The Blue Straggler Population of the Globular Cluster M5
The Blue Straggler Population of the Globular Cluster M5 1
B. Lanzoni1,2, E. Dalessandro1,2, F.R. Ferraro1, C. Mancini3, G. Beccari2,4,5, R.T. Rood6,
M. Mapelli7, S. Sigurdsson8
1 Dipartimento di Astronomia, Università degli Studi di Bologna, via Ranzani 1, I–40127
Bologna, Italy
2 INAF–Osservatorio Astronomico di Bologna, via Ranzani 1, I–40127 Bologna, Italy
3 Dipartimento di Astronomia e Scienza dello Spazio, Università degli Studi di Firenze,
Largo Enrico Fermi 2, I– 50125 Firenze, Italy
4 Dipartimento di Scienze della Comunicazione, Università degli Studi di Teramo, Italy
5 INAF–Osservatorio Astronomico di Collurania, Via Mentore Maggini, I–64100 Teramo,
Italy
6 Department of Astronomy and Astrophysics, The Pennsylvania State University, 525
Davey Lab, University Park, PA 16802
7 S.I.S.S.A., Via Beirut 2 - 4, I–34014 Trieste, Italy
8 Astronomy Department, University of Virginia, P.O. Box 400325, Charlottesville, VA,
22904
20 March, 07
ABSTRACT
By combining high-resolution HST and wide-field ground based observations,
in ultraviolet and optical bands, we study the Blue Stragglers Star (BSS) popula-
tion of the galactic globular cluster M5 (NGC 5904) from its very central regions
up to its periphery. The BSS distribution is highly peaked in the cluster center,
decreases at intermediate radii and rises again outward. Such a bimodal dis-
tribution is similar to those previously observed in other globular clusters (M3,
47 Tucanae, NGC 6752). As for these clusters, dynamical simulations suggest
that, while the majority of BSS in M5 could be originated by stellar collisions,
a significant fraction (20-40%) of BSS generated by mass transfer processes in
primordial binaries is required to reproduce the observed radial distribution. A
candidate BSS has been detected beyond the cluster tidal radius. If confirmed,
this could represent an interesting case of an ”evaporating” BSS.
http://arxiv.org/abs/0704.0139v1
– 2 –
Subject headings: Globular clusters: individual (M5); stars: evolution – binaries:
general - blue stragglers
1. INTRODUCTION
In globular cluster (GC) color-magnitude diagrams (CMD) blue straggler stars (BSS)
appear to be brighter and bluer than the Turn-Off (TO) stars and lie along an extension of
the Main Sequence. Since BSS mimic a rejuvenated stellar population with masses larger
than the normal cluster stars (this is also confirmed by direct mass measurements; e.g. Shara
et al. 1997), they are thought to be objects that have increased their initial mass during
their evolution by means of some process. Two main scenarios have been proposed for their
formation: the collisional scenario suggests that BSS are the end-products of stellar mergers
induced by collisions (COL-BSS), while in the mass-transfer scenario BSS form by the mass-
transfer activity between two companions in a binary system (MT-BSS), possibly up to the
complete coalescence of the two stars. Hence, understanding the origin of BSS in stellar
clusters provides valuable insight both on the binary evolution processes and on the effects
of dynamical interactions on the (otherwise normal) stellar evolution.
The relative efficiency of the two formation mechanisms is thought to depend on the en-
vironment (Fusi Pecci et al. 1992; Ferraro et al. 1999a; Bellazzini et al. 2002; Ferraro et al.
2003). COL-BSS are expected to be formed preferentially in high-density environments (i.e.,
the GC central regions), where stellar collisions are most probable, and MT-BSS should
mainly populate lower density environments (the cluster peripheries), where binary systems
can more easily evolve in isolation without suffering exchanges or ionization due to gravita-
tional encounters. The overall scenario is complicated by the fact that primordial binaries
can also sink to the core due to mass segregation processes, and “new” binaries can be formed
in the cluster centers by gravitational encounters. The two formation mechanisms are likely
to be at work simultaneously in every GC (see the case of M3 as an example; Ferraro et al.
1993, 1997), but the identification of the cluster properties that mainly affect their relative
efficiency is still an open issue.
