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The pure non-collisional Blue Straggler population in the giant stellar system omega Centauri

We have used high spatial resolution data from the Hubble Space Telescope and wide-field ground-based observations to search for blue straggler stars (BSS) over the entire radial extent of the large stellar system omega Centauri. We have detected the largest population of BSS ever observed in any stellar system. Even though the sample is restricted to the brightest portion of the BSS sequence, more than 300 candidates have been identified. BSS are thought to be produced by the evolution of binary systems (either formed by stellar collisions or mass exchange in binary stars). Since systems like Galactic globular clusters (GGC) and omega Cen evolve dynamically on time-scales significantly shorter than their ages, binaries should have settled toward the center, showing a more concentrated radial distribution than the ordinary, less massive single stars. Indeed, in all GGCs which have been surveyed for BSS, the BSS distribution is peaked at the center. Conversely, in omega Cen we find that the BSS share the same radial distribution as the adopted reference populations. This is the cleanest evidence ever found that such a stellar system is not fully relaxed even in the central region. We further argue that the absence of central concentration in the BSS distribution rules out a collisional origin. Thus, the omega Cen BSS are the purest and largest population of non-collisional BSS ever observed. Our results allow the first empirical quantitative estimate of the production rate of BSS via this channel. BSS in omega Cen may represent the best local template for modeling the BSS populations in distant galaxies where they cannot be individually observed.

largest population of non-collisional BSS ever observed. Our results allow the first empirical quantitative estimate of the production rate of BSS via this channel. BSS in ω Centauri may represent the best local template for modeling the BSS populations in distant galaxies where they cannot be individually observed.

Introduction
Blue straggler stars (BSS) define a sparsely populated sequence extending to higher luminosity than the turnoff (TO) point of normal hydrogen-burning main sequence (MS) stars in the color-magnitude diagrams (CMD) of stellar aggregates like Galactic globular clusters (GGC). Superficially they appear to be MS stars with masses larger than expected for the cluster age as determined from the MSTO. There are two mechanisms thought to produce BSS: the first is mass exchange in a binary system, the second is merger of two stars induced by stellar interactions, i.e., collisions (either between single or binary stars) in a dense stellar environment. In either scenario, BSS are significantly more massive than normal cluster stars. Because of the high stellar density, collisions are more frequent in cluster centers. In addition, clusters evolve dynamically, and more massive objects sink toward the center on a time scale known as the relaxation time. Both factors suggest that BSS should be more concentrated in the central regions with respect to the other cluster stars, and this has been found to be the case in all GGCs with adequate observations of the center.
Among the stellar systems that populate the Galactic Halo, ω Centauri is beyond any doubt the most surprising. The entire body of evidence collected so far -kinematics, spatial distribution and chemical composition peculiarities -make ω Centauri a unique object if classed as usual among the globular clusters of the Milky Way (GGCs). It is more massive (M ≈ 2.9 × 10 6 M ⊙ , ) and luminous than any other GGC. It differs dynamically from ordinary GGCs: it is one of the most flattened cluster (Meylan 1987;White & Shawl 1987) and is partially supported by rotation (Meylan 1987;Merritt et -3al. 1997). The most astonishing peculiarity of ω Centauri is its metallicity spread measured both spectroscopically (Norris et al. 1996;Suntzeff & Kraft 1996) and photometrically (Lee et al. 1999;Pancino et al. 2000;Hilker & Richtler 2000;Frinchaboy et al. 2002;Sollima et al. 2005). Being the only halo stellar system with such a large chemical inhomogeneity, it has been suspected not to be a "genuine" globular cluster but the remnant of a dwarf galaxy that merged with the Milky Way in the past. Because of its proximity, this giant system is a cornerstone in our understanding of the formation, chemical enrichment and dynamical evolution of stellar systems.

