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3. The colour-magnitude diagram of NGC 6528

3.1. General remarks, structural parameters, field

Figure 2 (click here) shows a CMD of all 24999 stars that were identified independently on at least 4 frames. The V-I scale is so compressed by the occurence of very red stars that features like the horizontal branch are not easily recognizable. It is clear, however, that the dominant population is not the cluster but the field population of the Galactic bulge. To maintain good legibility by avoiding an inconvenient compression of the V-I scale, Fig. 5 (click here) plots only stars bluer than V-I =3.0. No selection with respect to photometric errors was done to maintain completeness. The full list of stars (Table 6 (click here)) is published in electronic form at CDS.

Figure 2:   The colour-magnitude diagram shows all 24999 stars, for which photometry could be obtained. The V-I scale is strongly compressed by the occurence of very red stars, so that the CMD features of the cluster itself are hard to isolate. The cluster is embedded in a very rich population of field stars, which is the main obstacle for a proper derivation of cluster parameters. Triangles are those stars redder than V-I = 3.0, which are closer than tex2html_wrap_inline2358 to the cluster center and accordingly have a high probability of being cluster members. They show that the cluster is situated in the foreground with respect to the red field star population

This CMD shows why the investigation of this cluster is so difficult. Practically all interesting structures in the CMD of NGC 6528 are contaminated by the field population. While the horizontal branch (HB) can be distinguished, the exact shape of the red giant branch (RGB) cannot be clearly identified. OBB found strong differential reddening in the area of NGC 6528 which further broadens the RGB and other features of the CMD. Also a "tilted'' HB is recognisable, although less striking than in other clusters like NGC 6553 (Subramaniam et al. 1997) or Terzan 5 (Grebel et al. 1995).
The turn-off region is reached in our photometry. It is at most 3.5 mag fainter than the HB, which is typical of old clusters (e.g., Buonanno et al. 1989), but it is completely hidden in the rich field. Here we skip any discussion of the TO location due to the superior data of HST (Ortolani et al. 1995).

The problem now consists of properly selecting stars with a high probability of being cluster members and statistically subtracting field stars in order to isolate the significant features of the CMD. A useful approach is to determine the radial distance beyond which the field population starts to dominate over the cluster population.

A King model reads (King 1962)

where K is the central stellar surface density (or surface brightness, respectively), tex2html_wrap_inline2362 the core radius and tex2html_wrap_inline2364 the tidal radius. Webbink (1985) quotes tex2html_wrap_inline2366 arcmin and tex2html_wrap_inline2368 arcmin, while Trager et al. (1993) give tex2html_wrap_inline2370 arcmin and tex2html_wrap_inline2364 = 16.6 arcmin. These differences illustrate the difficulties of deriving structural parameters for a cluster so strongly embedded in the field population. Our attempts to fit a King profile confirmed the degeneracy of the parameters. However, our area is quite limited so we cannot improve the parameter values.

We determined the radial profile of the surface density, selecting stars brighter than V = 19 mag, where we are complete except perhaps for the very center. We find for the central surface density tex2html_wrap_inline2376 stars/square arcmin. The observed density profile together with two King models with the above parameters are given in Fig. 3 (click here). From this plot, the cluster and field densities are equal at a radius 1.1 arcmin. At 2.5 arcmin, the field density is already 8 times higher than the cluster density in case of the Trager et al. (1993) model.

The CMD of the field population (Fig. 4 (click here)) therefore consists of stars with radial distances larger than 2.5 arcmin.

Figure 3:   This plot shows the observed stellar surface density (logarithm of stars brighter than V=19 mag per square arcmin) as a function of radius (triangles). Also plotted are two King models with parameters quoted by Webbink (1985, solid line) and Trager et al. (1993, dotted line). At a radial distance of about 1.1 arcmin, the field population equals the cluster population. For constructing the field CMD, one should select stars with radial distances larger than 2.5 arcmin, where the cluster population is negligible

Figure 4:   The CMD of the field population around NGC 6528. We selected stars with radial distances larger than 2.5 arcmin to plot this diagram. Triangles are those stars redder than V-I = 3.0, which are closer than tex2html_wrap_inline2358 to the cluster center and accordingly have a high probability of being cluster members. One recognizes that the giant branch of NGC 6528 constitutes a brighter envelope to this star distribution and thus is located in the foreground

