The Cluster Velocity Dispersion of the Abell 2199 cD Halo of NGC 6166
Hobby-Eberly Telescope (HET) spectroscopy is used to measure the velocity dispersion profile of the nearest prototypical cD galaxy, NGC 6166 in the cluster Abell 2199. We also present composite surface photometry from many telescopes. We confirm the defining feature of a cD galaxy; i. e., a halo of stars that fills the cluster center and that is controlled dynamically by cluster gravity, not by the central galaxy. Our HET spectroscopy shows that the velocity dispersion of NGC 6166 rises from km s in the inner to km s at 100 in the cD halo. This extends published observations of an outward increase and shows for the first time that rises all the way to the cluster velocity dispersion of km s. We also observe that the main body of NGC 6166 moves at km s with respect to the cluster mean velocity, whereas the velocity of the inner cD halo is 70 km s closer to the cluster velocity. These results support our picture that cD halos consist of stars that were stripped from individual cluster galaxies by fast tidal encounters.
However, our photometry does not confirm the widespread view that cD halos are identifiable as an extra, low-surface-brightness component that is photometrically distinct from the inner, steep-Sérsic-function main body of an otherwise-normal giant elliptical galaxy. Instead, all of the brightness profile of NGC 6166 outside its core is described to 0.037 mag arcsec by a single Sérsic function with index . The cD halo is not recognizable from photometry alone. This blurs the distinction between cluster-dominated cD halos and the similarly-large-Sérsic-index halos of giant, core-boxy-nonrotating ellipticals. These halos are believed to be accreted onto compact, high-redshift progenitors (“red nuggets”) by large numbers of minor mergers. They belong dynamically to their central galaxies. Still, cDs and core-boxy-nonrotating Es may be more similar than we think: Both may have outer halos made largely via minor mergers and the accumulation of tidal debris.
We construct a main-body cD-halo decomposition that fits both the brightness and dispersion profiles. To fit , we need to force the component Sérsic indices to be smaller than a minimum- photometric decomposition would suggest. The main body has 30 % of the total galaxy light. The cD halo has , 1/2 mag brighter than the brightest galaxy in the Virgo cluster. A mass model based on published cluster dynamics and X-ray observations fits our observations if the tangential dispersion is larger than the radial dispersion at to 60. The cD halo is as enhanced in element abundances as the main body of NGC 6166. Quenching of star formation in 1 Gyr suggests that the center of Abell 2199 has been special for a long time during which dynamical evolution has liberated a large mass of now-intracluster stars.
Matthews, Morgan, & Schmidt (1964) and Morgan & Lesh (1965) introduced the
This paper presents two new observational results:
1 – Section 2 demonstrates that the velocity dispersion of the stars in the nearest, prototypical cD galaxy – NGC 6166 in the cluster Abell 2199 – rises from values typical of giant elliptical galaxies near the center to the cluster dispersion in the cD halo. The halo also shifts toward the velocity of the cluster, which is different from that of NGC 6166. Thus the halo shares the dynamics of individual galaxies in the cluster. We interpret this as evidence that the stars in the cD halo of NGC 6166 were stripped from the galaxies by fast collisions.
2 – We measure the brightness profile of NGC 6166 to make quantitative Morgan’s point (d) that cDs consist of a central elliptical plus a distinct, shallow-brightness-gradient halo. Photometry by Oemler (1976) suggested that NGC 6166 has such two-component structure. Our ideas about cD halos are based in large part on this result. However, we find that NGC 6166 is described by a single Sérsic (1968) profile at all radii outside the core. The cluster-dominated halo that is obvious in the kinematics is not obvious in the photometry. We need to rethink our understanding of how we recognize cDs and of whether cD galaxies are fundamentally different from other giant, core-boxy-nonrotating elliptical galaxies.
2. HET Spectroscopy: Velocity And Velocity Dispersion Profiles of NGC 6166
2.1. History and Motivation
To distinguish between competing theories about the origin of cD galaxies (Section 8), a particularly powerful diagnostic is their internal kinematics. Does the velocity dispersion profile increase to the cluster velocity dispersion as one looks farther out into the part of the halo that encompasses many non-central cluster members? Is the systemic velocity of the halo similar to that of the central galaxy or is it similar to that of the cluster as a whole? Are these velocities ever different? This subject has a long history, and partial answers to these questions have been known for several decades:
Zabludoff, Huchra, & Geller (1990) find that NGC 6166 has km s for galaxy and cluster velocities of km s and km s (71 galaxies). Zabludoff et al. (1993) find that km s; km s; km s for 68 cluster galaxies. Oegerle & Hill (2001) get peculiar velocities of 258 69 to 346 73 km s, depending on how is calculated and on how far out in the cluster the ( 132) galaxies are counted. The derived peculiar velocity gets smaller as more galaxies get averaged. Among recent determinations, Coziol et al. (2009) get km s; km s; = 156 km s for 471 cluster galaxies. The most up-to-date study by Lauer et al. (2014) gets = 9317 10 km s; = 9088 38 km s; = 229 39 km s for 454 cluster galaxies.
Many cDs are essentially at rest at their cluster centers (e. g., Quintana & Lawrie 1982; Zabludov et al. 1990; Oegerle & Hill 2001). Generally, cDs are more nearly at rest in their clusters than are non-cD first-ranked galaxies (e. g., Oegerle & Hill 2001; Coziol et al. 2009). But a significant fraction move at several hundred km s with respect to their clusters, often in association with cluster substructure which suggests that a merger of two clusters is in progress (e. g., Oegerle & Hill 2001; Pimbblet, Rosebloom, & Doyle 2006; see also Beers & Geller 1983; Zabludoff et al. 1990, 1993). Proof of concept is provided by the Coma cluster. It is in the process of a cluster merger (White, Briel, & Henry 1993; Briel et al. 2001; Neumann et al. 2001, 2003; Gerhard et al. 2007; Andrade-Santos et al. 2013; Simionescu et al. 2013). The NGC 4839 group is falling into the main Coma cluster, which itself has two central galaxies, NGC 4874 and NGC 4889, with different velocities (by about 680 km s) and their own X-ray halos. NGC 4889 has a velocity of km s with respect to the Coma cluster. Only NGC 4874 is within 250 km s of the cluster velocity. NGC 4874 and NGC 4889 are weak cDs, and NGC 4839 also shows signs of cD structure.
NGC 6166’s velocity with respect to Abell 2199 is typical. The diagnostic question is: Does the halo of NGC 6166 have the same systemic velocity as its central galaxy or as its cluster? We find that the cD halo shows velocities between that of the galaxy and that of the cluster. The observation that NGC 6166 is not centered in velocity in its cD halo is evidence that that halo does not belong dynamically to the galaxy.
Velocity Dispersion Profiles
In a paper that fundamentally shaped our concept of cD galaxies, Dressler (1979) pushed measurements of velocity dispersions to then-unprecedented low surface brightnesses and showed that for IC 1101, the brightest galaxy in Abell 2029, rises with increasing radius from 375 km s at the center to 500 km s at 71 kpc. (The distance has been converted to the WMAP 5-year cosmology distance scale, Komatsu et al. 2009; NED.) Thus the dispersion rises toward but does not reach the cluster of km s (Coziol et al. 2009) or km s (Lauer et al. 2014). Dressler interpreted this in the context of suggestions (Gallagher & Ostriker 1972; White 1976; Ostriker & Tremaine 1975; Richstone 1976; Merritt 1983; Richstone & Malumuth 1983) that cD halos consist of accumulated debris of stars stripped from cluster members by tidal encounters and by dynamical friction against the growing halo. Thus a cD consists of “a luminous but normal elliptical galaxy sitting in a sea of material stripped from cluster galaxies” (Richstone 1976; Dressler 1979). Dressler concludes: “The results of this study confirm an [outward] increase in velocity dispersion, which is a necessary (but not sufficient) condition in the proof of the stripped debris hypothesis”. Sembach & Tonry (1996) and Fisher et al. (1995) confirm these results.
Among the nearest galaxies, M 87 is marginally a cD in the sense of having extra light at large radii with respect to an Sérsic fit (Figure 50 in Kormendy et al. 2009; hereafter KFCB). This is a normal Sérsic index for a core-boxy-nonrotating elliptical, but the amount of extra light is small, and in fact, an Sérsic function fits the whole profile outside the core. This Sérsic index is outside the range normally observed for core-boxy-nonrotating Es. Nevertheless, a cD halo cannot securely be identified as an outer component that is photometrically distinct from the main body of the galaxy. At the same time, it is clear that the Virgo cluster does contain intracluster stars, from broad-band surface photometry (Mihos 2011; Mihos et al. 2005, 2009), from spectroscopy of individual stars (Williams et al. 2007b), and from the detection of intracluster globular clusters (Williams et al. 2007a) and planetary nebulae Arnaboldi et al. 1996, 2002, 2004; Castro-Rodriguéz et al. 2009; see Arnaboldi & Gerhard 2010 and Arnaboldi 2011 for reviews). The intracluster light is irregular in its spatial distribution and defined largely by (tidal?) streams. It is reasonable to conclude that the intracluster light is in early stages of formation. Nevertheless, it pervades the cluster and must feel the cluster gravitational potential. And the outer halo of M 87 merges seamlessly with this intracluster light (Mihos papers). Do we observe that the velocity dispersion of stars in M 87 increases toward the cluster dispersion?
The answer – tentatively – is yes. The integrated light shows an outward drop in from km s in the central few arcsec to km s at and then an outward rise to km s at (Murphy et al. 2011, 2014). This is subtle and not easily interpreted. But the upward trend in continues in the globular cluster population, which reaches 400–470 km s by (Wu & Tremaine 2006; see Côté et al. 2001 for earlier results). Planetary nebula data in Doherty et al. (2009) reveal both M 87 halo and intracluster stars, but the data are too sparse to determine a profile. Also, though they do not overlap greatly in radial leverage, stellar dynamical models and mass profile measurements from the X-ray gas give essentially consistent results (e. g., Gebhardt & Thomas 2009; Churazov et al. 2008). Thus, M 87 is the nearest galaxy where various test particles have been used to probe the dynamics of a marginal cD from its center out to radii where the cluster dominates. The problem is that the test particles are heterogeneous enough and the statistics for point particles are poor enough so that we cannot securely see the transition from the galaxy’s main body to any halo that is controlled by cluster gravity. Nevertheless, as a proof of concept, M 87 is important. And it provides a hint that proves to be prescient: The dispersion profile starts to rise at kpc, well interior to the radii where any plausible argument identifies the beginning of a cD halo based on photometry alone.
Ooutward rises in cD or cD-like galaxies are reported by Carter et al. (1981, 1985) and by Ventimiglia et al. (2010). Still, the only prototypical cD in which the velocity dispersion of the stellar halo is robustly seen to rise toward larger radii by several authors is NGC 6166 in the cluster Abell 2199. From a central velocity dispersion of 300 km s, the dispersion first drops outward and then rises to 400 km s (Carter, Bridges, & Hau 1999) at about 30 and 600 100 km s (Kelson et al. 2002) at 50– 60. No velocity dispersion measurments of any cD galaxy reach large enough radii to show that increases all the way up to the cluster dispersion.
The first purpose of this paper is to push the measurements of NGC 6166 far enough out in radius to see whether or not reaches the cluster dispersion.