The central BSS population
The results presented here are based on a high-resolution sample obtained with the Advanced Camera for Surveys (ACS) onboard the Hubble Space Telescope (HST) through B and R filters. The observations are organized in a mosaic of 3 × 3 pointings covering the 9 ′ × 9 ′ around the cluster center. Preliminary results based on this data-base (dealing with the complex structure of the SGB-TO region), have been already published (Ferraro et al. 2004a). Here we focus our attention on the BSS population. Figure 1 shows the zoomed color-magnitude diagram (CMD) in the BSS region. As can be seen, there is a clear, well defined and populated BSS sequence, cleanly separated from the MSTO stars. As reference, the isochrones for 2 and 3 Gyr old populations (Cariulo et al. 2004) are also shown. We adopted a global metallicity [M/H] ≃ −1.5 (or Z = 0.0006) (Ferraro et al. 2004a) for the metal-poor, dominant population, a distance modulus (m − M) 0 = 13.70 (i.e. ≈ 5500 pc, Bellazzini et al. 2004) and reddening E(B − V ) = 0.11 (Lub 2002). The mean interstellar extinction coefficients listed in Table 2 by Savage & Mathis (1979) have been adopted. As can be seen, the observed BSS sequence is nicely marked by the theoretical sequences. In particular, the positions of the brightest BSS are similar to MSTO stars in 2-3 Gyr old clusters, while the MSTO of the dominant metal-poor population, has been fitted by 14-15 Gyr isochrone (Ferraro et al. 2004a). The MSTO masses in 2 and 3 Gyr old clusters are 1.4 and 1.2 M ⊙ , respectively. This is in agreement with the recent finding of Rey et al. (2004) who also concluded that the masses of BSS detected in ω Centauri do not exceed 1.4 M ⊙ . Indeed, the MSTO mass in ω Centauri itself is estimated to be 0.74M ⊙ , so mass transfer in an equal mass binary should not produce a BSS more massive than 1.5M ⊙ .
The detection of such a clean BSS sequence allows us to perform a direct comparison with normal cluster stars. In order to be conservative and to avoid any possible contamination from spurious blended objects (mainly due to MS stars), we selected only the brighter (16 < B < 18.4) and bluer (B − R < 0.7) portion of the BSS sequence.
Generally, in order to study the BSS radial distribution, both horizontal branch (HB) and red giant branch (RGB) stars are used as reference stellar populations. In most clusters the HB is clearly defined, and it has been used in previous papers  for cluster to cluster comparisons. On the other hand, the RGB in the same magnitude range of the BSS, including the lower RGB, is much more populous than the HB, hence star counts are less affected by statistical fluctuations. For this reason we used both populations as reference. Particular care has been devoted in selecting the RGB sample: (i) AGB stars are clearly separable from the RGB and are excluded from the reference sample; (ii) we accurately selected RGB stars belonging to the dominant metal-poor branch. The possible contamination by the more metal rich RGB stars is always negligible. Figure 2 shows the cumulative radial distributions of BSS and reference populations. A Kolomogorov-Smirnoff (KS) test indicates that the radial distribution of the BSS has a ∼ 67% probability to be extracted from the same parent distribution as HB and RGB stars. This is the very first time that BSS have been found to share the same radial distribution of normal stars of the parent cluster. Eventually, as can be seen from Figure 2, BSS in ω Centauri appear, if anything, even less concentrated than HB and RGB stars, contrary to what observed in any other cluster.
To further investigate the radial distribution of BSS, we computed the BSS relative frequency F BSS HB = N BSS /N HB (where N BSS and N HB are the number of BSS and HB stars, respectively) and studied its behavior as a function of the distance from the cluster center. In doing this, we divided the sampled area into a set of concentric annuli, and in each annulus we have counted the number of BSS and HB stars. The relative frequencies as a function of distance from the cluster center are shown in Figure 3 (upper panels): the distribution appears nicely constant over the entire extent of the ACS sample 2 . Moreover, the constancy of the BSS frequency at different distances from the cluster center is independent of the choice of the reference population. We emphasize that this is the first time that a such behavior has been found in a globular cluster. In all previously surveyed clusters the central regions show a significant overabundance of BSS. Typically, the specific frequency drops by a factor of 4 or more over a few core radii 3 . Note that BSS have been found to be significantly 2 Errors in the relative frequency (F BSS HB ) of BSS with respect to the HB, are given by where σ H B and σ B SS are errors derived from the Poisson statistics.
more centrally concentrated even in clusters with significantly lower central density than ω Centauri (see for example NGC 288, Bellazzini et al. 2002).

Sampling the entire cluster
In order to further investigate the peculiar distribution of BSS in this stellar system, we extended the analysis to the external regions. To do this we used the wide-field photometry previously obtained by our group using the Wide Field Imager (WFI) on the 2.2m ESO-MPI telescope at the European Southern Observatory (ESO) on La Silla. The WFI is a mosaic of 8 CCD chips (each with a field of view of 8 ′ × 16 ′ ) giving a global field of view of 33 ′ × 34 ′ . The images were obtained using B and I filters and have been previously used (Pancino et al. 2000) to identify an additional, previously unknown RGB (RGB-a, see also Lee et al. 1999). Since the cluster center is slightly off-center in that data-set, we also supplement the sampling of the external region with the B, V catalog by Rey et al. (2004) 4 . The final catalog samples the cluster population over (complete) concentric annuli up to a distance of 20 arcmin from the cluster center.
