WIDE-FIELD WASHINGTON PHOTOMETRY OF THE NGC 5128 GLOBULAR CLUSTER SYSTEM. II. LARGE-SCALE PROPERTIES OF THE SYSTEM

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1 The Astronomical Journal, 128: , 2004 August # The American Astronomical Society. All rights reserved. Printed in U.S.A. A WIDE-FIELD WASHINGTON PHOTOMETRY OF THE NGC 5128 GLOBULAR CLUSTER SYSTEM. II. LARGE-SCALE PROPERTIES OF THE SYSTEM Gretchen L. H. Harris 1 Department of Physics, University of Waterloo, Waterloo ON N2L 3G1, Canada; glharris@astro.uwaterloo.ca William E. Harris Department of Physics and Astronomy, McMaster University, Hamilton ON L8S 4M1, Canada; harris@physics.mcmaster.ca and Doug Geisler 1 Universidad de Concepción, Departamento de Fisica, Casilla 160-C, Concepción, Chile; doug@kukita.cfm.udec.cl Received 2003 October 14; accepted 2004 April 12 ABSTRACT BuildingontheCMT 1 photometric database presented in Paper I, in this paper we derive the large-scale properties of the globular cluster system (GCS) in NGC 5128, the nearest giant elliptical and the dominant galaxy in the Centaurus group. In global terms, it has a smaller total population than previously thought: we estimate clusters over all magnitudes, yielding a specific frequency S N ¼ 1:4 0:2, with a steep projected radial distribution r 2. The luminosity distribution of the clusters resembles that of an old, normal GC luminosity function (Gaussian-like with peak at M V 7:4 and dispersion of 1.3 mag), but these parameters are unfortunately quite uncertain because of the system s low population and the heavy field contamination. Using the metallicity-sensitive C T 1 color index, we discuss the metallicity distribution function (MDF) for a subsample of 211 previously identified clusters, all on a homogeneous photometric system. We find the MDF to be strongly bimodal, with metallicity peaks at ½Fe=HŠ ¼ 1:55 and 0.55 and with nearly equal numbers of clusters in each of the metal-poor and metal-rich modes. The combined evidence from the system s low specific frequency, the MDF, and the isophotal shell features in the halo light make a major merger a plausible model for the formation history of this giant E galaxy. However, the progenitor galaxies must have been more gas-rich than in any present-day mergers or starbursts. Finally, we present a list of 327 new cluster candidates not identified in any previous surveys; most of these are in the less well studied bulge region of the galaxy and along the minor axis. Key words: galaxies: individual (NGC 5128) galaxies: stellar content globular clusters: general On-line material: machine-readable table 1. INTRODUCTION In Harris et al. (2004, hereafter Paper I), we describe a new database of wide-field photometry in the Washington CMT 1 system covering the region around NGC 5128, the centrally dominant giant at the center of the Centaurus group. Here we use this material to derive the large-scale properties of this galaxy s globular cluster system (GCS), including the spatial distribution, globular cluster luminosity function (GCLF), and metallicity distribution function (MDF). 2. SELECTION OF OBJECTS FOR STATISTICAL ANALYSIS As we will show in the next section, more than 99% of the 100,000 objects measured in our study are likely to be either foreground stars in the Milky Way or faint, slightly nonstellar background galaxies. The globular cluster population we are seeking is only a trace component within this heavy contamination. As discussed in Paper I, we cannot objectively cull out contaminating objects on the basis of image morphology alone because, at a seeing of 1 00 or worse, only the largest clusters 1 Visiting Observer, Cerro Tololo Inter-American Observatory, which is operated by the Association of Universities for Research in Astronomy, Inc., under cooperative agreement with the National Science Foundation. 723 will appear nonstellar, while many, perhaps most, will be starlike. Thus, selecting globular clusters on the basis of image morphology alone would miss a large fraction of the cluster population. The fastest way to reduce the contamination significantly is by object color. In Figures 1 and 2 we show color-magnitude diagrams (CMDs) of the measured objects within the 1N2 box centered on NGC 5128; within this region we have almost complete areal coverage (see Paper I). The Washington photometric colors of known old globular clusters in this galaxy (Harris et al. 1992, hereafter H92) fall well within the range 0:45 < M T 1 < 1:1 and 0:30 < C T 1 < 2:5. Initially, therefore, we eliminate objects redder or bluer than these ranges (marked with dotted lines in the figures). Using of the two-color diagram in the Washington system (C M vs. M T 1 ) provides no significant additional help in this case, since many G and K field stars fall in just the same region as do normal globular clusters. This is illustrated in Figure 3. In the first of this pair of diagrams the color indices are plotted for objects previously known to be globular clusters in NGC 5128 (this sample will be defined in x 5 below). In the second graph the two-color plot is shown for all objects in our 1N2 box brighter than T 1 ¼ 20:5, the magnitude range containing the brighter half of the GCLF. It is apparent that the contaminating field objects strongly overlap the known

2 724 HARRIS, HARRIS, & GEISLER Fig. 1. CMD for 81,000 measured objects within a 1N2 box centered on NGC 5128, as described in Paper I. Lines at M T1 ¼ 0:45 and 1.10 show the upper and lower color boundaries of the region containing the old globular cluster population. clusters, and even with exquisitely precise two-color data we would not be able to cull out the field contaminants using only this combination of indices. We must also address the issue that the limiting magnitude of our data set differs from filter to filter, as well as from place to place across the 1N2 field of our survey, because the database is spliced together from many different mosaic pointings taken under a wide variety of observing conditions (see the extensive discussion of Paper I). Although some of the frames have detectable objects as faint as T1 ¼ 24, most do not, and the outer frames in particular have brighter limits. Careful inspections of (, ) plots for narrow magnitude ranges indicate that we cannot usefully employ data fainter than T1 ¼ 22:0; fainter than this, the data start becoming patchy in many places across the field. In Figure 4 we show the locations of the 38,600 objects from which we will derive the main properties of the GCS: namely, ones in the range 16 < T1 < 22 and 0:45 < M T1 < 1:1. The main impression from this figure is a nearly uniform spread of detected objects across the entire field, but with a faint hint of a concentration of objects around the galaxy center. It is this small signal that we are attempting to measure. 