude Total
number of globular clusters Sp eci c frequency Metallicity of p eaks indicates marginal signi cance
Percentage of globular clusters in each p eak Number of metal rich to metal p o or globular clusters
Source of metallicity i e broad band colors or Washington photometry Metallicity reference Bridges
et al b Ostrov et al Forbes et al Secker et al Zepf et al Geisler et al
Lee Geisler Forbes et al Harris et al Forbes Bro die Huchra
30
Table Metallicity Gradient Slop es
Galaxy m M M S Fe H logR Ref
V N
N
N
N
N
N
y
N
N
N
N
N
N
Notes to Table
Galaxy name Distance mo dulus from Harris Absolute V magnitude Sp eci c
00
y
frequency Globular cluster mean metallicity gradient in dex statistical error includes
data from Ajhar et al Reference Ostrov et al Secker et al Forbes et
al Zepf et al Unpublished metallicity gradients using data from Forbes et al
Geisler et al Lee Geisler Forbes Bro die Huchra
31
Table Tidally Stripp ed Candidates
Galaxy Type m M R V M log M Fe H
V pr edict pr edict
N E
N B cE
N A cE
Notes to Table
Galaxy name Galaxy type Distance mo dulus from Harris Pro jected separation
1
from primary in kp c Radial velocity di erence from primary in km s Observed absolute V
magnitude Central velocity disp ersion Predicted V magnitude from column Predicted
GC metallicity from column
Table Fornax Cluster Galaxies
Galaxy Type R M S Fe H
V N pr edict
N E
N E
N S
N E
N E
Notes to Table
Galaxy name Galaxy type Pro jected separation from NGC in kp c Absolute V magnitude Sp eci c frequency Predicted GC metallicity from column
32