All the catalogs were transformed to the same absolute astrometric system by using the Space Telescope Guide Star Catalog II, and then photometrically matched by using the stars in common. Note that all three datasets have the B passband in common. To be conservative and in order to prevent any possible incompleteness due to poor spatial resolution, we excluded the innermost region (r < 8 ′ ) of the WFI sample and used the Rey et al. catalog only in the most external region not sampled by our WFI observations.

The star density profile
As a first step, we used the selected sample of RGB/SGB in the magnitude range 16 < B < 18.4 to obtain a new density profile for the cluster. The analysis has been performed using almost 23,000 stars. The surveyed region has been divided in 27 concentric annuli. 5 Each annulus has been split in a number of subsectors (generally octants or quadrants). Then 4 The total region sampled by Rey et al. (2004) covers 40 ′ × 40 ′ . 5 The use of circular annuli is a good approximation. Even though ω Centauri is the most elliptical GGC, Pancino et al (2003) showed that the axial ratio of the entire RGB distribution changes only a few percent in the central region (r < 5 ′ ). In the outer region statistical fluctuations overwhelm any errors caused by the ellipticity. the number of stars lying within each sub-sector has been counted and averaged. Star density has been obtained by dividing the average star number by the corresponding sub-sector area (in arcsec 2 ). The resulting profile can be nicely reproduced by a King model characterized by two parameters: the core radius, r c and the concentration, c. We redetermined these parameters using our observations and the projected star density from a standard isotropic single-mass King model (Sigurdsson & Phinney 1995). The result is shown in Figure 4: we found r c = 153 ′′ = 2. ′ 57 and c = 1.31 in good agreement with the values (c = 1.24 and r c = 155 ′′ ) listed by Trager et al. (1995).

The BSS population over the entire cluster
We selected the BSS following the criteria adopted for the ACS sample. Only BSS in the magnitude range 16 < B < 18.4 were considered. We used isochrones (Cariulo et al. 2004) to convert the red edge of the selection box from the (B − R) color into (B − I) and (B − V ) colors. We found that the adopted red edge in the ACS sample (B − R = 0.7) corresponds to (B − I) ≈ 1.0 and (B − V ) ≈ 0.42. Figure 5 shows the zoomed CMDs in the BSS region for the ACS and WFI samples. The selection boxes are shown in each panel. The BSS population appears as a clear sequence diagonally crossing each box. To be conservative, we selected only stars within the two dashed lines. This accounts for the bulk of the BSS populations with only a few objects being excluded from the selection. Following these criteria 6 we have identified 158 BSS in the HST-ACS and 155 in the external sample (r > 8 ′ , with 117 found in the region sampled by the WFI and 38 in the Rey et al. sample), for a total of 313 BSS. This is the cleanest and largest BSS sample detected in the cluster. However, the global BSS population in ω Centauri probably is significantly larger (> 400) since (i) we sampled ≈ 70% of the cluster light; (ii) we limit our selection to the brightest and cleanest portion of the BSS sequence.
Of course each photometric sample has its own incompleteness, and for this reason we analyzed them separately. According to section 2, each BSS population is referred to the RGB population of the corresponding catalog in the same magnitude range (16 < B < 18.4) of the selected BSS.
To compare the samples we used the doubly normalized ratio R BSS which gives the fraction of BSS counted in concentric annuli at different distance from the cluster center with respect to the fraction of light sampled in each annulus. The sample presented here covers the cluster extent up to 20 ′ from the center. Hence, the entire surveyed area has been divided into six concentric annuli excluding only the region between 320 ′′ and 480 ′′ . The cluster light sampled in each annulus has been computed by integrating the King profile (King 1966) fitted to the observed cluster density profile shown in Figure 4. The same ratio has been computed for the reference population. As emphasized in Ferraro et al. (1993), for a stellar population which is distributed accordingly with the integrated cluster light, this ratio is 1. The result is shown in Figure 6, and as can be seen the normalized BSS ratio is nicely constant and fully consistent with the reference population.