3. THE RADIAL DISTRIBUTION AND SPECIFIC FREQUENCY The first stage of our analysis is to gain a quantitative idea of the radial distribution and total population of clusters embedded within our subselected sample. Table 1 gives the relevant density distribution versus projected galactocentric radius, calculated for annuli of width 10. The only exception to this is the innermost ring, whose outer boundary is a circle of radius 4A2 and whose inner boundary is the exclusion ellipse of semimajor axis 4A2 and eccentricity 0.5. In all the rings the mean radius is defined empirically rather than geometrically, as the mean r-value of all the objects lying within Vol. 128 Fig. 2. CMD in T1 vs. C T1, for the same objects as in Fig. 1. Lines at C T1 ¼ 0:3 and 2.5 enclose the color range of the old globular cluster population. the ring boundaries. The successive columns of the table give the mean galactocentric radius of all objects in each annulus; the number of objects in the annulus that lie within the magnitude and color ranges specified above; the mean density (number of objects per unit area) and its statistical uncertainty; and finally the number of residual objects (by assumption, globular clusters) and their uncertainty after subtraction of the adopted background density b (defined below). The last four rows of the table give the -values for the separate outermost regions north, south, east, and west of the main body of data (see Fig. 4 and Paper I). In Figure 5 the number density is shown as a function of radius. In constructing this plot, we have ignored the effects of various tiny gaps in the areal coverage left behind by the multiply overlapping CCD fields (see again Fig. 4); these are several arcseconds wide at most and have no important effects on our results for the large-scale properties of the system. A clear rise in in toward the galaxy center is evident, but it is noteworthy that even in the innermost zones the residual signal is not strong relative to the background. The GCS around NGC 5128 is clearly not a populous one. To define the mean background level, we take the average of the rings within 20 0 < r < 30 0 and adopt b ¼ 6:90 0:10 objects per square arcminute. This level is consistent with the outermost fields covering r , which average b ¼ 6:97 0:27 (see Fig. 4 and the four outer points in Fig. 5). We use these outermost points only for rough confirmation of the background, since they have rather different effective limiting magnitudes from one to the other and are in any case so far from the center of NGC 5128 (and each other) that real field-to-field differences in the stellar population densities are possible. We conclude from Figure 5 that we have detected the clear presence of a population of clusters around NGC 5128, with colors in the appropriate range for old halo and bulge globular clusters. It is apparent from the numbers in the last two

3 No. 2, 2004 THE NGC 5128 GLOBULAR CLUSTER SYSTEM. II. 725 Fig. 3. Left: Two-color diagram for individually known globular clusters in NGC The colors plotted have been corrected for the foreground reddening of NGC 5128 (see text). Large solid dots show the ½C M ; (M T1 )0 -values for 62 clusters from the previous photometry of Harris et al. (1992); small crosses show the values for 155 additional clusters measured in the present BTC photometry, taken from Table 2 of this paper. Right: Two-color diagram for all objects in our BTC database brighter than T1 ¼ 20:5 and in the 1N2 box centered on the galaxy (see text). Fig. 4. Positions of objects in the color range 0:45 < M T1 < 1:10 and magnitude range 16 < T1 < 22 centered on NGC This list is expected to include the majority of old globular clusters around the galaxy. columns of Table 1 that the majority of the GCS lies within r ¼ 15 0, with almost negligible statistical signal beyond that. The main GCS contribution, in fact, lies within 100 (equivalent to a linear radius of 11 kpc for the 3.9 Mpc distance to NGC 5128), making its GCS globally a more compact system than found in many other giant ellipticals (Harris 2001; Harris 1993). Model fits of simple power laws to the backgroundsubtracted densities cl ¼ b r lead to a best-fit slope ¼ 2:0 0:2. Because of the bright bulge and wide central dust lane, our data do not reach inward far enough to delineate the flattening off of the profile and the core radius of the system, but from the innermost two data points in Table 1 we would estimate rc P 4 0 or about 4.5 kpc. For comparison, the Milky Way GCS follows a radial profile with 2:5 and most giant ellipticals are in the range 1:3 1:8. For ellipticals, a rough correlation of with galaxy luminosity exists in the sense that more luminous ellipticals have shallower GCS profiles, (0:29 0:03)MVT þ 8:00 (Harris 1993; Kaisler et al. 1996). With MVT ¼ 22:1 for NGC 5128, we would then expect from this relation a mean slope ¼ 1:6 0:3, somewhat shallower than our observed value of 2.0. However, there is quite a bit of galaxy-to-galaxy scatter in this trend, and the residuals may be correlated with the balance between metalrich and metal-poor clusters (see below). The very most extended GCSs with < 1:5 tend to be the cd-type central supergiants in rich Abell-type clusters. Other giant ellipticals with luminosities only slightly lower than NGC 5128 and with

4 726 HARRIS, HARRIS, & GEISLER Vol. 128 TABLE 1 Radial Profile of the Star Counts r (arcmin) N N(resid) : : : : k 2 include NGC 1404 and 1379 in Fornax, and field ellipticals NGC 1052 and 4278, among others (Harris 1986). Another case for comparison is that of the giant IC 4051 in the Coma Cluster (Woodworth & Harris 2000), which has a GCS with a profile shape 2:0 and a core radius r c ¼ 4:8 kpc, all of which resemble NGC 5128; however, for IC 4051 there is a strong possibility that its halo has been truncated by the huge tidal field of Coma, in contrast to NGC 5128, which is within the much smaller Centaurus group. Thus, while the steep profile and small core radius of the NGC 5128 GCS are not what we think of as typical for giant ellipticals, they are also far from unique. In Figure 6 the GCS profile is compared with the surface brightness profile of the underlying galaxy light, with photometry from van den Bergh (1976), which extends to r 10A5. Within the admittedly large internal uncertainties of the GCS data the slopes of both profiles match up extremely well. In the later discussion we suggest that this close agreement may be connected with the large numbers of metal-rich clusters in the galaxy, which are usually more centrally concentrated around the parent galaxy than metal-poor clusters and which arguably would have formed along with the main body of the galaxy. Using the data in Table 1, we find that the total residual cluster population after background subtraction, within the radial region of our counts, is N resid brighter than T 1 ¼ 22:0. To this we need to add estimates of the cluster numbers both outside r ¼ 20 0 and within the innermost Fig. 5. Radial distribution profile for objects in our database within the magnitude and color ranges that include globular clusters in NGC Data for the plotted points are taken from Table 1. The adopted far-field background level is shown by the dashed line at b ¼ 6:90 objects per square arcminute (see text). The solid line through the data points is a power-law model profile with cl r 2:0 (see text). elliptical region excluded from our data, which has an area of arcmin 2 and a semimajor axis 4A2. 