2.2. HET Spectroscopy
We obtained spectra at three slit positions (Figure 1) along and near the major axis of NGC 6166 with the 9.2 m Hobby-Eberly Telescope (HET) and Low Resolution Spectrograph (LRS: Hill et al. 1998). The slit width was 1.5, the reciprocal dispersion was 116 km s pixel, and the resolution expressed as a velocity dispersion was km s. The slit positions had exposure times of 8 900 s (“center”, with NGC 6166 centered well inside the slit), 4 900 s (“offset” position along the major axis, centered on the bright, elongated galaxy NGC 6166A visible in Figure 1), and 6 900 s 1 800 s (“alternate” position offset by 11.5 from the major axis but on the other side of the center, positioned to miss star and galaxy images). All individual exposures were taken on different nights. The standard star spectrum used is a combination of HD74377 and HR2600. Our experience is that this combination fits old elliptical-galaxy stellar populations well and give kinematic results that are relatively free from template bias. In any case, the kinematics were measured with Bender’s (1990) Fourier correlation quotient method, which is designed to eliminate template bias.
Figure 2 shows an unsharp-masked version of the sum of the best spectra along the central slit position (white line in Figure 1). By dividing out the brightness profile of the galaxy, we can see absorption lines and qualitatively judge from the center out to the largest radii. The strongest lines in NGC 6166, Mg b, Na D, and H, are visible all the way to the companion galaxy on the slit. Even Fe 5270 Å and 5335 Å are visible quite far out (see also Figure 3). They are used in Section 6 to measure [/Fe] overabundance out into the part of the halo where the velocity dispersion is large. Most important, Figure 2 already shows that all lines except Na D get very wide in the cD halo of NGC 6166.
The Na D line is narrow at all radii and shows little gradient in velocity. It gives a dispersion of km s at all radii. We assume that the line is produced by interstellar gas and do not include it in the wavelength region from redward of the iron lines to blueward of H that we use for and measurements. Dust is seen near the center in Figure 8. There may be a more smoothly distributed ISM at larger radii, as suggested also by the fact that H absorption in our spectra is significantly weaker than even a very old stellar population would show. However, it is not obvious that its kinematics should be a simple as we measure with the Na D line. Interpretation of this line in the context of the X-ray gas halo of the galaxy is beyond the scope of this paper.
The offset and alternate slit positions yielded poorer spectra. We discard one spectrum taken with too much moonlight, so the alternate slit position has only 6 good spectra. Of these, one is fainter than normal by 14 % and two more are fainter by 23 %, presumably due to clouds. (The observations are
queue-scheduled, so we cannot personally monitor the observing conditions. However, we checked that the galaxy was centered on the slit. Seeing is relatively unimportant.)
Offset sky spectra were taken after all NGC 6166 exposures. For the center slit position, these were cleaned of bad pixels and averaged to give high and then used for all sky subtractions. Each spectrum was individually sky-subtracted before the spectra were added. The sky subtraction of the central slit spectra is good (Figure 2). However, for the other two slit positions, most sky spectra could not be used for sky subtraction because too many night sky emission lines changed in strength in the short time between exposures. For the offset sky positions, the sky was measured as far from the galaxy as possible; since even the NGC 6166 end of the slit is far from the galaxy (Figure 1), these sky spectra should be essentially free of galaxy light. However, for the alternate slit position, sky spectra taken from the galaxy images do subtract a little halo light. For this reason – as well as problems with moonlight and with clouds – the alternate slit position does not reach as far out as the primary slit position illustrated in Figure 2. In addition, we found that we got the best results to the largest radii in the alternate slit position by using only the four best spectra.
Figure 3 shows sample spectra for five radial bins in NGC 6166 and for the optimized template star. This binning is used in Section 6 to measure line strengths for the Mg b and Fe lines. Reliable line strength measurements are possible out to the bin at . Velocity dispersions are easier – they are measureable for the bin and for several others at large radii in the center, alternate, and offset slit positions.
2.3. Kinematic Results
The summed center, alternate, and offset spectra were reduced with the Fourier correlation quotient program of Bender (1990). This gives velocity , velocity dispersion , the higher-order Gauss-Hermite coefficients and , and nonparametric line-of-sight velocity distributions (LOSVDs). At some radii near (see Figure 2), the LOSVDs show a main peak at the systemic velocity of NGC 6166 and smaller peak in its wings associated with another of the multiple nuclei. We omitted the corresponding velocity bins from the LOSVD fit. Since neither the center nor the radii where starts to climb are affected, this cleaning does not affect our conclusions. However, many published and measurements show contamination from the multiple nuclei.
The instrumental velocity dispersion was measured in our reduced spectra to be km s, easily adequate for the galaxy dispersions 300 km s studied in this paper.
The kinematics are listed in Table 1 and shown in Figure 4.
2.4. The Velocity Profile of NGC 6166
The systemic velocity of NGC 6166 is km s higher than the velocity km s of 494 cluster galaxies (Lauer et al. 2014). Here we use our measure of the systemic velocity of NGC 6166, km s. It is consistent within errors with values in Zabludoff et al. (1993) and in Coziol et al. (2009). Other, inconsistent published measurements may be affected by contamination by the multiple nuclei. Using our , NGC 6166 moves at , typical of the values found by Lauer et al. (2014).
If the cD halo consists of tidal debris, then we expect that its systemic velocity should shift toward that of the cluster at the radii where rises toward the cluster value. Figure 4 shows that the velocity at large radii on both sides does decrease from toward the cluster velocity. The average of the large-radius points is only km s. Still, the inner part of the cD halo of NGC 6166 is – as far as we can measure it – more nearly at rest within the cluster than is the central galaxy.
2.5. The Velocity Dispersion Profile of NGC 6166
Figure 5 compares our kinematic results on NGC 6166 with published dispersion profiles. Carter et al. (1999) and Kelson et al. (2002) observed much of the rise in to the cluster value. However, our observations are the first to reach deep enough to see for the intergrated starlight in a cD halo rise all the way to the cluster dispersion in any galaxy cluster.
The Carter et al. (1999) data are not shown in Figure 5, because they did not publish a table of their results. Their outermost measurements at radii of 30 to 36 are , 361, and 438 km s. These are consistent with our results and with Kelson’s. (However, Carter et al. 1999 derive velocities that increase as increases; they interpret this as “modest major-axis rotation”. Kelson et al. 2002 also see “systematic rotation  in the intracluster stars beyond 20 kpc”. We do not see rotation; rather, the halo velocity decreases toward the cluster velocity on both sides of the center.)
Tonry (1984, 1985) measured the multiple nuclei of NGC 6166 but did not reach far enough out to see an outward increase in . Similarly, Fisher, Illingworth, & Franx (1995) and Loubser et al. (2008) measured only a slight outward drop in in the main body of the galaxy.
Figure 5 illustrates the most important result in this paper: The velocity dispersion in NGC 6166 increases outward to a weighted mean of km s for the four data points at to 107. This equals the velocity dispersion km s for 454 galaxies in Abell 2199 (Lauer et al. 2014). The rise in to the cluster velocity dispersion is seen in all three of our slit positions. This result is the strongest evidence that the cD halo of NGC 6166 is made of stars that have been accreted in minor mergers or stripped from cluster galaxies by dynamical harassment.
3. Surface Photometry: Does NGC 6166 Have a Photometrically Distinct Halo?
3.1. The Standard Picture of cD Halos
Our standard picture of the nature of cD halos and the way in which we identify cD galaxies are based in large part on photometry of NGC 6166 and other cD galaxies by Oemler (1976). Oemler’s procedures and conclusions were later made quantitative by Schombert, as discussed below. But the iconic, two-component structure suggested by Oemler’s photometry of cDs – particularly NGC 6166 – firmly cemented in our minds the notion that cDs consist of an elliptical-galaxy-like central body plus a photometrically distinct, shallower-surface-brightness halo that is not present in normal giant ellipticals. Oemler’s profile of NGC 6166 – augmented by Hubble Space Telescope (HST) photometry to improve the central spatial resolution – is shown in Figure 6.
The clearly two-humped profile in Figure 6 decisively quantifies Morgan’s description of his visual impression of two-component structure. Other cDs in Oemler (1976), in Schombert (1986, 1987, 1988), and in other papers from the same era behave similarly. The picture of cD halos that has been in our minds ever since is made still more concrete using modern profile analysis machinery by decomposing the profile into two Sérsic (1968) functions. Several recent papers have done this and suggested that the inner components are normal ellipticals whereas the cD halos have exponential profiles (Seigar, Graham, & Jerjen 2007; Donzelli, Muriel, & Madrid 2011). In fact, the Sérsic-Sérsic decomposition in Figure 6 requires that the cD halo have , between an exponential () and a Gaussian () in its outer cutoff. A worrying hint is that the inner profile has , smaller than we have found for any other elliptical (KFCB). Note that, in making this fit, we have been very conservative about excluding the inner, shallow-power-law core (see Kormendy et al. 1994; Lauer et al. 1995, and KFCB for the definition of cores and Gebhardt et al. 1996; Kormendy 1999, and Lauer et al. 1995 for a demonstration that they are features of the unprojected and not just the projected profiles). We also omit the central AGN from the fit. About 2/3 of the light of the profile in Figure 6 is in the cD halo.
The ideas summarized above were made more quantitative by Schombert (1988). Schombert (1986, 1987) measured average surface brightness profiles of non-first-ranked ellipticals as functions of galaxy absolute magnitude in seven bins from to ( km s Mpc). Schombert (1988) then used these template profiles to define cD galaxies. First, the template profile is found that best matches the inner profile of the candidate galaxy over the largest possible radius range. If this profile fits all of the candidate galaxy to within the scatter seen among the individual profiles that were used to make the template, then this galaxy is an ordinary elliptical. In contrast (Figure 1 in Schombert 1988; cf. Figure 6 here), if the galaxy in question has a giant outer halo above the template profile fitted to the inner parts, then the galaxy is a cD and the integrated difference between its observed profile and the best-fitting template is the cD halo. This definition is similar in spirit to one used by Oemler (1976) but has the advantage of allowing the profiles of ellipticals to depend on luminosity. And it has the virtue of being nonparametric – it does not depend on describing the inner profile with an analytic fitting function.
The profile decomposition shown in Figure 6 is nothing more nor less than Schombert’s procedure in parametric form, using Sérsic functions for the inner and outer components. Much experience in recent years has shown that Sérsic functions are excellent fits to elliptical-galaxy profiles (see KFCB for data and review) and hence also to Schombert’s template profiles. However:
We find a problem with our canonical picture of cD halos (§ 3.2). The photometry shown in Figure 6 is in error. Our composite profile measurements of NGC 6166 are very well fitted by a single Sérsic function at all radii outside the central core. In contrast to our kinematic results, there is no photometric hint of two-component structure.
3.2. Composite -Band Brightness Profile of NGC 6166
We have measured the - and -band surface brightness profiles of NGC 6166 using CCD images from four ground-based telescopes and four cameras (WFPC1 PC, WFPC2 WF, ACS, and NICMOS2) on HST. Parameters of the images are listed in Table 2. This section discusses the -band profile.
The central profile is from an HST WFPC1 measurement by Lauer et al. (1995), from our measurement of an HST WFPC2 F555W image (GO program 7265; D. Geisler, P. I.), and from our high-resolution (Gaussian dispersion radius .) -band image from the Canada-France-Hawaii Telescope Cassegrain camera. The CFHT observing run is discussed in KFCB. The three images give independent -band zeropoints that agree (fortuitously) to much better than mag arcsec. The three zeropoints have been averaged.