In other clusters (M3: Ferraro et al. 1997, 47 Tuc: Ferraro et al. 2004b, NGC 6752: Sabbi et al 2004, M55: Zagggia et al. 1997) not only does the normalized BSS population peak at the center, but it again rises in the cluster periphery. This effect probably arises because primordial binaries whose orbits are confined to the cluster exterior have low collision rate and thus remain in the exterior where some eventually become BSS. BSS progenitors at intermediate radii drift inward where they are either "ionizated" or driven to merge by collisions. Dynamical simulations have demonstrated this phenomenon in 47 Tuc (Mappeli et al. 2004).
Since this result seems to be quite peculiar, a natural question arises: Is there any other evidence that can support our findings? First we note that there is no convincing evidence of equipartion effects in the radial distribution of MS stars of ω Centauri (Anderson 2002). What about the radial distribution of other sub-populations significantly more massive than normal cluster stars? Interacting binaries containing a compact object (like a neutron star or white dwarf) in which mass transfer is occurring are expected to show X-ray emission. These are among the most massive objects we can currently find in an old stellar population. We used the recent XMM-Newton observations (Gendre et al. 2003) in order to check it. There are 42 X-ray sources lying in the field-of-view covered by the ACS observations. The radial distribution of these sources is compared with that of BSS in the right panel of Figure 2. As can be seen the radial distribution of the X-ray sources is the same or even less centrally concentrated than the BSS. 7 A KS test shows that BSS and X-ray sources have a ∼ 12% probability of being extracted from the same parent distribution. However due to the small sample of detected X-ray sources, we can conservatively conclude that the difference between the two distributions is not significant. In conclusion neither of the two most massive star populations in the cluster (BSS and X-ray sources) shows any signature of radial segregation with respect to the normal stars in the cluster.
Since relaxation is the major dynamical process which differentiates stars according to their mass (Meylan & Heggie 1997), the observational facts presented here represent the cleanest evidence found so far that the system is still far from being completely relaxed even in the core region.

Discussion
The evidence presented here demonstrates that ω Centauri, the largest stellar system of the galactic halo, does not share the same dynamical characteristics of GGCs. Indeed because of its mass, previous hints and arguments have led, in the past, to the suspicion that ω Centauri is not completely relaxed . Since ω Centauri is one order of magnitude more massive than a typical halo GGC, its relaxation time-scale is also expected to be significantly larger. Previous estimates (Djorgovski 1993) of the central relaxation time for ω Centauri, give log t crt = 9.76 i.e., t crt ≈ 5.7 Gyr, adopting a distance of d ≈ 4900 pc. This distance is shorter than that currently adopted here (d ≈ 5500 pc). For this reason we re-computed the relaxation time for this stellar system, adopting the new distance and the structural parameters obtained in Section 3.1. Under these assumptions, the physical size of the core radius turns out to be r c ∼ 4.1 pc. By adopting the observed integrated magnitude V t = 3.68 (Harris 1996), the resulting central relaxation time is log t crt = 9.82, i.e t crt ∼ 6.6 Gyr. Although larger, this new determination of the relaxation time is still a factor of two shorter than the cluster age (12-14 Gyr), hence some segregation should be visible at least in the central regions. A good indication of the expected segregation time-scale for an object of mass m * can be obtained by considering the half-mass relaxation time-scale (Davies et al. 2004). Using Equation (10) of Davies et al. (2004) we found that roughly half of the m * = 1.0 M ⊙ BSS should have been sunk to the core after only 2 Gyr.

Is the lack of segregation connected to the origin of this peculiar stellar system? Can the history of a stellar system influence its internal dynamical status? Is it due to rotation?
According to Davies et al (2004) heavier stars tend to take longer to sink in the core of more massive systems. If we accept the hypothesis that ω Centauri is the relic of a dwarf galaxy partially disrupted by the tidal field of the Milky Way, then this system was significantly more massive in the past than what we observe today. According to Tsuchiya et al. (2004) the initial mass of ω Centauri could have been M ≈ 10 8 M ⊙ and the cluster should have remained more massive than 10 7 M ⊙ for a few Gyrs. Equation (10) of Davies et al (2004) suggests that the sinking-time is a factor 7-8 longer for a system 10 2 times larger than the current ω Centauri. In such a scenario, the stormy past of this stellar system could have extended the sinking-time needed for the heaviest population to reach the cluster center to a time larger than the cluster age.