2 Using the first two annuli in Figure 5 as a guide, we conservatively assume resid 4 1arcmin 2 for the central zone, giving an additional clusters. Extrapolating beyond r ¼ 20 0 requires a certain amount of judgement, since the residual counts become lost in the noise (Fig. 5), and thus the outermost part of the radial profile of the GCS is quite ill determined. We know from the list of individually found clusters (see x 5 below) that some certainly do exist at very large galactocentric distance: of the 211 known clusters listed in Table 2 below, 25 of them (or 11%) are at radii between 20 0 and Since the outer halo along its major axis has been rather carefully searched for cluster candidates and the crowding and background-light problems are lower there (see the discussions of Peng et al. 2004a; Harris et al. 1984; Hesser et al. 1984), this fraction is likely to be, if anything, an overestimate of the true ratio of clusters beyond 20 0 relative to the numbers within The method we adopt to estimate the total population is to use the surface brightness profile shown in Figure 6 (van den Bergh 1976) and extend it outward. Some additional confidence that this is a plausible procedure comes from the published HST/WFPC2 photometry of the halo stars in NGC 5128 (Harris & Harris 2000b), which shows that the same de Vaucouleurs r 1/4 profile established at r < 10 0 continues on outward to more than 27 0 (31 kpc). Numerical integration of this profile shows that the expected number of clusters beyond r ¼ 20 0 should equal about 7% of the total from 3 0 to 20 0, a ratio that agrees tolerably well with the relative numbers of individually known outer clusters. The exact percentage is insensitive to the outermost cutoff radius since the r 1/4 profile becomes quite steep at large radius and the integral converges rapidly. 2 This ellipse encloses the bulk of the large dust lane and the brightest areas of the galaxy bulge within which photometry could not be successfully carried out; see Paper I.

5 No. 2, 2004 THE NGC 5128 GLOBULAR CLUSTER SYSTEM. II. 727 Fig. 6. Radial distribution profile of the globular cluster system in NGC 5128, cl vs. radius after background subtraction. The plotted error bars include the uncertainties in both the object counts in each annulus and the background level. The solid curved line is the surface brightness profile of the galaxy light, from van den Bergh (1976), which covers r < 10A5. It has been arbitrarily shifted vertically to match the GCS data points. The dashed line shows the extension of the same halo light profile to r ¼ 25 0, as suggested by HST/WFPC2 photometry of the halo red giant stars (Harris & Harris 2000b). This extrapolation corresponds to an additional 50 clusters brighter than T 1 ¼ 22. The uncertainty in this excess we estimate very roughly as 30. Adding the directly counted clusters to the estimates for the innermost and outermost regions, we find that the net total cluster population over all radii and brighter than T 1 ¼ 22 is then N cl ¼ The last step to estimating the true total cluster population is to correct for the number of clusters fainter than T 1 ¼ 22. Anticipating the results of our later discussion on the luminosity function, we assume that the GCLF has a conventional Gaussian-like shape with a GCLF turnover, or peak frequency, at M 0 V 7:4 0:2 and dispersion V ¼ 1:3 0:1 (Harris 2001). An earlier direct verification of this assumption is given by Rejkuba (2001), who derived the GCLF from a small but deep cluster sample from VLT imaging. For an apparent distance modulus (m M ) V ¼ 28:30 and a typical cluster color (V T 1 ) 0 (V R) 0 0:5 (Harris et al. 1992; Harris & Harris 2002), we then have for the turnover magnitude T 0 1 ¼ 20:35, which means that T 1(lim ) ¼ 22 should include 90% 4% of the clusters. Our final estimate of the total globular cluster population is then N t ¼ The only previously published estimates of the total GCS population we can compare with are from the wide-field photographic star counts of Harris et al. (1984) and van den Bergh et al. (1981). Both those sets of counts extended to radii r 24 0, and both indicated as do our BTC data that the cluster system did not extend detectably above the background count level past radii of In their star counts from blue-sensitive IIIaJ plates, van den Bergh et al. (1981) found a net population of about 600 clusters brighter than a very roughly determined limiting photographic magnitude J 22 (or about V 21:5, half a magnitude fainter than the GCLF turnover and effectively about a magnitude brighter than our BTC limit). Harris et al. (1984) obtained star counts in three filters (U, V, andr) and to somewhat different limiting magnitudes in each case, but their data for the R filter have a quoted limit of R lim ¼ 22:0 0:25, fortuitously similar to our BTC limit. To that limit, they found a residual population of about objects over the radial range 1A38 < r < 12A1, a total only about 10% 20% higher than ours. Given the traditional difficulties of establishing a consistent detection limit from star counts done by eye inspection on wide-field photographic plates, we believe that the mutual agreement of the previous photographic surveys with our wide-field CCD photometry is good (for example, a correction to their plate limit estimate of just 0.2 mag would bring the photographic counts completely in line with our CCD data). The specific frequency S N of the GCS (Harris & van den Bergh 1981) is S N ¼ N t 10 0:4(M T V þ15) for galaxy luminosity M T V. If the total integrated magnitude for NGC 5128 is V T ¼ 6:2 (van den Bergh 1976), then M T V 22:1, and thus S N ¼ 1:4 0:2. This specific frequency is low even for field ellipticals and actually falls close to the range of S N values for cluster-rich giant spirals such as M31 (Harris 2001). We will discuss the implications of this result in x 7. In previous catalogs of specific frequency (Harris 1991, 2001), NGC 5128 has usually been listed as having S N 2 3. These earlier estimates were based on the star count totals from Harris et al. (1984). From their directly observed numbers of clusters to a given limiting magnitude, these authors tried several different methods to extrapolate the total to all radii and all magnitudes, including (1) the use of only the counts from 1A38 < r < 16A1, (2) a linear extrapolation inward to r ¼ 0, (3) fitting an r 1/4 halo profile to extrapolate to larger radii, and (4) three possible adopted distances (3, 5, and 8 Mpc) since the distance modulus was not well determined at that time. Combinations of these different approaches (see Cluster (hr) TABLE 2 Washington Photometry for Individual Globular Clusters (deg) T 1 M T 1 C M C T 1 V C C C C C C C Note. Table 2 is presented in its entirety in the electronic edition of the Astronomical Journal. A portion is shown here for guidance regarding its form and content.