Similar in resolution to the CFHT Cassegrain image is a image from the CFHT Megacam. We also include photometry of an image from SDSS; it is used over a larger radius range to derive the -band profile in the next subsection, but it is used here to help to tie together small and large radii, and it helps to measure the ellipticity and PA profile. The outer profile is obtained using a -band image from the Wendelstein Observatory’s new 2 m Fraunhofer Telescope (FTW) and a -band image from the McDonald Observatory 0.8 m telescope. The latter profile reaches , where mag arcsec. The -band profile of NGC 6166 is similar in accuracy and limiting surface brightness to the data in KFCB.
Figure 7 shows the raw profiles. Three kinds of profiles are shown. Most are based on isophote fits as in Bender (1987), Bender & Möllenhoff (1987), and Bender, Döbereiner, & Möllenhoff (1987, 1988). The algorithm fits ellipses to the galaxy isophotes; it calculates the ellipse parameters surface brightness, isophote center coordinates and , major and minor axis radii, ellipticity , and position angle PA of the major axis. Radial deviations of the isophotes from the ellipses are expanded in a Fourier series in the eccentric anomaly ,
The most important parameter is , expressed in the figures as a percent of the major-axis radius . If , the isophotes are disky-distorted; large at intermediate radii would indicate an S0 disk. If , the isophotes are boxy. The importance of these distortions is discussed in Bender (1987, 1988); Bender et al. (1987, 1988, 1989, 1994); Kormendy & Djorgovski (1989); Kormendy & Bender (1996), KFCB, Kormendy (2009), and below.
Some profiles were measured using Lauer’s (1985) program profile in the image processing system VISTA (Stover 1988). The interpolation algorithm in profile is optimized for high spatial resolution, so it is best suited to our high- images of the core of NGC 6166. The isophote calculation is Fourier-based, so it not well suited to measuring the outer parts of NGC 6166, where masking of other galaxies in the cluster results in very incomplete isophotes.
Finally, as discussed further below, we use a major-axis, (02-) two-pixel-wide cut profile to verify that the ellipse fitting was not adversely affected by the companion galaxies.
Seriously discrepant data in the profiles at small radii (usually because of inadequate spatial resolution) and at large radii (usually because of spatial variations in sky brightness) were pruned out before final averaging. Two additional complications require discussion:
(1) Three additional cluster galaxies lie in projection close to the center of NGC 6166 (e. g., Minkowski 1961; Burbidge 1962; Tonry 1984, 1985). Profile calculations need to correct for the light of these galaxies. Lauer (1986; see also Lachièze-Rey et al. 1985) decomposed the four galaxies using ground-based images and concluded that the two large companions are relatively undistorted, consistent with the hypothesis that they are not strongly interacting with NGC 6166. It was already known that the brighter two companions differ in velocity from NGC 6166 by and km s (e. g., Minkowski 1961); these velocity differences are consistent with true separations that are similar to the projected ones, but they do not clearly establish a close physical relationship. We follow Lauer and assume that NGC 6166 itself is not affected by the companions. We therefore calculate its profile by masking out the companions.
(2) There is patchy dust absorption near the galaxy center. We take this into account next.
Figure 8 illustrates both problems. The top image shows isophotes at average major-axis radii of 79, 117, 186, and 249. Above the center, all contours except the one at 249 are substantially affected by the closest companion. Various strageties were used to correct for the companions. For some profiles, the companions were masked; for others, contaminated pixels were replaced by pixels from the opposite side of the galaxy center. The same strategy was used on the dust contamination; the most reliable results were obtained by interpolating through the dust in the right-hand quadrants and then replacing the most strongly affected pixels
in the left quadrants by pixels from the opposite side of the center. All these procedures are somewhat vulnerable, because isophote fitting requires many pixels that need correction. So, as a check on the isophote fitting, we derive a major-axis cut profile along the vertical line in Figure 8. The cut is 2 pixels = 02 wide in the F555W WF image. The lower part of Figure 8 shows that the cut is minimally affected by dust (a few pixels were corrected). More importantly, we used pixels only from the bottom half of the image at radii where the top half is affected by the companions shown and only from the top half of the image at much larger radii where a companion not illustrated in the figure begins to be important.
Figure 9 shows that the average -band composite profile is robustly determined. We have enough different data sets with different problems (e. g., non-flatness of the sky brightness) so that agreement among data sets reliably identifies problem points. They are pruned. Near the center, the profiles that are corrected with Lucy (1974) - Richardson (1972) deconvolution – i. e., the ones from Lauer et al. (1995) and from the CFHT Cassegrain camera – agree with the much higher-resolution WF profile. In fact, since the -band cut profile is most free of dust effects, it is used at radii near 1 in preference to the Lauer et al. (1995) data. (The difference is only a few hundredths of a mag arcsec – see Figure 11.) Most important: The major-axis cut profile agrees with the isophote fit profiles to 0.02 mag arcsec. The success of this check is important to our confidence in the final profile.
The average -band photometry is tabulated in Table 3.
3.3. Composite -Band Brightness Profile of NGC 6166
An -band composite profile is derived in Figure 10, albeit from few sources. We need it primarily as another check of the -band profile, including the ellipticity and position angle. The central profile and VEGAMAG zeropoint are from an HST ACS F814W image (GO program 9293; H. Ford, P. I.). It helps that dust is less important at band. However, we can go further: the availability of an ACS F475W image (GO program 12238, W. Harris, P. I.) allows us to make a dust-corrected image, as follows.
First, the F475W -band image was rotated and registered to -pixel accuracy with the F814W -band image. Then a dust-corrected -band image was derived using the procedure described in Nowak et al. (2008, Appendix A) and summarized here. In the following, and are the F475W and F814W surface fluxes per square arcsecond; no subscript indicates magnitudes or fluxes as observed, a subscript ‘0’ refers to an extinction-corrected quantity. From the relation,
where is the absorption and is the reddening in the color , it follows that:
If the stellar population gradient in the inner regions of NGC 6166 is negligible, then constant and thus:
The parameter is determined by
where we have assumed a standard extinction curve to obtain the numerical value for the filters considered here (e. g. Savage and Mathis 1979).
The correction is not perfect, because it is based on the assumption that all of the dust is in a screen in front of the image. In NGC 6166, most of the dust is near the middle of the galaxy, in front of only about half of the stars. Then Equation (5) overcorrects for the dust. Better results are obtained if we adopt a smaller value for (a value of would imply no correction). After some experimenting, we adopt , which yields the smoothest appearance of the isophotes. Explicitly,
The residual dust contamination is small.
Then the brown circles in Figure 10 are derived from the dust-corrected image using Bender’s isophote fitting program. The red points are derived using VISTA profile on the dust-corrected image after 80 iterations of Lucy-Richardson deconvolution and after further cleaning of dust as discussed in § 3.2. These profiles agree essentially perfectly.
A final check is possible using an HST NICMOS2 F160W image (GO program 7453, J. Tonry, P. I.). There is no star in the field of view, so we do not attempt PSF deconvolution. But dust is essentially unimportant. The core profile calculated from this image also agrees very well with the -band results, when PSF blurring is taken into account. In particular, the F160W profile confirms that the core profile is cuspier at red wavelengths than it is in band.
3.4. Photometry Results. I. The Profile in the Core
Figure 11 illustrates our § 3.3 conclusion: The core profile of NGC 6166 is cuspier at red wavelengths than it is in band. We suggest that the difference is caused by -band absorption over the entire central arcsec of the galaxy. Clear hints of widespread, low-level absorption are visible in Figure 8.
It is difficult to measure the power-law cusp slope far inside the profile break radius (Lauer et al. 2007). The reason is that the nuclear source is spatially resolved and has an unknown profile. Whether it consists of stars or an AGN or some combination, we cannot subtract it robustly. However, the shallowest -band slope at to 10 corresponds to a Nuker function (Lauer et al. 1995) . This agrees with obtained in Lauer et al. (2007; any correction for the nuclear source is not discussed). Previous estimates, (Lauer et al. 1995) and (Byun et al. 1996), were determined from the Lauer et al. (1995) -band PC1 profile shown below. Our -band cut profile is even flatter than Lauer’s profile – it is less affected by patchy dust – so our composite -band profile is even less cuspy than .
The cuspiness of the central profile affects no conclusions of this paper. But it will be important to use the appropriate, dust-free profile if in future we obtain stellar kinematic data that allow a dynamical search for a supermassive black hole.
3.5. Photometry Results. II. The cD Structure of NGC 6166 is Not Recognizable from the Shape of the Brightness Profile
Our profile measurements in Figures 9 and 10 do not show the two-component structure that is so obvious in Figure 6. We believe that Oemler (1976) profile is in error; the most likely reason is the difficulty of correcting for the many cluster galaxies that overlap the cD halo. Modern ellipse-fit software copes more robustly with incomplete isophotes.
A single Sérsic (1968) function fits the complete profile of NGC 6166 outside the cuspy core. Both this result and the Sérsic index, in band or in band, are completely normal for core-boxy-nonrotating ellipticals. Figure 12 compares NGC 6166’s profile shape with the sample of elliptical galaxies studied by KFCB. They found that ranges from to 9 1 for their core ellipticals (red profiles in Figure 12). NGC 6166 is virtually indistinguishable from these galaxies; indeed, many core ellipticals have shallower outer profiles than does NGC 6166. It is especially interesting to contrast NGC 6166 with M 87. M 87 is by all arguments a more marginal cD than NGC 6166. But a Sérsic fit to its overall profile gives , larger than in NGC 6166. Plausible allowance for a cD halo in M 87 – i. e., exclusion of the outermost profile points – gave a marginally better fit with , consistent with our fit to NGC 6166 but with only a little extra light in the cD halo of M 87. Such a halo is less – not more – obvious in NGC 6166.
A two-component, Sérsic-Sérsic decomposition is allowed by our data (§ 4), but the fit is not significantly better than the one-component decomposition. There is no reason to believe that we detect two components from photometry alone.
This is a surprising result. We plan but have not yet carried out similar photometry of other cD galaxies. We therefore do not know that the present results on NGC 6166 apply more generally to all cD galaxies. Nevertheless:
We arrive at an ironic situation: The spectroscopy results resoundingly confirm our standard picture that the cD galaxy NGC 6166 in Abell 2199 has an outer halo that consists of debris from member galaxies. The halo stars are dynamically controlled by the cluster, not the central galaxy, and they have the kinematics (i. e., more nearly the systemic velocity and the velocity dispersion) of the other galaxies in the cluster, even when the cD is dynamically colder and in motion with respect to the sea of background stars. But the supposedly much easier task of recognizing the presence of a cD halo from two-component structure in the surface brightness distribution turns out to fail dramatically in the nearest, most prototypical cD galaxy, NGC 6166.
3.6. Photometry Results. III. Recognizing NGC 6166 as a cD Galaxy via Quantitative Differences in Structural Parameters
Is it possible to recognize cD galaxies by photometry alone? A photometric technique is desirable, because spectroscopy to look for an outward rise in velocity dispersion is expensive. Our results suggest a partial answer: The cD nature of NGC 6166 can be recognized via quantitative differences in structural parameters and parameter correlations. This helps but is not entirely satisfactory. Parameter distributions for cD galaxies and non-cD ellipticals overlap. There may be physics in this. The physical differences between cDs and core-boxy-nonrotating ellipticals may be smaller than we have thought. Figures 13 and 14 illustrate these points.