Figure 5:   This plot shows the CMD in V, V-I of the NGC 6528 region containing only stars bluer than V-I = 3.5. It is obvious that this diagram is overwhelmingly dominated by the population of the Galactic bulge. The HB of the cluster is discernible while the RGB cannot be clearly identified. The turn-off region is reached in our photometry but the turn-off itself is completely hidden in the field population

3.2. The red giant branch (RGB) and the horizontal branch (HB)


The location of the RGB is the basis for estimation of the metallicity and reddening. In Fig. 5 (click here), it is difficult to locate the cluster RGB among the numerous field stars. We found that statistical field star subtraction in the region of the RGB does not give useful results due to small number statistics. Selection according to distance from the cluster worked much better.

We found that selecting the stars with distances less than 100 pix (0.553 arcmin) gave the best representation of the RGB. We are using this selection to define a fiducial for the RGB and the HB (both given in Fig. 6 (click here)). Numerical values of the adopted fiducial are presented in Table 7 (click here).

It is interesting to note that the "tilt'' of the HB recognizable in Fig. 5 (click here) vanishes in Fig. 6 (click here). Such tilts have been repeatedly reported for metal-rich clusters. Armandroff (1988) proposed differential reddening as the cause, because the slope of the tilted feature resembles in the most striking cases the reddening vector. OBB already remarked that the HB gets more clumpy when HB stars near the cluster center are considered, which is confirmed by Fig. 5 (click here). This means that the tilt is mainly due to differential reddening in the field population.
On the other hand, NGC 6553 (Ortolani et al. 1991; Subramaniam et al. 1997) and Terzan 5 (Grebel et al. 1995) are examples, where the tilted HB branch is seemingly not induced by a smooth reddening gradient in the cluster field, but rather by extremely patchy reddening. This conjecture is strengthened by the fact that the "tilted'' HB is more striking with increasing foreground reddening. However, a deeper discussion must include better data of highly reddened clusters.

Figure 6:   CMD resulting from a selection with of stars closer than tex2html_wrap_inline2358 to the cluster center. This selection was used to define a fiducial for the RGB and the HB, which is overplotted (filled circles). It can be seen that the apparent "tilt'' of the HB, which may be present in Fig. 2 (click here) vanishes if the stars on the HB have a high probability of being cluster members


2.909 15.489
2.645 15.672
2.309 16.120
2.151 16.486
2.003 16.818
1.887 17.317
1.835 17.616
1.787 17.998
1.719 18.663
1.640 19.410
1.645 19.826
1.592 20.125
1.661 17.217
1.698 17.217
1.719 17.217
1.740 17.217
1.761 17.217
1.745 17.234
1.724 17.234
1.703 17.234
1.692 17.234
Table 7:   Our fiducial for the RGB of NGC 6528 derived by radius selection

3.3. The "curved giant branch''

Figure 2 (click here) shows many very red stars, the reddest ones have V-I = 7. They cannot be late M dwarfs since they are preferably found among the brighter stars, so they should be intrinsically bright. OBB interpret their occurence as a continuation of the RGB, caused by the high metallicity of the cluster and suggest that these stars are affected by strong blanketing.

The cluster membership of these stars is naturally of interest. If giants with such red V-I colours occur in metal rich populations, we expect these stars to be present among to the Galactic bulge field stars as well. Indeed, the large majority of the stars redder than V-I = 3.5 does not show any concentration towards the cluster center and therefore likely belong to the bulge population.

However, a subsample have a high probability of being cluster members, since a selection according to radial distance (less than tex2html_wrap_inline2416 from the cluster center) is equivalent to the selection of the outer envelope of the distribution of the red stars (see Fig. 4 (click here)). This implies that the cluster is not embedded in the field population, but is located in the foreground.

Garnavich et al. (1994) found a "curved giant branch'' in the V-I, V-diagram of NGC 6791, a metal-rich old open cluster. Their calculations show that the decreasing V brightness can be understood simply as a consequence of a rapidly increasing bolometric correction. Quantitative theoretical statements concerning colours are difficult to make, since no detailed atmospheric models exist for very cool giants (c.f. Gustafsson & Jørgensen 1994), but models for cool dwarf stars show that already in the black or grey body approximation the stellar continuum at or below temperatures of 3000 K easily produces colours as red as tex2html_wrap_inline2424.