6 728 HARRIS, HARRIS, & GEISLER Vol. 128 their Tables 4 and 5) led to the conclusion that the probable value of N t remains in the range In the subsequent literature values of N t near the high end of their range of possibilities tended to be adopted, leading to the S N 2:5 level that was usually quoted since the mid-1980s. Our new work improves on these previous photographic star counts in several obvious ways: the field size is even larger, the limiting magnitudes are more rigorously determined, the star counts are uniformly done to the same limits across the entire field, half the field contamination is objectively removed via color selection, and (not least important) the distance to the galaxy is accurately known. From the combination of the red giant branch tip magnitude, the planetary nebula luminosity function, and the bulge surface brightness fluctuation (see Harris et al for a summary), the uncertainty in the distance calibration is now 0.15 mag. Reinterpreting the older star counts with the benefit of hindsight, we find that their true background level had already been reached at r 12 0 and that little further extrapolation was needed either for the GC luminosity function or radii beyond the count area. In fact, our new totals agree rather well with the old photographic ones for their minimal case 1 method (see their Table 4), which corrects the raw counts only for the limiting magnitude and the small unobserved area around the galaxy center. 4. LUMINOSITY FUNCTION Studying the luminosity distribution of clusters makes the issue of differential limiting magnitude across the field more critical, since in this case we are explicitly interested in levels as faint as we can reach in an unbiased way. As indicated above, the BTC survey data reach a useful limit of T 1 ¼ 22, which is roughly 1.6 mag fainter than the expected GCLF turnover point. We use the same subsample as defined above (objects in the color range 0:45 < M T 1 < 1:1), and now derive the GCLF. Knowing (Fig. 3) that the majority of the clusters are within 10 0 of the galaxy center, we compare an inner region 4A2 < r < 10 0 with a region of similar area immediately outside it, 10 0 < r < 14 0 and plot the residual population as a function of magnitude. By tightly restricting the background region this way, we avoid, as much as possible, any changes in effective limiting magnitude with position in the field. The inner region is also selected to avoid the bright galaxy bulge almost totally. The results are shown in Figure 7. The overall residual population can be tolerably matched by a standard Gaussianlike GCLF with a turnover magnitude at T 1 ¼ 20:4 and a dispersion 1.3 mag, corresponding to V(turnover) 20:9 and thus M V ¼ 7:4 for our adopted distance modulus (m M ) V ¼ 28:3 and foreground reddening E B V ¼ 0:11. This Gaussian function is shown as the solid line on the graph. However, any more detailed conclusions are prevented by the internal uncertainties. That is, we are trying to deduce the features of the GCS population by subtracting a large amount of noise from a combined signal-plus-noise that is scarcely larger. In a statistical sense, there are essentially no clusters in the system brighter than T 1 ¼ 17 (corresponding to M V 11, near the limit for the very most luminous globular clusters known). Curiously, the residual population drops out near T 1 20:5, just where we would nominally expect to find the GCLF turnover point, but the internal error bars are so large that it is difficult to ascribe any real significance to this. Alternate choices of the background annulus in the range end up giving residual GCLFs that are different Fig. 7. Derivation of the GCLF of the cluster system. The solid histogram (top curve) shows the number of objects per 0.2 mag bin in the color range 0:45 < M T 1 < 1:1 and within the radial range 4A2 < r < 10 0, while the dashed histogram shows the background LF for the radial range , normalized to the same area as the inner region. The data points with error bars show the residual LF (inner minus background). A Gaussian GCLF is drawn in for comparison, with turnover at T 1 ¼ 20:4 and dispersion ¼ 1:3 mag. only in small details, but with the same basic features. As noted above, the characteristics of the raw data prevent us from making any deductions about the GCLF fainter than T Frustratingly, we are left with only the rather weak conclusion that the GCLF for NGC 5128 is roughly consistent with the standard one for giant E galaxies, having a peak at M V 7:4 and Gaussian dispersion near 1.3 mag. A considerably cleaner sample of clusters will be required to establish the GCLF characteristics in greater detail. 5. THE COLOR AND METALLICITY DISTRIBUTION One of the primary goals of our observations (Paper I) was to use the highly metallicity-sensitive Washington index C T 1 to construct the MDF for the system. Previous studies with samples of only a few dozen clusters (Harris et al. 1992; Zepf & Ashman 1993; Rejkuba 2001; Held et al. 2002) indicated that the MDF for the NGC 5128 clusters has the now classic bimodal form that is seen in many large galaxies (NGC 5128 was one of the first giant E systems claimed to have a bimodal MDF; see the discussion of Zepf & Ashman 1993 based on the H92 data). With a considerably larger sample of known clusters in hand, we are in a position to reevaluate the MDF, where all data are on one homogeneous photometric system. Our primary new look at the MDF is from the pure sample of bona fide globular clusters found individually from previous searches (Hesser et al. 1984; Sharples 1988; Rejkuba 2001; Peng et al. 2004a). Eliminating duplicates between these four studies, we find a total of 211 objects matching with our BTC database. We list our photometry and identifications in Table 2. In the first column the C-numbers refer to cluster identifications from Hesser et al. (1984), G-numbers from Sharples (1988), f1- and f2-numbers from Rejkuba (2001), and PFF-numbers from Peng et al. (2004a). The coordinates listed are all from our BTC data (see Paper I for an extensive comparison of our BTC coordinates with those measured in the previous studies).