Figure 13 compares the brightness profile of NGC 6166 to the Virgo cluster elliptical galaxies. Radii are plotted in kpc. NGC 6166 has a larger and fainter core than any elliptical in Virgo, including M 87. And its outer profile is shallower and it reaches larger radii than that of any elliptical in Virgo, including M87. Quantitatively, the extreme cD NGC 6166 is distinguishable from normal core ellipticals. However, the marginal cD M 87 (see KFCB) overlaps with other core ellipticals in its profile properties.
Figure 14 compares the structural parameters of NGC 6166 with parameter correlations from KFCB and from Kormendy & Bender (2012). These are projections of the “fundamental plane” correlations (Djorgovski & Davis 1987; Faber et al. 1987; Dressler et al. 1987; Djorgovski et al. 1988; Djorgovski 1992; Bender et al. 1992, 1993), between the effective radius that encloses half of the light of the galaxy, the effective brightness at , and (in this case) total absolute magnitude.
NGC 6166 parameters are based on an assumed distance of Mpc (NASA/IPAC Extragalactic Database “NED” (Local Group) for cluster Abell 2199 and the WMAP 5-year cosmology parameters, Komatsu et al. 2009). NGC 6166 is plotted twice in Figure 14:
To get the less extreme point, we integrate the brightness and ellipticity profiles (that is, the two-dimensional isophotes) to the outermost data point in Figure 9, i. e., where mag arcsec. This gives , , kpc, and mag arcsec.
Galactic absorption corrections are from Schlegel et al. (1998). This point in Fig. 14 is consistent with a slight extrapolation to higher luminosity of the correlations for other ellipticals.
The more extreme point is derived by extending the profile to Mpc using the overall Sérsic fit and keeping the outer ellipticity constant at the last observed value. The limiting surface brightness is 30.9 mag arcsec; this is an “integration to infinity” similar to those discussed in KFCB. Then , , kpc, and mag arcsec. Within the scatter, this point is consistent with a larger extrapolation of the correlations for normal ellipticals. It deviates slightly from linear correlations in having larger and fainter , but slightly curved fits to normal ellipticals would not show NGC 6166 as deviant.
We conclude that NGC 6166 is more extreme than the ellipticals in the combined sample in Figure 14 in the sense expected for a cD: It has larger effective radius and fainter effective brightness. In this sense, the cD structure is recognizable quantitatively in the parameter correlations.
cD and non-cD galaxies overlap in parameter distributions (Schombert 1986, 1987). And yet, the cD NGC 6166 is qualitatively different from non-cD ellipticals, even brightest cluster galaxies. This important, because cD and brightest cluster galaxies are often considered to be equivalent. But NGC 6166 is surrounded by an immense halo of stars that are controlled dynamically by the cluster potential, not by the central galaxy. Isolated ellipticals cannot have such halos, and observations of velocity dispersion profiles in non-cD core ellipticals show no rise in at large radii (e. g., Kronawitter et al. 2000; Proctor et al. 2009; Weijmans et al. 2009; Foster et al. 2011; Raskutti, Greene, & Murphy 2014).
We conclude (1) that cD structure is real and distinct from non-cD ellipticals but (2) that it is difficult to recognize the difference photometrically. Extreme structural parameters help (Figure 14). But in less extreme cases – and, to be certain, even in NGC 6166 – velocity dispersion data are required to identify cluster halos reliably. The fact that cD classification is difficult is our problem, not the galaxy’s.
4. A Photometric and Kinematic Decomposition of NGC 6166 Into an Elliptical Galaxy Plus a cD Halo
This section presents a decomposition of the inner, E-galaxy part of NGC 6166 and its cD halo that accounts for both the photometry and the velocity dispersion profile of the galaxy.
The best-fit two-component Sérsic-Sérsic decomposition is illustrated in the left part of Figure 15. We emphasize: the RMS deviations 0.034 mag arcsec of the profile from the fit within the fit range (vertical dashes across the and profiles) are not significantly better than the deviations (Figure 9 RMS = 0.037 mag arcsec) of a single-Sérsic fit.
The decomposition in Figure 15 is similar to those in Huang et al. (2013a, b) – it minimizes for two Sérsic components. Huang and collaborators interpret such decompositions as supporting a two-phase scenario of elliptical galaxy formation (Oser et al. 2010; Johansson et al. 2012) in which wet mergers rapidly build high-, compact “red nuggets” (Buitrago et al. 2008; van Dokkum et al. 2010; Papovich et al. 2012; Szomoru et al. 2012) that later grow high-Sérsic-index halos via minor mergers. The inner component(s) in the decomposition are interpreted as descendent(s) of the red nuggets, and the outer component is interpreted as a later-accreted debris halo. Such a picture may be correct. But (1) it is not compellingly supported by the conclusion that two components fit the data better than one, and more importantly, (2) NGC 6166, with its cD halo, is a clearcut example of essentially the above processes, and in it, a two-component decomposition made by minimizing fails to explain the kinematics. As follows:
The observed dispersion profile implies that the central galaxy contributes most of the light along the line of sight out to kpc ( Mpc). The brightness profile extends out to kpc in the cD halo. In the transition region, we look through a short line of sight through the galaxy and a long line of sight through the halo. This suggests a simple procedure to capture the essence of the profle. We assume that the components have independent Gaussian LOSVDs. To keep things simple, we assume that the galaxy has the brightness profile of Component 1 in Figure 15 and that it has km s at all radii. We assume that the cD halo has the brightness profile of Component 2 and km s at all radii. This is an oversimplification. But if the decomposition in Figure 15 is approximately correct, then it should approximately fit the dispersion profile. It fails. The components overlap too much in radius; i. e., the inner component contributes too much light at large radii for the dispersion profile to increase outward as quickly as we observe toward km s. Modifying the assumed inner and outer dispersions does not help.
So a two-Sérsic-component photometric decomposition that minimizes fails to explain the velocity dispersion profile of NGC 6166. This argues for caution in the increasingly popular practice of making minimum-, Sérsic-Sérsic decompositions of elliptical galaxies based on photometry alone. It does not work in NGC 6166, where the profile provides physically motivated guidance in how to interpret the results. This does not argue for confidence in decompositions of giant-boxy-coreless ellipticals that are well fit by single Sérsic functions and in which monotonically decreasing profiles provide no guidance about which decompositions measure something that is physically meaningful.
Figure 16 tries a different kind of photometric decomposition that has been used to estimate the properties of cD halos. E. g., Seigar et al. (2007) and Donzelli et al. (2011) fit cD halos with exponential profiles. Since NGC 6166 is well fitted by a single Sérsic function, a Sérsic-exponential decomposition has a larger with respect to the photometric observations. It is therefore necessary to apply some additional constraint to force the program to find an exponential halo. We tried various decompositions in which the central surface brightness was constrained. All such decompositions behave similarly if we require that the RMS of the fit be consistent with measurement errors. Figure 16 shows an example in which the exponential is forced to have a central surface brightness of 25 mag arcsec. The fit RMS = 0.052 mag arcsec is worse than RMS = 0.037 mag arcsec in Figure 9 but is not excluded by the data. However, this halo is much too faint. The main galaxy contributes essentially all the light at radii where we have kinematic data, so the dispersion profile fails to rise significantly toward the outer observed value.
Again, we conclude that Sérsic-exponential decompositions of cD galaxies – at least in the case of NGC 6166 – are not well constrained physically using photometry alone.
The “cure” is to make the two components be as separate as possible by decreasing both Sérsic indices. The resulting best fit gets worse – gets, in fact, increasingly inconsistent with the photometric measurement errors – but the fit to the dispersion profile gets better. Figure 17 shows the decompositions (two of many that we tried) that best fit . Given the crude assumptions, it makes no sense to look for further improvement; the way to get a better fit is to make a full Schwarzschild (1979, 1982) model of the photometry and the kinematics. We save this exercise for a future paper. Here, we conclude that NGC 6166 and its cD halo are more distinct than a minimum- photometric decomposition suggests.
Figure 17 shows that, to fit the profile of NGC 6166, we need to make a photometric decomposition that does not minimize . This is no disaster: We chose Sérsic functions for each component, and our experience that they fit non-cD ellipticals well (KFCB) may not be relevant here.
Support for our photometric kinematic decomposition is provided by the result in Figures 15 and 17 that the predicted agrees with the observations at radii . At larger radii, the predicted remains positive but trends toward zero. The spectra there are too noisy to provide reliable constraints.
Fig. 15–17 suggest that the main body of NGC 6166 contains 30 2 % and the cD halo contains 70 2 % of the total luminosity. The formal error is probably an underestimate.
For the assumptions made in § 3.6 to get total absolute magnitudes of out to the last photometric data point or extrapolated to infinity, the main body of NGC 6166 has or . These are essentially identical to the absolute magnitudes of M 87 and NGC 4472 in the Virgo cluster (KFCB). The cD halo of NGC 6166 has or , i. e., 0.3 – 0.6 mag brighter than the brightest galaxy in the Virgo cluster.
5. Spherical Jeans Models
Our kinematic measurements allow a detailed study of the velocity distribution of the galaxy plus stellar halo and of the total mass distribution including X-ray gas and dark matter. Orbit-superposition models (Schwarzschild 1979, 1982) are postponed to a future paper. Here, we explore the stellar velocity anisotropy using spherical Jeans models.
Figures 18 and 19 show results for Jeans-model fits to our photometry and data. We assume that dark matter (“DM”, including X-ray gas) is distributed as a non-singular pseudo-isothermal (Kormendy & Freeman 2015) or as an NFW density profile (Navarro, Frenk, & White 1996, 1997). We choose the outer, circular-orbit rotation velocity km s of massless test particles in the halo to be consistent with the cluster dispersion of km s. Next, we assume that the stars have a Kroupa (2001) initial mass function with mass-to-light , based on the metallicity and age estimated in the next section and on stellar population models of Maraston et al. (2003). Then the only free parameter left is the scale length of the NFW profile or the core radius of the isothermal. We vary this scale length until the mass density profile matches the one derived from the X-ray gas by Markevitch et al. (1999). In this way, we derive a density profile over the full radius range (Figure 19) without yet using our kinematic data on NGC 6166. Finally, we vary the velocity anisotropy as a function of radius (middle panel of Figure 18) until we reproduce the observed velocity dispersion profile (bottom panel of Figure 18). Although the isothermal sphere and the NFW DM profiles are quite different, especially at kpc, the anisotropy profiles are qualitatively similar. That is, the total density profile and the dispersion profile together determine the anisotropy profile.
The important result is observed at radii to , where rises from the galaxy value of 300 km s to the cluster value of 800 km s. In this radius range, the tangential velocity dispersion is larger than the radial one, . We were unable to change this result by varying the DM profile. The observed dispersion rises so rapidly that it is necessary to “boost” the line-of-sight component by increasing . Our conclusion that in the inner part of the cD halo of NGC 6166 is consistent with the suggestion that cD halo stars are the debris torn off of individual cluster galaxies by fast collisions (see, e. g., Puchwein et al. 2010).
In recent years, the growth of Sérsic halos of giant-core-boxy-nonrotating elliptical galaxies (Kormendy et al. 2009) has also been attributed to accumulated debris from minor mergers (e. g., Naab et al. 2009; Hopkins et al. 2010; Oser et al. 2010, 2012; Hilz et al. 2012; Hilz, Naab, & Ostriker 2013). The relationship between these halos – which manifestly belong to the galaxy – and the halo of NGC 6166 – which manifestly belongs to the cluster – is a puzzle addressed in the following sections.