Rotation can also play a role. Relaxation time is expected to be longer for rotating systems, since angular momentum tends to keep stars out of the core, counteracting mass segregation (Spurzem 2001). There is evidence for large rotation of the ω Centauri system (Meylan & Mayor 1986;van Leeuwen & Le Poole 2002). Moreover, while the metal poor giants ([Ca/H] < −1.2), belonging to the dominant cluster population form a rotating system, the metal rich ones do not show any significant rotation (Norris et al. 1997). Such evidence suggests a clear correlation between dynamical and chemical properties of ω Centauri. Perhaps, the lack of equipartition observed could be related to the origin of the multiple populations inside the cluster.
Although, we do not fully understand why ω Centauri is not relaxed yet, we still feel safe in concluding that stellar collisions have played a minor role in generating exotic binary systems in ω Centauri. We have previously made detailed studies of nine clusters. Our major conclusion is that the overall population of BSS in a typical GGC turns out to be a complex conglomeration of collisional BSS and mass exchanging binaries. Still, while several puzzles remain, some firm facts are emerging. Among these are that, except for possibly NGC 288, all clusters have some collisional BSS, and those BSS are strongly centrally concentrated. This is hardly surprising since the collision rate is higher at the center, and since the BSS are more massive than the typical cluster star they will not migrate outwards. It is possible that a "kick" produced by the collision ejects them from the core (Sigurdsson et al. 1994, but see the discussion in Mapelli et al. 2004). However, Figure 6 shows that the number of BSS nicely scales with the sampled luminosity in the central region as well as in the outer regions. It would be most remarkable if some combination of collision rate and kicks managed to produce the observed flat BSS distribution. Applying Occam's razor, the simplest explanation is that the BSS observed today in ω Centauri are the progeny of primordial binaries whose radial distribution has not been yet significantly altered by collisions.

Recently Piotto et al (2004) and Davies et al (2004) from a large survey of BSS in 56
GGCs showed that the total number of BSS is largely independent of the cluster luminosity and collision rate, suggesting that no single process produces BSS in GGCs (accordingly with Baylin 1995, Fusi Pecci et al 1992, hence both binary evolution and collisions both could be active processes in producing BSS in different environments. In particular, Davies et al. (2004) developed a model for the production of BSS in GGCs. In the low mass systems (M V > −8) BSS arise mostly from mass exchange in primordial binaries. In more massive systems collisions produce mergers of the primordial binaries early in the cluster history. BSS resulting from these mergers long ago evolved away. Once the primordial binaries were used up, BSS produced via this channel disappeared. In the cores of the most massive systems (M V < −9) collisional BSS are produced (Fig. 6 of Davies et al.).
Detailed cluster-to-cluster comparison has shown that the scenario is much more complex than that proposed by Davies et al (2004). The dynamical history of each cluster apparently plays a role in determining the origin and radial distribution BSS content. For example, (i) clusters with comparable luminosity show vastly different BSS population, see for example the case of M3/M13 (M V ∼ −8.7 and −8.5 respectively, see Ferraro et al, 2003) and NGC6752/M80 (M V ∼ −7.7 and −7.9 respectively, see Sabbi et al 2004 andFerraro et al 1999) pairs; (ii) studies of the BSS population over the entire cluster extension (M3, Ferraro et al 1997;47 Tuc, Ferraro et al 2004b;NGC 6752, Sabbi et al 2004) have further supported the complexity of the emerging scenario showing that the central population of BSS is only a component of the entire BSS population of each cluster. In addition to a strongly centrally segregated population, these GGCs also have a population of external BSS which we argue arise from primordial binaries. At intermediate radii there is a paucity of BSS of either type, because the collision rate is low and primordial binaries settled to the center and probably became BSS long ago.