7 No. 2, 2004 THE NGC 5128 GLOBULAR CLUSTER SYSTEM. II. 729 Fig. 8. Locations of previously identified globular clusters in NGC Coordinates, are plotted relative to galaxy center. The central circle enclosing most of the dust lane and bulge light has radius 4 0. Crosses denote clusters from Hesser et al. (1984) or Sharples (1988); triangles are candidates from Rejkuba (2001); and solid dots are clusters from Peng et al. (2004a). In Figure 8 we show the locations of these objects on the sky. Most of them lie along the isophotal major axis of the galaxy (northeast to southwest). The objects from Hesser et al. (1984) and Sharples (1988) (crosses) and Peng et al. (2004a) (dots) were selected by radial velocity and thus are nearly certain clusters belonging to NGC The objects from Rejkuba (2001) (triangles) are in two smaller regions on the sky to the south and northeast and were selected only by image morphology and color; the majority should be clusters but a few background galaxies or stars may be present as well. A more subtle feature of this sample of clusters is that they may be biassed in favor of ones with large galactocentric distance, since only about half are within the 10 0 zone that our global radial profile shows to enclose most of the system (Fig. 5). However, any such radial bias should not lead to a bias in estimating the MDF, as long as the cluster system contains little or no global metallicity gradient (see below). Despite any selection effects, Figure 8 shows unequivocally that the GCS extends detectably to r gc k 40 kpc as does the halo light of the galaxy (Peng et al. 2002) even though most of the system is well inside that distance. The color-magnitude arrays for this selected set of objects are shown in Figure 9. Half a dozen objects have very red colors, M T 1 k 1:3; these are all from the R01 sample, and we suggest that these are probably faint background galaxies. In addition, a few faint objects with clearly blue colors, C T 1 P 1, and with radial velocities consistent with NGC 5128 membership also appear in the diagrams. These probably represent clusters that are much younger than conventional old-halo globulars (see Peng et al. 2002), and as such they deserve closer follow-up analysis. Fig. 9. Color-magnitude distributions in M T 1 and C T 1 for previously identified globular clusters. Symbol types are the same as in the previous figure.

8 730 HARRIS, HARRIS, & GEISLER Vol. 128 The color index (C T 1 ) 0 ¼ (C T 1 ) 1:97E(B V )isa sensitive indicator of metallicity [Fe/H] for old globular clusters. As noted in Paper I, we use E(B V ) ¼ 0:11 for the foreground reddening of NGC Geisler & Forte (1990) first employed this index and derived a linear transformation of (C T 1 ) 0 to metallicity using the [Fe/H] data then available for the Milky Way clusters (Harris & Canterna 1977). A revised and more accurate transformation, employing the most recent Milky Way cluster data (Harris & Harris 2002) is ½Fe=HŠ ¼ 6: :82(C T 1 ) 0 þ 0:162(C T 1 ) 2 0 : ð1þ This curve is modestly nonlinear relative to the old Geisler/ Forte equation. The top panel of Figure 10 shows the color distribution for the individual cluster sample in (C T 1 ) 0. The selections of these objects in the studies cited above have carefully avoided areas projected on the major dust lane of the galaxy, so we use a uniform reddening E(B V ) ¼ 0:11 for all of them. The shaded part of the histogram includes clusters confirmed by radial velocity membership, while the unshaded part (the objects from Rejkuba) were selected by image morphology; as noted above, the reddest of these may not be genuine clusters. The bottom panel of Figure 10 shows the deduced MDF binned in steps of ½Fe=HŠ ¼0:10, corresponding roughly to the internal precision of the data. We see that in either index, C T 1 or [Fe/H], the MDF is strongly bimodal, with a dividing line at (C T 1 ) 0 1:4 or½fe=hš 1:0. The bluer mode has 52.5% of the clusters (see Table 3), and the redder mode 47.5%. A simple double-gaussian fit to the MDF (shown superposed on the histogram in Fig. 10) provides a tolerable match to the MDF as a whole, with peaks at ½Fe=HŠ ¼ 1:55 0:05 [or (C T 1 ) 0 ¼ 1:20] and 0:55 0:05 [(C T 1 ) 0 ¼ 1:63]. The standard deviations of the two modes are ½Fe=HŠ ¼0:30 (metal-poor) and 0.22 (metal-rich). 3 This numerical fit should, however, be considered useful only as a rough description of where the two subgroups of clusters are concentrated and not as a model with any physical basis. The total range of metallicities is restricted to the interval 2:2 P ½Fe=HŠ P 0:0 already very familiar from the Milky Way and M31; as yet, we have no clear evidence that clusters exist in NGC 5128 that are significantly more metal-rich than solar abundance (see also the discussion of H92). Nevertheless, most of the inner-halo, presumably metal-rich component remains to be definitively explored. It should also be noted that the nonlinearity of the transformation from (C T 1 ) 0 to [Fe/H] has the effects of both slightly compressing the red (metal-rich) end of the distribution in going from the top panel of Figure 10 to the bottom panel; and stretching the blue (metal-poor) end. Thus, the small excess of clusters on the extreme metal-poor tail is likely to be due simply to photometric scatter, which has more of an effect at the blue end. For example, at ½Fe=HŠ ¼ 2, an error in color of 0.05 in C T 1 generates a spread of ½Fe=HŠ ¼ 0:15, whereas at solar metallicity the spread is only onethird as large. 