At large , the data hint that . This as a preliminary result. If it is correct, it could be a sign that even at 100 kpc, we reach radii where infall from the filaments of the cosmic web affect the velocity distribution (cf. Biviano et al. 2013 and Wu et al. 2014).
The isothermal halo parameters derived here, kpc and pc for , deviate from the DM parameter correlations found by Kormendy & Freeman (2015). The DM halo of NGC 6166 is more compact (e. g., higher in projected surface density) than expected from halos of late-type galaxies. However, it is consistent with scaling relations for cluster halos (Chan 2014), and its parameters agree with those derived for Abell 2199 by Chen et al. (2007).
6. Heavy Element Abundances
Our high -ratio spectra also allow us to probe stellar population diagnostics out into the part of the cD halo where the velocity dispersion is climbing to the cluster value. In Figure 20, we use the Lick Observatory spectral line indices (Faber et al. 1985; Gorgas et al. 1993; Worthey et al. 1994; Trager et al. 1998; Lee & Worthey 2005; Lee et al. 2009) to estimate Fe abundances and [Mg/Fe] – i. e., -element – overabundances in the main body and cD halo of NGC 6166.
Overabundances with respect to solar values of elements such as Mg imply short star formation time scales. Rapid enrichment of elements follows starbursts when high-mass stars die as supernovae of Type II. Alpha elements then get diluted by Fe once there is time for lower-mass stars to die as white dwarfs and subsequently blow up as Type I supernovae. After that, [/Fe] can never be enhanced again. Therefore super-solar [/Fe] abundances imply that essentially all star formation was completed in 1 Gyr. (Worthey et al. 1992; Terndrup 1993; Matteucci 1994; Bender & Paquet 1995; Thomas et al. 1999, 2002, 2005).
Kormendy et al. (2009) show that [/Fe] (over)abundance participates in the dichotomy (see Kormendy & Bender 1996; Kormendy 2009 for brief reviews) between giant, nonrotating, anisotropic ellipticals that have boxy isophotes and cuspy cores and lower-luminosity ellipticals that rotate enough to be more nearly isotropic and that have disky isophotes and (in general) central extra light components. They argue that rotating-coreless-disky ellipticals formed via at least one wet merger in which a starburst constructed the central extra component. And they argue that nonrotating-core-boxy ellipticals – which are embedded in large amounts of X-ray-emitting gas – formed most recently via dry major mergers (plus, we now believe, minor-merger addition of outer halos), protected from late star formation by their X-ray gas halos. Kormendy et al. (2009) found that [/Fe] is enhanced in nonrotating-core-boxy ellipticals but not in rotating-coreless-disky ellipticals (cf. Thomas et al. 2005, 2010). Essentially all star formation was completed very early in these galaxies. NGC 6166 is a giant core elliptical (Figures 9 – 13).
This machinery provides a partial test of our picture that cD halos consist of tidal debris torn from cluster galaxies. If [/Fe] is super-solar in the main body of NGC 6166 but near-solar in its cD halo and in smaller cluster galaxies, then this strongly supports the idea that cD halos consist of tidal debris. In contrast, if [/Fe] is super-solar in both the main body and the cD halo of NGC 6166, then this is consistent with our picture but does not prove it. Rather, that result is interesting because it suggests that star formation was switched off early in all galaxies that contribute to any part of NGC 6166. If so, then this result predicts that many (not necessarily all) smaller galaxies in the cluster are [/Fe] enhanced, too. We do not have such data. But if spectroscopy of the smaller galaxies shows that they have solar [/Fe] abundances whereas the cD halo has super-solar [/Fe], then this argues against our picture and instead supports a picture in which all of the cD including its halo forms early via some special process. We carry out the first part of the test, measuring only NGC 6166.
Figure 20 shows our measurements in NGC 6166 of the Fe mean equivalent width versus that of Mg b. The iron lines used are Fe 5270 and 5335 Å. Colors encode radii whose corresponding velocity dispersions are given in the key. Thus, the red and orange points are dominated by light from the central galaxy, whereas the green point and especially the blue point increasingly measure stars in the cluster- cD halo.
Also shown are black points at specific metallicities and population ages (lower key) for three [/Fe] abundance ratios. The points are connected by solid lines for ages of 10 Gyr and by dashed lines for ages of 3 Gyr. The models are from Thomas, Maraston, & Bender (2003); Maraston et al. (2003); and Thomas & Maraston (2003).
We conclude that the central, km s parts of NGC 6166 are old and substantially more metal-rich than solar. They have [/Fe] . These observations are consistent with the E – E dichotomy as discussed in KFCB.
At radii to 18, where begins to rise, the abundance is more nearly solar but [/Fe] remains high.
In the inner cD halo, where rising indicates that we see substantial (green point) and mostly (blue point) cluster halo ( km s) stars, the metallicity remains at least as high as at intermediate radii and [/Fe] remains at 0.3. This is consistent with but does not prove that the cD halo consists of tidally liberated galaxy debris.
Similar tests have been carried out in normal Es (e. g., Coccato et al. 2010). Greene et al. (2012, 2013) study 33 ellipticals with central km s, not quite high enough to single out core galaxies. Quoting from the latter paper : “the typical star at is old (10 Gyr), relatively metal-poor ([Fe/H] ), and -element enhanced ([Mg/Fe] ). … Stars at large radii have different abundance ratio patterns from stars in the center of any present-day galaxy, but are similar to average Milky Way thick disk stars. Our observations are consistent with a picture in which the stellar outskirts are built up through minor mergers with disky galaxies whose star formation is truncated early ( 1.5 – 2).”
7. Evolutionary History of NGC 6166 and Abell 2199
If the cD halo of NGC 6166 had its star formation quenched in 1 Gyr, then the environs of NGC 6166 have been special for a long time. This has implications for cD formation:
In recent years, there has been a substantial convergence from many lines of research on a consistent and plausible picture of what quenches star formation in general and especially in giant galaxies such as NGC 6166. The essential idea is often called “ quenching:” a total galaxy or cluster mass is required to hold gravitationally onto large amounts of hot, X-ray-emitting gas, and the hot gas quenches star formation. Essentially equivalent pictures have been reached (1) via theoretical studies of cosmological gas accretion onto large potential wells (Dekel & Birnboim 2006, 2008); (2) via semi-analytic modeling (Cattaneo et al. 2006, 2008, 2009); (3) via studies of galaxies in the high-redshift universe (e. g., Faber et al. 2007; Peng et al. 2010; Knobel et al. 2014); (4) via studies of physical differences between the two kinds of elliptical galaxies (KFCB; Kormendy & Bender 2012), and (5) via studies of AGN feedback in relation to the demographics of supermassive black holes and the properties of their host galaxies (Kormendy & Ho 2013). Note that the value of is somewhat higher at higher because of higher cold gas fractions there (see the Dekel & Birnboim papers).
Peng et al. (2010) provide the clearest description: They distinguish mass-driven quenching from environmentally-driven quenching and quenching related to bulge formation. Like Knobel et al. (2014), we suggest that mass-driven and environmentally-driven quenching are fundamentally the same process; in mass-driven quenching, the quenched galaxy owns its own hot gas, whereas in environmentally-driven quenching of satellite galaxies, the gas that does the work belongs to the parent giant galaxy or cluster. Both processes together are equivalent to the “maintenance-mode AGN feedback” discussed in Kormendy & Ho (2013). Again, the quenching is done by the hot gas, and the process that keeps it hot (AGN feedback is one possibility) is somewhat secondary. Quenching by hot gas is the essential process that is relevant here. (Peng’s “quenching associated with bulge formation” is equivealent to Kormendy & Ho’s “quasar-mode feedback”.)
The X-ray halo needed for quenching is present in Abell 2199 (e. g., Markevitch et al. 1999; Johnstone et al. 2002; Kawaharada et al. 2010). However, the implications of our results are broader than this:
In general, we expect that a cluster grows as galaxies and galaxy groupings fall into it that are sufficiently sub- to have had prolonged star formation histories. As they and their stars get added to NGC 6166, it is natural to expect that the resulting halo would not be as -element enhanced as the main body of the galaxy. Simulations suggest that cD halo stars are somewhat older than typical stars in the galaxies that contribute to the halo (Murante et al. 2004; Puchwein et al. 2010). Also, simulations by Murante et al. (2007) suggest that the inner parts of cD halos – this certainly includes the parts of NGC 6166 that we have measured – “come from the [merger] family tree of the [parent galaxy]”; that is, from galaxies that share the immediate history of the central galaxy. And simulators agree that the halo tends to be contributed by the most massive cluster galaxies; their star rofmation was presumably quenched early. Still, if even the debris halo of NGC 6166 is -element enhanced, then this suggests that the environs of the galaxy – including that of the progenitors that contributed to its cD halo – constituted a deep enough gravitational potential well to allow star formation to be quenched rapidly. And this suggests a solution to the following puzzle:
Why does NGC 6166 have such a high-surface-brightness halo of intracluster stars when apparently richer and denser clusters such as Coma have weaker cD characteristics? Note that the velocity dispersion has already risen significantly in NGC 6166 at (Figures 4 and 5), where the surface brightness is 22.5 mag arcsec (Table 3). Evidently the processes freed the intracluster stars happened less strongly or for a shorter time in Coma than in Abell 2199. Why?
Coma may have formed relatively recently – is, in fact, still forming now, with the imminent accretion of the NGC 4839 grouping. In contrast, Abell 2199 looks less dense than Coma does now, but the central few hundred kpc evidently has been a massive enough environment to allow the early quenching of star formation. It may also have been dense enough to allow cD halo formation processes to operate efficiently for an unusually long time.
A more speculative remark follows from the large core radius of NGC 6166 (Figure 13). There is a tight correlation between the light and mass “deficit” that defines the core phenomenon and the measured mass of supermassive black holes (Kormendy & Bender 2009). The canonical interpretation is that cores are created when supermassive black hole binaries produced in major, dry mergers fling stars away from the center as they decay toward an eventual merger (e. g., Ebisuzaki et al. 1991; Faber et al. 1997; Milosavljević & Merritt 2001; Milosavljević et al. 2002; Merritt 2006). If the –core correlation is valid for NGC 6166, then the core light deficit corresponds to a BH mass . The core radius is unusually large, but the core surface brightness is unusually small. So the light deficit and are almost the same as those of M 87. Still, Abell 2199 is one of the most plausible environments in which episodic AGN feedback could help to keep its hot gas hot (Fabian 2012). And the early quenching of star formation together with the long history of cluster dynamical evolution may be connected with the unusual properties (large radius but low surface brightness) of the core of NGC 6166.
8. Implications for cD Formation Mechanisms
This observational paper does not fully review the large literature on possible formation mechanisms for cD galaxies. We restrict ourselves to the most basic conclusions from our new results and concentrate on formation of the cD halo. Suggested mechanisms are divided into three categories:
8.1. Star Formation in Cooling Flows in X-Ray Gas
Are cD halos made of stars that rain out of cooling flows in hot gas (see Fabian et al. 1991; Fabian 1994 for reviews)? This idea was entertained in the heyday of the cooling-flow problem, when we observed large amounts of X-ray-emitting, hot gas in clusters but could not measure temperature profiles. Absent heating processes, hot-gas cooling times near the centers of many clusters and individual galaxies are short. In clusters, – yr of baryons should rain out of the hot gas, presumably by star formation. To escape detection, the initial mass function would have to be truncated above 1 (Fabian et al. 1991). We have never directly observed such star formation in any environment (Bastian et al. 2010).