In Fig. 7 (click here) a theoretical isochrone from Bertelli et al. (1994) is overplotted. The isochrone has Z=0.02 and an age of 13.2 Gyr. The solid line is the RGB while the filled circles trace the AGB. It turns out that, if the reddest stars are indeed cluster members, the discrepancy between the theoretical isochrone and the actually observed locus starts at tex2html_wrap_inline2428. The evolutionary status of the red stars is not clear, but one may suspect in analogy to the theoretical isochrones that they are on the AGB rather than on the RGB.
To assess at least roughly their effect on integrated colours, we first calculate an integral V-I colour of all stars in Fig. 2 (click here) and get 1.82 mag, while omitting the stars with V-I > 3.5 (as in Fig. 5 (click here)) results in an integrated colour of 1.70 mag. This effect is plausibly even stronger in the infrared. So this red extension of the AGB/RGB stars must be taken into account, when modelling spectra of elliptical galaxies. See the contributions by Bruzual (1996), Worthey (1996), Chiosi (1996) for a deeper discussion of related problems.

Figure 7:   This plot shows the full colour range of stars with the selection of being closer than tex2html_wrap_inline2358 to the cluster center. Overplotted is a theoretical isochrone from Bertelli et al. (1994) with an age of 13.2 Gy and solar metallicity, including the Red Giant Branch (RGB, solid line) and the Asymptotic Giant Branch (AGB, filled circles). While the isochrone traces the RGB and the AGB well to tex2html_wrap_inline2428, they deviate at redder colours and also do not cover the full extension of the CMD. Whether the reddest stars are RGB or AGB stars, is unclear

3.4. The main sequence turn-off of NGC 6528


The question for the age of NGC 6528 is of high importance for the metal-rich bulge population of which NGC 6528 is a representative. Is NGC 6528 placed among those few globular clusters that are younger than the bulk of clusters, like Pal 12 (Gratton & Ortolani 1988) or Ruprecht 106 (Buonanno et al. 1990)? That "disk'' clusters may be younger than halo clusters has been suggested for instance by Richtler et al. (1992) and Hatzidimitriou (1991). Ortolani et al. (1991) cited NGC 6553 as possibly younger. However, Richtler et al. (1994) showed that three other clusters, which had been classified as disk clusters (but in fact have properties more consistent with halo clusters), could not be distinguished in age from the bulk of globular clusters.
After the first version of this paper had been submitted, Ortolani et al. (1996) published a CMD of NGC 6528, based on HST-data (WFPC2). Although the present data are the deepest so far obtained from the ground for NGC 6528, they are definitely less reliable concerning the turn-off location due to the strong crowding at the relevant magnitude level. Therefore we renounce a deeper discussion and simply refer to Ortolani et al. who found that NGC 6528 is indistinguishable in age from the metal-poorer clusters in the galactic halo.

3.5. Reddening & Metallicity

Sarajedini (1994) devised a simple method to derive metallicity and reddening simultaneously from a V, V-I diagram by using two metallicity indicators, tex2html_wrap_inline2448, the magnitude difference between HB and RGB at the unreddened V-I colour 1.2, and tex2html_wrap_inline2452, the (unreddened) V-I colour of the RGB at the HB magnitude level. These metallicity indicators depend in different ways on the reddening. Thus we can regard reddening and metallicity to be determined, if both indicators give the same metallicity for the same reddening. However, a direct application is not possible since in Sarajedini's calibration, 47 Tuc appears with -0.7 dex as the most metal-rich cluster, whereas one wants to have a calibration extended to at least solar metallicity.

At present this is practically impossible to do in a purely empirical way. There are simply no V, V-I CMDs for clusters suitable for inclusion in the metal-rich domain as a supplement to Sarajedini's calibration. But there are also only a few theoretical isochrones available for clusters this metal-rich. The only published V-I isochrones for old and metal-rich clusters that we are aware of are those of Tripicco et al. (1995) and from Bertelli et al. (1994). Our approach is therefore to use these isochrones like empirically determined CMD loci and supplement the calibration set of Sarajenini.