3 The means and standard deviations of the two Gaussian curves differ slightly from those quoted for the two modes in Table 3, particularly for the metal-rich end. This is because the entries in the table are the raw totals including the outlying red or blue objects, which increase the -values particularly. For any comparative purposes we recommend using the means and standard deviations quoted in the text above. Fig. 10. Metallicity distribution function for individual previously detected globular clusters in NGC The sample of objects is the same as in the previous figure. Top:Histogramof(C T 1 ) 0 -values. The shaded section represents clusters known to be members of NGC 5128 by radial velocity, while the unshaded portion represent objects selected by image morphology alone. Bottom: Histogram for the same sample, converted to [Fe/ H] with the equation in the text. The two Gaussian curves have centroids at ½Fe=HŠ 0 ¼ ( 1:55; 0:55) and standard deviations (0.30, 0.22). Our C T 1 histogram resembles the ones published by Peng et al. (2004b) from UBVI color indices (see their Fig. 1), though it is constructed from a slightly different list of objects. We find a relatively higher number of red clusters than they do, but these numbers should not be overinterpreted. The samples are still only a fraction of the total GCS, and further rigorous culling of the sample by radial velocity measurement may change the ratio since more than half of the field contamination is from redder objects. In addition, we note that the conversion from any of these broadband color indices to [Fe/H] is expected to be mildly nonlinear, and so any comparisons of histograms in color versus those in metallicity, such as from our Figure 10 or the directly spectroscopic one of Held et al. (2002), must keep this effect in mind. The range of metallicities, the form of the MDF, and the locations of the two modes are similar to what Held et al. (2002) has found from a sample of 57 clusters in NGC 5128 with spectroscopically determined metallicities. The MDF strongly resembles the distinct bimodal form found in many other large ellipticals in Virgo, Fornax, and other nearby galaxy clusters (Gebhardt & Kissler-Patig 1999; Neilsen & Tsvetanov 1999; Larsen et al. 2001; Kundu & Whitmore 2001). Thus, NGC 5128 supplies further evidence that ge galaxies in relatively sparse environments can contain large numbers of quite metal-rich globular clusters, as well as metal-poor ones. For giant ellipticals, galaxy environment does not seem to be a dominant factor in determining the shape of the MDF. The metallicity derived from the (C T 1 ) 0 color index as a function of galactocentric radius is shown in Figure 11 for the

9 No. 2, 2004 THE NGC 5128 GLOBULAR CLUSTER SYSTEM. II. 731 TABLE 3 Numbers of Metal-Poor and Metal-Rich Clusters Radius (arcmin) n blue n red hfe/hi b (Fe/H) b hfe/hi r (Fe/H) r hfe/hi all (Fe/H) all All radii same group of objects. Among other features, it shows clearly that metal-rich clusters are not restricted only to the bulge regions, but can be found at quite large distances out in the halo. This feature of the GCS is now realized to be routinely present in many large ellipticals (e.g., Geisler et al. 1996; Rhode & Zepf 2001; Dirsch et al. 2003) and can be interpreted either within a major-merger formation scenario (Zepf & Ashman 1993; Bekki et al. 2003) or a hierarchical-merging scheme (Beasley et al. 2003). Figure 11 can also be used to briefly investigate the distribution for any trace of a metallicity gradient. A rough quantitative test is summarized in Table 3, where we list the mean metallicities of the blue (½Fe=HŠ < 1:0) and red (½Fe=HŠ > 1:0) clusters in various radial bins, plus the mean of the whole sample. No significant change with radius occurs within either the blue or red clusters, and the mean of the whole GCS remains nearly constant until r k 20 0, beyond which the metal-richer clusters drop off in numbers relative to the metal-poor ones. The lack of metalrich objects at these outermost radii is likely to be real, since the sample of known clusters was not selected on the basis of metallicity in any of the discovery studies. In other words, any overall metallicity gradient within the entire system is due mainly to changing proportions of metal-poor versus metalrich clusters. The overall effect is strikingly similar to what has been found in other giant ellipticals (e.g., Geisler et al. 1996; Larsen et al. 2001). In Figure 12 we show the MDF of our cluster sample along with the metallicity distribution for field stars in the NGC 5128 halo (Harris & Harris 2002). For purposes of this comparison, we use the field-star data at 8 kpc from the galaxy center, the innermost and most metal-rich of four fields along the southwest major axis that have been studied with deep WFPC2 photometry (see Harris & Harris 2000a; Harris & Harris 2002). For the field stars, V I colors were used with a grid of globular cluster red giant fiducial tracks to construct the MDF. An uncertainty in the comparison that needs to be noted is that the field-star MDF is measured in terms of ½m=HŠ log (Z=Z ), whereas the globular cluster