This possibility is now regarded as a non-starter. The main reason is that we now can measure gas temperature profiles, and we find that temperatures decrease only modestly to a floor at keV. In particular, we do not see the strong emission lines from Fe XVII that would be our signal that gas has cooled below 0.7 keV (see Fabian 2012 for review). So the cooling flow problem has morphed into a different question: What keeps the gas hot? At least three heating processes are hard to avoid. Most popular is heating by AGN feedback (Fabian 2012; Kormendy & Ho 2013; and Heckman & Best 2014 provide reviews). Also, gas from the cosmological web that falls into objects with masses accelerates so much that a shock forms where it impacts the static intergalactic or intracluster medium; this heats the hot gas from the outside inward (Birnboim & Dekel 2003; Kereš et al. 2005; Dekel & Birnboim 2006, 2008). This is an aspect of quenching of star formation. Finally, dying stars eject large amounts of mass into the intracluster medium at the kinetic temperatures of stars in galaxies and galaxies in clusters (e. g., Ostriker 2006). All three mechanisms are likely to be important. In this picture, episodic cooling fuels the AGN and switches it on long enough to allow it to keep the center of the hot gas hot (Fabian 2012). Small amounts of star formation may be connected with these events, and small amounts of star formation are seen in brightest cluster galaxies (e. g., Liu et al. 2012). But no compelling argument suggests that large amounts of star formation occur in clusters at radii where we see cD halos. Also, our observation that the cD halo of NGC 6166 is -element enriched precludes the idea that prolonged, in-situ star formation made a significant fraction of the light that we see in the halo.
8.2. Processes Intrinsic to the Origin of the Central Galaxy
Do cD halos originate as an integral part of the formation of the central galaxy? For example, could a specialized history of galaxy mergers make both the central and halo parts of a cD galaxy together?
Our phrasing is somewhat different from the question that dominated work on brightest cluster galaxies (“BCGs”) in the 1970s – 1990s (see Tremaine 1990 for a review). Then, the emphasis was on observational hints that BCGs in general (i. e., including but not limited to cDs) are inconsistent with statistical expectations based on the luminosity functions of fainter galaxies in the cluster. If with characteristic luminosity (Schechter 1976), then BCGs with are statistically too bright to be drawn from the populations of other galaxies in the clusters (see Figure 1 in Binggeli 1987 for an evocative illustration). In many papers, cDs and non-cD BCGs were discussed together. Given the observation that cD halos are approximately as bright as or brighter than the central parts of the galaxies (e. g., Seigar et al. 2007), this essentially ensures that BCGs as a class will look especially luminous (Tremaine & Richstone 1977).
As some authors have done since the beginning of this subject, we differentiate between the main bodies of cDs and their halos. In NGC 6166, we separate them operationally as having 300 km s and km s, respectively. How the main bodies of BCGs form and how cD halos form may be separate questions.
When their halos are inventoried separately, it is much less obvious that the main bodies of cDs are unusual enough to imply formation physics that is different from that of other cluster galaxies. The new observations in this paper do not speak strongly to this issue, and we do not discuss it in detail. Ways in which the main body of NGC 6166 is not unusual are the subject of Sections 3.2 – 3.5. Except for its unusually large and low-surface-brightness core (discussed in the previous section), the main body of NGC 6166 is rather like M 87 (a marginal cD) but also like the other giant-core-boxy ellipticals in the Virgo cluster. Quantitative differences (Section 3.6) are mainly due to the cD halo of NGC 6166. However, we note here one additional observation that does imply something special about cD-like galaxies:
Prototypical of a compelling but mysterious phenomenon, M 87 has an unusually large number of globular clusters for its galaxy luminosity. Harris & van den Bergh (1981) introduced the specific globular cluster frequency as the number of globulars per unit absolute magnitude = of galaxy luminosity. Measurement of is tricky for many reasons (e. g., galaxy distances are uncomfortably large, so we see less deeply into cluster luminosity functions than we would like), but the conclusion that is factors of several larger for M 87 and for some other BCGs (e. g., NGC 1399: Hanes & Harris 1986; Harris & Hanes 1987; NGC 3311: Harris 1986; see Harris, Harris, & Alessi 2013 for the most recent summary) has withstood the test of time. The number of globular clusters in NGC 6166 is (Harris et al. 2013). With respect to the absolute magnitude of the main E-like part of NGC 6166, this implies that . If instead we normalize by the total luminosity including the cD halo, then . This is still slightly higher than the canonical numeber of 1 – 2 for ellipticals. As discussed, for example, in Burkert & Tremaine (2010), this is one indication that the early evolution of the objects that later assembled into these BCGs (some of which are clearly cDs and others of which are just giant ellipticals) was already special. This theme of an early, special environment in which NGC 6166 and its cD halo formed was discussed in § 7.
8.3. cD Halo Formation by Stellar-Dynamical Processes Inherent to Clusters
Our observations are most consistent with the now favored picture that cD halos are constructed by stellar-dynamical processes that are inherent to cluster evolution. The main body forms by the usual hierarchical clustering and galaxy merging, especially in smaller group precursors to present-day, rich clusters. In the process, violent relaxation splashes some stars to large radii. But the cD halo is added as a result of cluster-related processes such as the stripping of stars off of member galaxies by dynamical harassment and the cannibalism and destruction of dwarf galaxies in minor mergers. This picture was originated by Gallagher & Ostriker (1972) and by Richstone (1975, 1976) and has now been greatly elaborated in many papers, both observational (see the earlier papers on cluster background light and, e. g., Bernstein et al. 1995; Gonzalez, Zabludoff, & Zaritsky 2005; Arnaboldi et al. 2012; Montes & Trujillo 2014) and theoretical (e. g., Dubinski 1998; Murante et al. 2004, 2007; De Lucia & Blaizot 2007; Puchwein et al. 2010; and Cui et al. 2014).
8.4. Blurring the Distinction Between cD Galaxies and Elliptical Galaxies with Cores
Our observations (1) that the cD halo of NGC 6166 is more nearly at rest in Abell 2199 than is its central galaxy and (2) that this halo has the same velocity dispersion as the cluster galaxies support the idea that it consists of stars that were liberated from cluster members. The high velocity dispersion implies that the cD halo is controlled by cluster gravity. It is only by convention – and not because this is physically meaningful – that we call it the halo of NGC 6166.
On the other hand, the outer parts of NGC 6166 and the intracluster light merge seamlessly such that the brightness profile outside the central core is well described by a single Sérsic function with index . In this sense, NGC 6166 qualitatively resembles other core-boxy-nonrotating elliptical galaxies such as those studied in KFCB and emphasized in the SAURON/Atlas series of papers (see Cappellari 2015 for a review). The Sérsic halos of core-boxy-nonrotating ellipticals that are not brightest cluster galaxies manifestly belong to the galaxy – their velocity dispersions generally decrease monotonically outward.
This blurs the distinction between cDs and giant elliptical galaxies. Perhaps they are more similar than we thought. The central puzzle about both kinds of galaxies is why . In contrast, many numerical simulations of major mergers of two similar galaxies robustly show that the scrambled-up remnants of the stars that were already present before the mergers have Sérsic profiles with (e. g., van Albada 1982; Mihos & Hernquist 1994; Springel & Hernquist 2005; Naab & Trujillo 2006; Hopkins et al. 2009a, b). These are precisely the Sérsic indices observed for coreless-disky-rotating ellipticals, which are thought to be formed in wet mergers during which starbursts grew the central extra light components (see Kormendy 1999 and KFCB for observations and review and Hopkins et al. 2009a for the most detailed simulations).
Maybe the main difference between cDs and core-boxy-nonrotating (but not cD) ellipticals is the degree to which clusters are dynamically old enough to have liberated enough stars from individual galaxies to make a detectable intracluster population. It may also matter whether the large- halos formed in subgroups such that the central galaxy controls their dynamics or conversely in high-, rich clusters at radii controlled by the cluster rather than the central galaxy. An important goal of future work is to explore the reasons why cD galaxies and core-boxy-nonrotating ellipticals look so similar when their halo velocity dispersions point to significant differences in formation history.
- affiliation: Based on observations obtained with the Hobby-Eberly Telescope, which is a joint project of the University of Texas at Austin, the Pennsylvania State University, Stanford University, Ludwig-Maximilians-Universität München, and Georg-August-Universität Göttingen. Submitted to ApJ.
- affiliation: Max-Planck-Institut für Extraterrestrische Physik, Giessenbachstrasse, D-85748 Garching-bei-München, Germany; email@example.com
- affiliation: Universitäts-Sternwarte, Scheinerstrasse 1, München D-81679, Germany
- affiliation: Department of Astronomy, University of Texas at Austin, 1 University Station C1400, Austin, Texas 78712-0259; firstname.lastname@example.org
- affiliation: Max-Planck-Institut für Extraterrestrische Physik, Giessenbachstrasse, D-85748 Garching-bei-München, Germany; email@example.com
- affiliation: Universitäts-Sternwarte, Scheinerstrasse 1, München D-81679, Germany
- affiliation: Department of Astronomy, University of Texas at Austin, 1 University Station C1400, Austin, Texas 78712-0259; firstname.lastname@example.org
- affiliation: Present Address: MIT Lincoln Laboratory, ETS Field Site, P. O. Box 1707, Socorro, NM 87801; Mark.Cornell@ll.mit.edu
- affiliation: Department of Astronomy, University of Texas at Austin, 1 University Station C1400, Austin, Texas 78712-0259; email@example.com
- affiliation: Present Address: Centre for Astrophysics and Supercomputing, Swinburne University of Technology, Mail Stop H30, P. O. Box 218, Hawthorn, Victoria 3122, Australia; firstname.lastname@example.org
- The name “cD” has created some confusion. It has been interpreted to mean “cluster dominant” or “central dominant” or “central diffuse”. All are correct descriptions, but they are not the origin of the name. Morgan (1958) introduced the “D” form classification for galaxies that are like ellipticals but with distinct, outer halos with shallow brightness gradients. The “D” class has not been as useful as Hubble classes (Hubble 1936; Sandage 1961), because it includes several different physical phenomena, (a) S0 galaxies, in which the outer halo is the disk; (b) giant ellipticals with high Sérsic (1968) indices , and (c) the subjects of this paper: giant ellipticals whose distinct outer halos consist of intracluster stars that have been stripped from cluster galaxies. Because this involves important physics, the name “cD” has survived even though the name “D” has not. But “c” does not mean “central” or “cluster”. Rather, it is a historical anachronism that survives from stellar spectral classes that are no longer used. Quoting Mathews et al. (1964): “These very large D galaxies observed in clusters are given the prefix ‘c’ in a manner similar to the notation for supergiant stars in stellar spectroscopy.”