3.5.1. The isochrones of Tripicco et al. (1995)

The 12 Gyr, solar metallicity isochrone of Tripicco et al. (1995) agrees very well with a linear extrapolation of Sarajedini's relations and its inclusion has only a little effect on the coefficients. For this isochrone, tex2html_wrap_inline2468 and tex2html_wrap_inline2470. Measuring these indicators, we adopted for the brightness of the HB MV = 1.05 mag, as indicated by the red part of the theoretical HB of Tripicco et al. (1995). A linear regression then leads to

These relations can now be used individually to plot reddening against metallicity, using the proper values for NGC 6528, for which we adopted tex2html_wrap_inline2474 and the numbers given in Table 8 (click here). These two lines then intersect at a certain pair of values that we regard to be the cluster reddening and metallicity. It is apparent that the main uncertainty is introduced by tex2html_wrap_inline2476, which loses sensitivity in the metal-rich domain. We conclude from this method that the reddening of NGC 6528 is E(V-I) = 0.8 and a metallicity of tex2html_wrap_inline2480.


V-I tex2html_wrap_inline2172
2.62 0.06
2.72 0.39
2.80 0.64
2.95 1.00
1.16 1.16
Table 8:   This table gives the magnitude difference between the corresponding points on the RGB and the HB for different V-I colours. Note that V-I = 1.2 + E(V-I), so the table gives tex2html_wrap_inline2448 for different values of the reddening

3.5.2. The Bertelli et al. (1994) isochrones

If Sarajedini's relation is supplemented with the aid of the Bertelli et al. (1994) isochrones, the linearity is lost and a significant curvature is introduced. We chose the [13.2 Gy, Z=0.02] isochrone for extrapolating to solar metallicity.

For this isochrone, tex2html_wrap_inline2498 and tex2html_wrap_inline2500. Including this point in Sarajedini's relation, leads to


Applying the above relations, our result for reddening and metallicity is E(V-I) = 0.6 and [M/H]=-0.4 dex. The uncertainty expressed by this difference can presently not be avoided.

3.6. Distance


tex2html_wrap_inline2510 tex2html_wrap_inline2512 tex2html_wrap_inline2514 [M/H] (T) tex2html_wrap_inline2518 [M/H] (B)
tex2html_wrap_inline2522 tex2html_wrap_inline2524 0.8 +0.1 dex 0.6 -0.4 dex
Table 9:   Measured parameters for NGC 6528 using an extrapolation of Sarajedini's (1994) calibration with isochrones of Bertelli et al. (1994) (B) and Tripicco et al. (1995) (T)

Calculating the distance based on HB brightness, we have to adopt a relation between brightness and metallicity for HB stars (basically for RR Lyrae stars). A recent compilation of related work has been given by Chaboyer et al. (1996). Their preferred relation is

To span a representative range of values occurring in the literature, we also adopted the relation favoured by Nemec et al. (1994):

Moreover, the absorption depends on the colour of the star considered. Assuming tex2html_wrap_inline2530 (Grebel & Roberts 1995) as the appropriate relation for our value of tex2html_wrap_inline2532, and tex2html_wrap_inline2534, we get the values for MV(HB), absorption, distance modulus and corresponding distance listed in Table 10 (click here). This table shows the full dilemma of distance determination of reddened globular clusters. We note that the ration AV/E(B-V) might also be discussed. In case of NGC 6528, a change of 0.3 would cause a shift of 0.25 mag in the distance modulus. Formally, there is no hard reason to consider any of these values as being more reliable than the others. However, since the structure of the CMD sets the cluster in the foreground, the high values are not probable. The recent investigation by Fusi Pecci et al. (1996) of M 31 globular clusters argues in favour of the Chaboyer et al. relation, so the distance value of 6.6 kpc might be preferred. One notes that the external error is much more dependent on the uncertain metallicity, absorption and HB brightness than on the photometric errors.


E(V-I) tex2html_wrap_inline1594 MV(HB) AV m-M D[kpc] z[kpc] tex2html_wrap_inline2554
0.8 +0.1 1.22 (N) 2.08 13.91 6.0 0.4 2.0
0.8 +0.1 1.02 (C) 2.08 14.11 6.6 0.5 1.5
0.6 -0.4 1.06 (N) 1.55 14.60 8.3 0.6 0.7
0.6 -0.4 0.90 (C) 1.55 14.76 8.9 0.7 1.1
Table 10:   Adopted and derived parameters for NGC 6528 using different relations (see text) between HB brightness and metallicity. It is indicated whether the Chaboyer et al. (1996) (C) or the Nemec et al. (1994) (N) relation is used. It becomes clear that the external distance uncertainty is completely dominated by these relations and the uncertain absorption. For the calculation of the galactocentric distance, we assumed 8 kpc as the distance from the sun to the Galactic Center

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