How does ω Centauri fit into the Davies et al. picture? Following that scenario, since ω Centauri is very luminous, its primordial binaries would have been destroyed long ago, while a population of several hundred collisional BSS should populate the core of the cluster (see Figure 6 by Davies et al 2004). Hence we would expect to see a population of collisional BSS in the core and possibly a population of surving primordial binaries in the outer cluster. This would possibly produce a bimodal distribution as observed in other clusters, or even a single centrally-peaked distibution but in any case it cannot produce the flat distribution shown in Fig For sake of comparison, we can compare the expected collision rate for ω Centauri (under the assumption that rotation or the cluster history has not made the estimate invalid) to that expected in other massive GGCs such as 47 Tuc. According to Davies et al (2004) the collision rate is Γ coll = 5 × 10 −15 × (r c Σ 3 0 ) 1/2 in units of collisions per year (where Σ 0 is the central surface brightness in units of L ⊙V pc −2 (equivalent to µ V = 26.41) and r c is the core radius. By assuming r c = 0.42 pc (Mapelli et al 2004) and r c = 4.1 pc for 47 Tuc and ω Centauri , respectively, we find that the expected collision rate is only 5-8 times larger in 47 Tuc as compared to ω Centauri (approssimately 50-60 collisions per Gyr in the core of 47 Tuc with respect to 6-11 in the core of ω Centauri ) depending on the adopted central surface brightness values (Harris 1996or Djorgovski 1993. In any case this difference does not seem sufficient to explain the quite different BSS content in the two clusters 9 . This evidence further supports the complexity of the BSS formation and evolution scenario and suggests that the dynamical state of the cluster and its history plays an important role in determining the BSS content. Using the well populated BSS sequence in ω Centauri we have the possibility to determine the production rate of BSS in a non-collisional system. From the upper panel of Figure 3 it is easy to derive that the primordial-binary BSS specific frequency F BSS HB is 0.08. Note that values up to one order of magnitude larger (F BSS HB ≈ 1) have been detected in the central region of GGCs ). An even more useful quantity for the study of an unresolved stellar population is the ratio of the number of BSS to the sampled luminosity in units of 10 4 L ⊙ , S4 BSS = N BSS /L 4 (Ferraro et al. 1995). S4 BSS varies from cluster to cluster (ranging from 0.2 to 30) possibly tracing the generation/destruction effect of stellar interactions in collisional systems like the dense cores of GGCs. The bottom panel of Figure 3 shows the radial behavior of S4 BSS within ω Centauri. The rate of production is constant throughout the system despite the large variation of stellar density. If S4 prim BSS refers to primordial binary BSS, our results suggest that S4 prim BSS ∼ 2. This means that in absence of collisions ∼ 2 BSS can be found for each ∼ 10 4 L ⊙ of sampled light. This value can be used to estimate the expected number of BSS generated by primordial binaries for each fraction of sampled light in any stellar system.
Of course the specific frequency of primordial BSS is expected to depend on the fraction of primordial binaries contained in the considered stellar system. This fraction is still largely unknown, it has been estimated only in a few traditional GGCs (NGC 6752- Rubenstein & Bailyn 1997, NGC 288-Bellazzini et al. 2002, 47 Tuc-Albrow et al. 2001) and some lower density systems (Pal 13- Clark et al. 2004, E3-Veronesi et al. 1996. In one of those (47 Tuc) the determination has been obtained in the very dense central region, where binaries are created, destroyed and hardened by stellar collisions. It would be highly desirable to directly measure the binary fraction in ω Centauri. Sadly the technique used by Bellazzini et al. (2002) and Rubenstein & Bailyn (1997) it can hardly be applied to ω Centauri because its MS is broadened by the abundance spread.
How does the BSS specific frequency in ω Centauri compare with the field? The BSS specific frequency (F ∼ 0.1) measured here turns out to be 40 times lower than that (F ∼ 4) found by Preston & Sneden (2000) in the field. Indeed, such a large frequency has never been observed in any GGC. One possible origin of this discrepancy is that GGCs never harbored the sort of primordial binaries that produce the long period BSS binaries found in the field by Preston & Sneden. However, it must be noted that their value is only an indirect estimated derived from a chain of assumptions based on the observed fraction (60%) of the blue metal poor field stars.
Indeed the low BSS frequency compared to the field and the lack of segregation in ω Centauri are hard to reconcile. Strong cluster rotation, for example, could depress the frequency of collisional BSS but would probably not affect the primordial binary population of the cluster. Then, if the data in the field are eventually confirmed, the observations presented here strongly point at a substantial difference between the binary population in the field and in the cluster.
As final consideration, we note that there is now a growing interest in determining the age of distant stellar systems from their integrated spectra. The main age discriminants arise from the stars at the MSTO. Since BSS are both hotter and more luminous than MSTO stars, a significant population of BSS can make a system appearing younger. Because of the stellar density in galaxies is low enough that collisional BSS are probably not important, BSS in galaxies should have properties similar to those detected in ω Centauri. Hence, this BSS population can represent a suitable template for estimating the contribution of BSS to the global spectral energy distribution of distant galaxies.