- Andrade-Santos, F., Nulsen, P. E. J., Kraft, R. P., et al. 2013, ApJ, 766, 107
- Arnaboldi, M. 2011, Paper presented at the ESO Workshop on Fornax, Virgo, Coma et al.: Stellar Systems in High Density Environments, ed. M. Arnaboldi, http://www.eso.org/sci/meetings/2011/ fornax_virgo2011/talks_pdf/Arnaboldi_Magda.pdf
- Arnaboldi, M., Aguerri, J. A. L., Napolitano, N. R., et al. 2002, AJ, 123, 760
- Arnaboldi, M., Freeman, K. C., Mendez, R. H., et al. 1996, ApJ, 472, 145
- Arnaboldi, M., & Gerhard, O. 2010, Highlights Astron., 15, 97
- Arnaboldi, M., Gerhard, O., Aguerri, J. A. L., et al. 2004, ApJ, 614, L33
- Arnaboldi, M., Ventimiglia, G., Iodice, E., Gerhard, O., & Coccato, L. 2012, A&A, 545, A37
- Baggett, W. E., Baggett, S. M., & Anderson, K. S. J. 1998, AJ, 116, 1626
- Bastian, N., Covey, K. R., & Meyer, M. R. 2010, ARA&A, 48, 339
- Beers, T. C., & Geller, M. J. 1983, ApJ, 274, 491
- Bender, R. 1987, Mitt. Astr. Gesellschaft, No. 70, 226
- Bender, R. 1988, A&A, 193, L7
- Bender, R. 1990, A&A, 229, 441
- Bender, R., Burstein, D., & Faber, S. M. 1992, ApJ, 399, 462
- Bender, R., Burstein, D., & Faber, S. M. 1993, ApJ, 411, 153
- Bender, R., Döbereiner, S., & Möllenhoff, C. 1987, A&A, 177, L53
- Bender, R., Döbereiner, S., & Möllenhoff, C. 1988, A&AS, 74, 385
- Bender, R., & Möllenhoff, C. 1987, A&A, 177, 71
- Bender, R., & Paquet, A. 1995, in IAU Symposium 164, Stellar Populations, ed. P. C. van der Kruit & G. Gilmore (Dordrecht: Kluwer), 259
- Bender, R., Saglia, R. P., & Gerhard, O. E. 1994, MNRAS, 269, 785
- Bender, R., Surma, P., Döbereiner, S., Möllenhoff, C., & Madejsky, R. 1989, A&A, 217, 35
- Bernstein, G. M., Nichol, R. C., Tyson, J. A., Ulmer, M. P., & Wittman, D. 1995, AJ, 110, 1507
- Binggeli, B. 1987, in Nearly Normal Galaxies: From the Planck Time to the Present, ed. S. M. Faber (New York: Springer), p. 195
- Birnboim, Y., & Dekel, A. 2003, MNRAS, 345, 349
- Biviano, A., Rosati, P., Balestra, I., et al. 2013, A&A, 558, A1
- Briel, U. G., Henry, J. P., Lumb, D. H., et al. 2001, A&A, 365, L60
- Buitrago, F., Trujillo, I., Conselice, C. J., et al. 2008, ApJ, 687, L61
- Burbidge, E. M. 1962, ApJ, 136, 1134
- Burkert, A., & Tremaine, S. 2010, ApJ, 720, 516
- Byun, Y.-I., Grillmair, C. J., Faber, S. M., et al. 1996, AJ, 111, 1889
- Cappellari, M. 2015, ARA&A, in preparation
- Carter, D., Bridges, T. J., & Hau, G. K. T. 1999, MNRAS, 307, 131
- Carter, D., Efstathiou, G., Ellis, R. S., Inglis, I., & Godwin, J. 1981, MNRAS, 195, 15P
- Carter, D., Inglis, I., Ellis, R. S., Efstathiou, G., & Godwin, J. G. 1985, MNRAS, 212, 471
- Castro-Rodriguéz, N., Arnaboldi, M., Aguerri, J. A. L., et al. 2009, A&A, 507, 621
- Cattaneo, A., Dekel, A., Devriendt, J., Guideroni, B., & Blaizot, J. 2006, MNRAS, 370, 1651
- Cattaneo, A., Dekel, A., Faber, S. M., & Guideroni, B. 2008, MNRAS, 389, 567
- Cattaneo, A., Faber, S. M., Binney, J., et al. 2009, Nature, 460, 213
- Chan, M. H., 2014, MNRAS 442, L14
- Chen, Y., Reiprich, T. H., Böhringer, H., Ikebe, Y., & Zhang, Y.-Y., 2007, A&A, 466, 805
- Chiboukas, K., Karachentsev, I. D., & Tully, R. B. 2009, AJ, 137, 3009
- Churazov, E., Forman, W., Vikhlinin, A, et al. 2008, MNRAS, 388, 1062
- Coccato, L., Gerhard, O., & Arnaboldi, M. 2010, MNRAS, 407, L26
- Côté, P., McLaughlin, D. E., Hanes, D. A., et al. 2001, ApJ, 559, 828
- Coziol, R., Andernach, H., Caretta, C. A., Alamo-Martínez, K. A., & Tago, E. 2009, AJ, 137, 4795
- Cui, W., Murante, G., Monaco, P., et al. 2014, MNRAS, 437, 816
- Dekel, A., & Birnboim, Y. 2006, MNRAS, 368, 2
- Dekel, A., & Birnboim, Y. 2008, MNRAS, 383, 119
- De Lucia, G., & Blaizot, J. 2007, MNRAS, 375, 2
- Djorgovski, S. 1992, in Morphological and Physical Classification of Galaxies, ed. G. Longo, M. Capaccioli, & G. Busarello (Dordrecht: Kluwer), 337
- Djorgovski, S., & Davis, M. 1987, ApJ, 313, 59
- Djorgovski, S., de Carvalho, R., & Han, M.-S. 1988, in The Extragalactic Distance Scale, ed. S. van den Bergh & C. J. Pritchet (San Francisco: ASP), 329
- Doherty, M., Arnaboldi, M., Das, P., et al. 2009, A&A, 502, 771
- Donzelli, C. J., Muriel, H., & Madrid, J. P. 2011, ApJS, 195, 15
- Dressler, A. 1979, ApJ, 231, 659
- Dressler,A., Lynden-Bell, D., Burstein, D., et al. 1987, ApJ, 313, 42
- Dubinski, J. 1998, ApJ, 502, 141
- Ebisuzaki, T., Makino, J., & Okamura, S. K. 1991, Nature, 354, 212
- Faber, S. M., Friel, E. D., Burstein, D., & Gaskell, C. M. 1985, ApJS, 57, 711
- Faber, S. M., et al. 1987, in Nearly Normal Galaxies: From the Planck Time to the Present, ed. S. M. Faber (New York: Springer), 175
- Faber, S. M., Tremaine, S., Ajhar, E. A., et al. 1997, AJ, 114, 1771
- Faber, S. M., Willmer, C. N. A., Wolf, C., et al. 2007, ApJ, 665, 265
- Fabian, A. C. 1994, ARA&A, 32, 277
- Fabian, A. C. 2012, ARA&A, 50, 455
- Fabian, A. C., Nulsen, P. E. J., & Canizares, C. R. 1991, A&AR, 2, 191
- Ferrarese, L., Côté, P., Jordán, A., et al. 2006, ApJS, 164, 334
- Fisher, D., Illingworth, G., & Franx, M. 1995, ApJ, 438, 539
- Fisher, D. B., & Drory, N. 2008, AJ, 136, 773
- Foster, C., Spitler, L. R., Romanowsky, A. J., et al. 2011, MNRAS, 415, 3393
- Ebisuzaki, T., Makino, J., Okamura, S. K. 1991, Nature, 354, 212
- Gallagher J. S., & Ostriker J. P. 1972, AJ, 77, 288
- Gavazzi, G., Donati, A., Cucciati, O., et al. 2005, A&A, 430, 411
- Gavazzi, G., Franzetti, P., Scodeggio, M., Boselli, A., & Pierini, D. 2000, A&A, 361, 863
- Gebhardt, K., Adams, J., Richstone, D. et al. 2011, ApJ, 729, 119
- Gebhardt, K., & Thomas, J. 2009, ApJ, 700, 1690
- Gebhardt, K., et al. 1996, AJ, 112, 105
- Gerhard, O., Arnaboldi, M., Freeman, K. C. , et al. 2007, A&A, 468, 815
- Gerhard, O., Kronawitter, A., Saglia, R. P., & Bender, R. 2001, AJ, 121, 1936
- Gonzalez, A. H., Zabludoff, A. I., & Zaritsky, D. 2005, ApJ, 218, 195
- Gorgas, J., Faber, S. M., Burstein, D., et al. 1993, ApJS, 86, 153
- Greene, J. E., Murphy, J. D., Comerford, J. M., Gebhardt, K., & Adams, J. J. 2012, ApJ, 750, 32
- Greene, J. E., Murphy, J. D., Graves, G. J., et al. 2013, ApJ, 776, 64
- Hanes, D. A., & Harris, W. E. 1986, ApJ, 309, 564
- Harris, W. E. 1986, AJ, 91, 822
- Harris, W. E., & Hanes, D. A. 1987, AJ, 93, 1368
- Harris, W. E., Harris, G. L. H., & Alessi, M. 2013, ApJ, 772, 82
- Harris, W. E., & van den Bergh, S. 1981, AJ, 86, 1627
- Heckman, T. M., & Best, P. N. 2014, ARA&A, 52, 589
- Hill, G. J., Nicklas, H. E., MacQueen, P. J., Tejada, C., Cobos Duenas, F. J., & Mitsch, W. 1998, Proc. SPIE, 3355, 375
- Hilz, M., Naab, T., & Ostriker, J. P. 2013, MNRAS, 429, 2924
- Hilz, M., Naab, T., Ostriker, J. P., et al. 2012, MNRAS, 425, 3119
- Hopkins, P. F., Bundy, K., Hernquist, L., Wuyts, S., & Cox, T. J. 2010, MNRAS, 401, 1099
- Hopkins, P. F., Cox, T. J., Dutta, S. N., et al. 2009a, ApJS, 181, 135
- Hopkins, P. F., Lauer, T. R., Cox, T. J., Hernquist, L., & Kormendy, J. 2009b, ApJS, 181, 486
- Huang, S., Ho, L. C., Peng, C. Y., Zhao-Yu, L., & Barth, A. J. 2013a, ApJ, 766, 47
- Huang, S., Ho, L. C., Peng, C. Y., Zhao-Yu, L., & Barth, A. J. 2013b, ApJ, 768, L28
- Hubble, E. 1936, The Realm of the Nebulae (New Haven: Yale University Press)
- Jardel, J. R., Gebhardt, K., Shen, J., et al. 2011, ApJ, 739, 21
- Johansson, P. H., Naab, T., & Ostriker, J. P. 2012, ApJ, 754, 115
- Johnstone, R. M., Allen, S. W., Fabian, A. C., & Sanders, J. S. 2002, MNRAS, 336, 299
- Kawarharada, M., Makishima, K., Kitaguchi, T., et al. 2010, PASJ, 62, 115
- Kelson, D. D., Zabludoff, A. I., Williams, K. A., et al. 2002, ApJ, 576, 720
- Kereš, D., Katz, N., Weinberg, D. H., & Davé, R. 2005, MNRAS, 363, 2
- Knobel, C., Lilly, S. J., Woo, J., & Kovač, K. 2014, ApJ, submitted (arXiv:1408.2553)
- Komatsu, E., et al. 2009, ApJS,180, 330
- Kormendy, J. 1999, in Galaxy Dynamics: A Rutgers Symposium, ed. D. Merritt, J. A. Sellwood, & M. Valluri (San Francisco: ASP), 124