MDF is in terms of [Fe/H]. For the Milky Way halo and globular cluster stars, [m/h] is larger than [Fe/H] by amounts of typically (e.g., Shetrone, Côté, & Sargent 2001); we have used the same conversion factor for the NGC 5128 stars, although we have no direct information to support that assumption. Another unavoidable uncertainty in the comparison of Figure 12 is that the field-star sample is taken from only one projected location in the halo (8 kpc), whereas the GCS covers the entire halo. We would expect regions of the halo and bulge still farther inward to be at least slightly more metal-rich than our 8 kpc field, whereas field stars from farther out in the halo (Harris & Harris 2000a) are a bit more metal-poor. A global sample of the field stars should thus have a somewhat broader MDF than is shown here. Despite these caveats, the major peak and width of the field-star MDF and the metal-rich globular clusters are similar enough to suggest that both were Fig. 11. Color index vs. projected galactocentric distance, for previously identified globular clusters. Here we show the metallicity derived from the C T 1 color index vs. r in arcminutes, where 1 0 is approximately equal to 1 kpc. Symbol types are the same as in the previous figure. Fig. 12. Comparison of the MDF for the NGC 5128 globular cluster system (points with error bars from Fig. 10, rebinned in steps of 0.2 dex) with the MDF for the halo field stars at a projected distance of 8 kpc (dashed line) from Harris & Harris (2002). The field-star MDF, which was defined in terms of ½m=HŠ¼log (Z=Z ), has been converted with the assumption ½Fe=HŠ ¼ ½m=HŠ 0:3 (seetext).

10 732 HARRIS, HARRIS, & GEISLER Vol. 128 Fig. 13. Locations of new candidates for NGC 5128 globular clusters as listed in Table 4, identified by their nonstellar image structure and color indicesasdiscussedinthetext. formed in the same event, or events, which built the major part of this galaxy. By contrast, very little star formation seems to have accompanied the blue (presumably metal-poor and older) GC formation. 6. ADDITIONAL CLUSTER CANDIDATES Much recent progress has been made in defining the properties of the NGC 5128 GCS, but certain clear problems remain. Our BTC database provides a relatively unbiased sample but is hampered by major field contamination, whereas the smaller samples of confirmed clusters used in the previous section (Hesser et al. 1984; Sharples 1988; Rejkuba 2001; Peng et al. 2004a) are relatively uncontaminated but may be biased by their initial selection toward objects in the outer halo. The studies of Sharples and Rejkuba also surveyed only small spatial regions. We suggest that a necessary step forward to more detailed understanding of this difficult GCS is to strongly reduce the crippling problems of field contamination and bias by identifying many more of the clusters individually. The problem appears to us to be particularly acute in the inner regions of the galaxy, where the numbers of known clusters must be very incomplete. As a step toward increasing the sample of known clusters, we have returned to our BTC image database and more carefully examined selected frames to find new candidates for globular clusters. The exposures of best seeing (usually the short-exposure M images, which for the inner fields have 1B0 1B2 image quality) were blinked with their PSF-subtracted images generated by DAOPHOT/ALLSTAR. Any objects that are slightly nonstellar appear in the subtracted frames as small doughnut-shaped rings with oversubtracted cores and undersubtracted wings (see Rejkuba 2001 for a similar approach). These candidates were marked, their coordinates were determined, and they were matched with our BTC photometry database. Any objects with extremely red colors (M T 1 k 1:3, which must be background galaxies or foreground stars) were rejected. This procedure generated a list of 398 candidates. By Fig. 14. CMD for the new cluster candidates, with data from Table 4. comparing these with the 211 previously known clusters in our data, we found that we had successfully rediscovered 71 of them, leaving 327 new candidates not in any previous cluster list. In addition to some of the more spatially extended globular clusters in NGC 5128 this sample will include a mix of small, symmetrically shaped background galaxies and interloping stars. The locations of the candidates are shown in Figure 13, and their color-magnitude distribution is shown in Figure 14. Many of them (particularly the fainter, redder ones) are very likely not to be clusters, and it is quite possible that the success rate for the entire candidate list (given our rediscovery rate for real clusters, and the fraction of all measured objects within r 10 0 of galaxy center that are clusters) may be no larger than about one-third. In Table 4 we list the available data for these (positions and our Washington photometric indices). A handful (Fig. 14) can be seen to be faint and blue; these may represent genuinely young clusters. Note that some of the objects fall within our inner exclusion ellipse and thus do not have BTC photometry. (deg) TABLE 4 List of New Cluster Candidates (deg) T 1 M T 1 C T Note. Table 4 is presented in its entirety in the electronic edition of the Astronomical Journal. A portion is shown here for guidance regarding its form and content.