- Kormendy, J. 2009, in ASP Conf. Series 419, Galaxy Evolution: Emerging Insights and Future Challenges, ed. S. Jogee, I. Marinova, L. Hao, & G. A. Blanc (San Francisco, CA: ASP), 87
- Kormendy, J., & Bender, R. 1996, ApJ, 464, L119
- Kormendy, J., & Bender, R. 2009, ApJ, 691, L142
- Kormendy, J., & Bender, R. 2012, ApJS, 198, 2
- Kormendy, J., & Djorgovski, S. 1989, ARA&A, 27, 235
- Kormendy, J., Fisher, D. B., Cornell, M. E., & Bender, R. 2009, ApJS, 182, 216 (KFCB)
- Kormendy, J., & Freeman, K. C. 2015, ApJ, submitted
- Kormendy, J., & Ho, L. C. 2013, ARA&A, 51, 511
- Kormendy, J., et al. 1994, in ESO/OHP Workshop on Dwarf Galaxies, ed. G. Meylan & P. Prugniel (Garching: ESO), 147
- Kosyra, R., Gössl, C., Hopp, U., et al. 2014, Exp. Astron., DOI 10.1007/s10686-014-9414-1 (Springer) (arXiv:1408.2519)
- Kronawitter, A., Saglia, R. P., Gerhard, O., & Bender, R. 2000, A&AS, 144, 53
- Lachièze-Rey, M., Vigroux, L., & Souviron, J. 1985, A&A, 150, 62
- Lauer, T. R. 1985, ApJS, 57, 473
- Lauer, T. R. 1986, ApJ, 311, 34
- Lauer, T. R., Ajhar, E. A., Byun, Y.-I., et al. 1995, AJ, 110, 2622
- Lauer, T. R., Gebhardt, K., Faber, S. M., et al. 2007, ApJ, 664, 226
- Lauer, T. R., Postman, M., Strauss, M. A., Graves, G. J., & Chisari, N. E. 2014, ApJ, arXiv:1407.2260
- Lee, H.-C., & Worthey, G. 2005, ApJS, 160, 176
- Lee, H.-C., Worthey, G., Dotter, A., et al. 2009, ApJ, 694, 902
- Liu, F. S., Mao, S., & Meng, X. M. 2012, MNRAS, 423, 422
- Loubser, S. I., Sansom, A. E., Sánchez-Blázquez, P., Soechting, I. K., & Bromage, G. E. 2008, MNRAS, 391, 1009
- Lucy, L. B. 1974, AJ, 79, 745
- Maraston, C., Greggio, L., Renzini, A., et al. 2003, A&A, 400, 823
- Markevitch, M., Vikhlinin, A., Forman, W. R., & Sarazin, C. L. 1999, ApJ, 527, 545
- Martin, N. F., de Jong, J. T. A., & Rix, H.-W. 2008, ApJ, 684, 1075
- Mateo, M. 1998, ARA&A, 36, 435
- Matteucci, F. 1994, A&A, 288, 57
- Matthews, T. A., Morgan, W. W., & Schmidt, M. 1964, ApJ, 140, 35
- McConnachie, A. W., & Irwin, M. J. 2006, MNRAS, 365, 1263
- Merritt, D. 1983, ApJ, 264, 24
- Merritt, D. 2006, ApJ,, 648, 976
- Mihos, J. C. 2011, Paper presented at the ESO Workshop on Fornax, Virgo, Coma et al.: Stellar Systems in High Density Environments, ed. M. Arnaboldi, http://www.eso.org/sci/meetings/2011/fornax_virgo2011/ talks_pdf/Mihos_Chris.pdf
- Mihos, J. C., Harding, P., Feldmeier, J., & Morrison, H. 2005, ApJ, 631, L41
- Mihos, J. C., Janowiecki, S., Feldmeier, J. J., Harding, P., & Morrison, H. 2009, ApJ, 698, 1879
- Mihos, J. C., & Hernquist, L. 1994, ApJ, 437, L47
- Milosavljević, M., & Merritt, D. 2001, ApJ, ApJ, 563, 34
- Milosavljević, M., Merritt, D., Rest, A., & van den Bosch, F. C. 2002, MNRAS, 331, L51
- Minkowski, R. 1961, AJ, 66, 558
- Montes, M., & Trujillo, I. 2014, ApJ, 794, 137
- Morgan, W. W. 1958, PASP, 70, 364
- Morgan, W. W., & Lesh, J. R. 1965, ApJ, 142, 1364
- Murante, G., Arnaboldi, M., Gerhard, O., et al. 2004, ApJ, 607, L83
- Murante, G., Giovalli, M., Gerhard, O., et al. 2007, MNRAS, 377, 2
- Murphy, J. D., Gebhardt, K., & Adams, J. J. 2011, ApJ, 729, 129
- Murphy, J. D., Gebhardt, K., & Cradit, M. 2014, ApJ, 785, 143
- Naab, T., Johansson, P. H., & Ostriker, J. P. 2009, ApJ, 699, L178
- Naab, T., & Trujillo, I. 2006, MNRAS, 369, 625
- Navarro, J. F., Frenk, C. S., & White, S. D. M. 1996, ApJ, 462, 563 (NFW)
- Navarro, J. F., Frenk, C. S., & White, S. D. M. 1997, ApJ, 490, 493 (NFW)
- Neumann, D. M., Arnaud, M., Gaustad, R., et al. 2001, A&A, 365, L74
- Neumann, D. M., Lumb, D. H., Pratt, G. W., & Briel, U. G. 2003, A&A, 400, 411
- Nowak, N., Saglia, R. P., Thomas, J., et al. 2008, MNRAS, 391, 1629
- Oegerle, W. R., & Hill, J. M. 2001, AJ, 122, 2858
- Oemler, A. 1976, ApJ, 209, 693
- Oser, L., Naab, T., Ostriker, J. P., & Johansson, P. H. 2012, ApJ, 744, 63
- Oser, L., Ostriker, J. P., Naab, T., Johansson, P. H., & Burkert, A. 2010, ApJ, 725, 2312
- Ostriker, J. P. 2006, Paper Presented at the Thinkshop on The Role of Black Holes in Galaxy Formation and Evolution, Potsdam, Germany, 2006 September 10–13 (see Cattaneo et al. 2009)
- Ostriker, J. P., & Tremaine, S. D. 1975, ApJ, 202, L113
- Papovich, C., Bassett, R., Lotz, J. M., et al. 2012, ApJ, 750, 93
- Peng, Y.-J., Lilly, S. J., Kovač, K., et al. 2010, ApJ, 721, 193
- Pimbblet, K. A., Rosebloom, I. G., & Doyle, M. T. 2006, MNRAS, 368, 651
- Proctor, R. N., Forbes, D. A., Romanowsky, A. J., et al. 2009, MNRAS, 398, 91
- Puchwein, E., Springel, V., Sijacki, D., & Dolag, K. 2010, MNRAS, 406, 936
- Quintana, H., & Lawrie, D. G. 1982, AJ, 87, 1
- Raskutti, S., Greene, J. E., & Murphy, J. D. 2014, ApJ, 786, 23
- Richardson, W. H. 1972, JOSA, 62, 52
- Richstone D. O., 1975, ApJ, 200, 535
- Richstone D. O., 1976, ApJ, 204, 642
- Richstone D. O., & Malumuth, E. M. 1983, ApJ, 268, 30
- Rines, K., Geller, M. J., Diaferio, A., et al. 2002, AJ, 124, 1266
- Sandage, A. 1961, The Hubble Atlas of Galaxies (Washington: Carnegie Institution of Washington)
- Savage, B. D., & Mathis, J. S. 1979, ARA&A, 17, 73
- Schechter, P. 1976, ApJ, 203, 297
- Schlegel, D. J., Finkbeiner, D. P., & Davis, M. 1998, ApJ, 500, 525
- Schombert, J. M. 1986, ApJS, 60, 603
- Schombert, J. M. 1987, ApJS, 64, 643
- Schombert, J. M. 1988, ApJ, 328, 475
- Schwarzschild, M. 1979, ApJ, 232, 236
- Schwarzschild, M. 1982, ApJ, 263, 599
- Seigar, M. S., Graham, A. W., & Jerjen, H. 2007, MNRAS, 378, 1575
- Sembach, K. R., & Tonry, J. L. 1996, AJ, 112, 797
- Sérsic, J. L. 1968, Atlas de Galaxias Australes (Cordoba: Observatorio Astronomico, Universidad de Cordoba)
- Simionescu, A., Werner, N., Urban, O., et al. 2013, ApJ, 775, 4
- Springel, V., & Hernquist, L. 2005, ApJ, 622, L9 (arXiv:astro-ph/041379)
- Szomoru, D., Franx, M., & van Dokkum, P. G. 2012, ApJ, 749, 121
- Stover, R. J. 1988, in Instrumentation for Ground-Based Optical Astronomy: Present and Future, ed. L. B. Robinson (New York: Springer-Verlag), 443
- Terndrup, D. M. 1993, in The Minnesota Lectures on the Structure and Dynamics of the Milky Way, ed. R. M. Humphreys (San Francisco, CA: ASP), 9
- Thomas, D., Greggio, L., & Bender, R. 1999, MNRAS, 302, 537
- Thomas, D., & Maraston, C. 2003, A&A, 401, 429
- Thomas, D., Maraston, C., & Bender, R. 2002, ApSpSci, 281, 371
- Thomas, D., Maraston, C., & Bender, R. 2003, MNRAS, 339, 897
- Thomas, D., Maraston, C., Bender, R., & Mendez de Oliveira, C. 2005, ApJ, 621, 673
- Thomas, D., Maraston, C., Schawinski, K., Sarzi, M., & Silk, J. 2010, MNRAS, 404, 1775
- Tonry, J. L. 1984, ApJ, 279, 13
- Tonry, J. L. 1985, AJ, 90, 2431
- Trager, S. C., Worthey, G., Faber, S. M., Burstein, D., & González, J. J. 1998, ApJS, 116, 1
- Tremaine, S. 1990, in Dynamics and Interactions of Galaxies, ed. R. Wielen (Berlin: Springer), p. 394
- Tremaine, S., & Richstone, D. O. 1977, ApJ, 212, 311
- van Albada, T. S. 1982, MNRAS, 201, 939
- van Dokkum, P. G., Whitaker, K. E., Brammer, G., et al. 2010, ApJ, 709, 1018
- Ventimiglia, G., Gerhard, O., Arnaboldi, M., & Coccato, L. 2010, A&A, 520, L9
- Weijmans, A.-M., Cappellari. M., Bacon, R., et al. 2009, MNRAS, 398, 561
- White, S. D. M. 1976, MNRAS, 174, 19
- White, S. D. M., Briel, U. G., & Henry, J. P. 1993, MNRAS, 261, L8
- Williams, B. F., Ciardullo, R., Durrell, P. R., et al. 2007a, ApJ, 654, 835
- Williams, B. F., Ciardullo, R., Durrell, P. R., et al. 2007b, ApJ, 656, 756
- Worthey, G., Faber, S. M., & Gonzalez, J. J. 1992, ApJ, 398, 69
- Worthey, G., Faber, S. M., González, J. J., & Burstein, D. 1994, ApJS, 94, 687
- Wu, X., Gerhrd, O., Naab, T., et al. 2014, MNRAS, 438, 2701
- Wu, X., & Tremaine, S. 2006, ApJ, 643, 210
- Zabludoff, A. I., Geller, M. J., Huchra, J. P., & Vogeley, M. S. 1993, AJ, 106, 1273
- Zabludoff, A. I., Huchra, J. P., & Geller, M. J. 1990, ApJS, 74, 1