11 No. 2, 2004 THE NGC 5128 GLOBULAR CLUSTER SYSTEM. II. 733 We emphasize that this candidate list is not in any sense homogeneous or complete to a particular magnitude limit or area of coverage. The candidates are simply objects that our raw, somewhat heterogeneous image data have allowed us to identify as having image morphology, magnitudes, and colors like those of known clusters, and which give emphasis to the less well studied inner regions of the galaxy and the regions along the minor axis. The next stage should be to obtain radial velocities, which would not only confirm or reject their cluster nature in a definitive way, but would also contribute to a much improved dynamical analysis of the NGC 5128 GCS (Hesser et al. 1986; Peng et al. 2004b). In a new observational program started with the Anglo-Australian Telescope, some of the candidate list has now been covered (Beasley et al. 2004), and we are continuing the investigation. 7. COMMENTS AND DISCUSSION In many of the previous studies of NGC 5128 an ongoing goal has been to relate the GCS properties to the possible formation history of the galaxy. Clearly, the dust lane, gas, and young star formation that are obvious in the inner 5 kpc of the galaxy indicate a relatively recent satellite accretion event (less than 500 Myr ago; see Ebneter & Balick 1983). In addition, the many faint arcs and ripples superposed on the smooth halo light extending out 15 kpc or more (see Peng et al. 2002) are strongly suggestive of mergers from much longer ago. But were these events relatively minor ones superposed on a basic hierarchical-merging history that built the main body of the giant elliptical? Or could it have formed from a very small number of major mergers of preexisting disk galaxies? A key number in such discussions has always been the GCS specific frequency S N : disk galaxies invariably have low S N P 1:5, while E galaxies can have values anywhere from 1 to 10 or even higher, depending on a wide variety of illunderstood factors, including luminosity, environment, and presence of hot halo gas (Harris 2001; McLaughlin 1999; Kavelaars 1999; Blakeslee 1999). High-S N ellipticals, such as the Virgo giants or cd-type galaxies, are unlikely to have formed from low-s N spirals and dwarfs since it is extraordinarily difficult for the relative number of globular clusters to increase during a merger. For such an increase to happen, both huge amounts of gas and very high cluster efficiency would be necessary (see Harris 2001 for more comprehensive discussion). Field ellipticals, however, tend to have low specific frequencies, S N P 2, as do large galaxies that are clearly the results of recent mergers (e.g., Whitmore et al. 2002; Goudfrooij et al. 2001; Miller et al. 1997; Schweizer et al among others). For the first time we have a reliable specific frequency measurement for NGC 5128, with well-determined estimates for both the cluster population and the galaxy distance, and we findthatitfitsintothelow-s N category. The sheer number of globular clusters, then, would not present a barrier to formationofthisgalaxybymajormergers(harris2003)andwould be consistent with the isophotal properties of the halo mentioned above. If we take the proportions of clusters discussed in x 5atface value, NGC 5128 has 510 blue, metal-poor clusters and 470 red, metal-rich ones. Its blue clusters (see below) could have been accumulated from three or four disk galaxies the size of M31 or the Milky Way, or equally well from a larger number of dwarfs. The metal-richer clusters now present in the galaxy would, by hypothesis, be a combination of bulge clusters from the progenitor disk galaxies, plus ones built in the subsequent mergers. Model calculations of major mergers indicate that under the right conditions the bimodal MDF, weak metallicity gradient and highly concentrated metal-rich component will closely resemble real cases such as NGC 5128 (Bekki et al. 2003). Additional features of the globular cluster MDF must, however, also be carefully considered. In major spirals such as M31 and the Milky Way, roughly two-thirds of the globular clusters are metal-poor (½Fe=HŠ < 1) and for dwarf galaxies the MDF is even more strongly weighted to metal-poor clusters (Harris 2001). In contrast, for NGC 5128 there are nearly equal numbers of metal-poor and metal-rich clusters. In the standard merger formation picture (Ashman & Zepf 1992) we would therefore want the progenitor galaxies to be very gas-rich to permit many new metal-rich clusters to be built during the mergers and thus tilt the balance of the resulting MDF further toward the metal-rich end. A sample calculation will illustrate the sequence. If the premerger disk galaxies are like the Milky Way or M31, all 510 of the metal-poor ones would have come directly from the progenitor galaxies, as well as no more than about 250 of the metal-rich ones. Thus, at least 220 metal-rich clusters need to be built from the relatively enriched gas in the mergers. At a normal cluster formation efficiency of M cl =M gas 0:0015 and a mean cluster mass of 3 ; 10 5 M (Harris 2001), we then require the progenitors to have brought in 4:4 ; M of gas. For comparison, the total stellar mass in NGC 5128 at the present time is near 2 ; M (Mathieu et al. 1996). However, if the progenitors were mostly dwarf galaxies instead of large spirals, then almost all of their original clusters would have been metal-poor, which would mean that virtually all of the metal-rich clusters now in the system should have formed during their progressive mergers, requiring an even higher mass of gas to be involved. Moreover, the normal cluster formation efficiency quoted above corresponds to a specific frequency S N 3:5; for lower S N,the amounts of initial gas would have to be proportionately higher to produce the same final number of clusters. For NGC 5128 (with S N ¼ 1:2) the gas used for star and cluster formation in the merger(s) would have to total more than M.No galaxies in present-day mergers or starbursts are this gas-rich. The suggestion from this logic is that the progenitor galaxies needed to be primarily gaseous at the time of merger an overly strong condition for today s universe, but one easier to satisfy at much earlier epochs. Somewhat paradoxically, it therefore seems appropriate from the GCS evidence to argue that formation via major mergers is possible for NGC 5128, but that the mergers needed to be so gas-rich that they resemble aspects of hierarchical merging in the early universe. Obviously, there are no sharp boundaries between these formation regimes. For galaxies in lower density environments, including small groups like Centaurus, the sequence of mergers could have taken longer to complete than in rich Abell-type clusters, so relatively gas-rich systems could have continued to merge and accrete for many gigayears. The metal-rich clusters in NGC 5128 could therefore span a large range of ages, as is suggested by Peng et al. (2004b). On the other hand, if more detailed spectroscopic investigation shows that they are all conventionally old (10 Gyr or more), this would argue that the major part of the formation took place rather quickly at early times, more like the ge galaxies in rich environments. Preliminary analyses of spectroscopic age-sensitive line indices by Held et al. (2002); Peng et al. (2004b) indicate that the metal-richer clusters do appear to be