One possibility for distinguishing between the two types of BSS is offered by high-
resolution spectroscopic studies. Anomalous chemical abundances are expected at the surface
1Based on observations with the NASA/ESA HST, obtained at the Space Telescope Science Institute,
which is operated by AURA, Inc., under NASA contract NAS5-26555. Also based on WFI observations
collected at the European Southern Observatory, La Silla, Chile, within the observing programs 62.L-0354
and 64.L-0439.
– 3 –
of BSS resulting from MT activity (Sarna & de Greve 1996), while they are not predicted
in case of a collisional formation (Lombardi, Rasio & Shapiro 1995). Such studies have just
become feasible, and the results found in the case of 47 Tucanae (47 Tuc; Ferraro et al.
2006a) are encouraging. The detection of unexpected properties of stars along standard
evolutionary sequences (e.g., variability, anomalous population fractions, or peculiar radial
distributions) can help estimating the fraction of binaries within a cluster (see, e.g., Bailyn
1994; Albrow et al. 2001; Bellazzini et al. 2002; Beccari et al. 2006), but such evidence does
not directly allow the determination of the relative efficiency of the two BSS formation
processes.
The most widely applicable tool to probe the origin of BSS is their radial distribution
within the clusters (see Ferraro 2006, for a review). This has been observed to be bimodal
(i.e., highly peaked in the cluster centers and peripheries, and significantly lower at inter-
mediate radii) in at least 4 GCs: M3 (Ferraro et al. 1997), 47 Tuc (Ferraro et al. 2004),
NGC 6752 (Sabbi et al. 2004), and M5 (Warren, Sandquist & Bolte 2006, hereafter W06).
Preliminary evidence of bimodality has also been found in M55 (Zaggia, Piotto & Capaccioli
1997). Dynamical simulations suggest that the bimodal radial distributions observed in M3,
47 Tuc and NGC 6752 (Mapelli et al. 2004, 2006) result from ∼ 40− 50% of MT-BSS with
the balance being COL-BSS. In this context, the case of ω Cen is atypical: the BSS radial
distribution in this cluster is flat (Ferraro et al. 2006b), and mass segregation processes have
not yet played a major role, thus implying that this system is populated by a vast majority
of MT-BSS (Mapelli et al. 2006). These results demonstrate that detailed studies of the BSS
radial distribution within GCs are very powerful tools for better understanding the complex
interplay between dynamics and stellar evolution in dense stellar systems.
In the present paper we extend this kind of investigation to M5 (NGC 5904). With HST-
WFPC2 and -ACS ultraviolet and optical high-resolution images of the core we have been
able to efficiently detect the BSS population even in the severely crowded central regions.
Moreover, with wide-field optical observations performed with ESO-WFI we sampled the
entire cluster extension. The combination of these two data sets allowed us to study the
dynamical properties of M5, accurately redetermining its center of gravity, its surface density
profile, and the BSS radial distribution over the entire cluster. The BSS population of M5
has been recently studied by W06, but we have extended the analysis to larger distances
from the cluster center, and we have used Monte-Carlo dynamical simulations to interpret
the observational results.
– 4 –
2. OBSERVATIONS AND DATA ANALYSIS
2.1. The data sets
The present study is based on a combination of two different photometric data sets:
1. The high-resolution set – It consists of a series of ultraviolet (UV) and optical images
of the cluster center obtained with HST-WFPC2 (Prop. 6607, P.I. Ferraro). To efficiently
resolve the stars in the highly crowded central regions, the Planetary Camera (PC, being the
highest resolution instrument: 0.′′046/pixel) has been pointed approximately on the cluster
center, while the three Wide Field Cameras (WF, having a lower resolution: 0.′′1/pixel) have
been used to sample the surrounding regions. Observations have been performed through
filter F255W (medium UV) in order to efficiently select the BSS and horizontal branch
(HB) populations, and through filters F336W (approximately corresponding to an U filter)
and F555W (V ) for the red giant branch (RGB) population and to guarantee a proper
combination with the ground-based data set (see below). The photometric reduction of the
high-resolution images was carried out using ROMAFOT (Buonanno et al. 1983), a package
developed to perform accurate photometry in crowded fields and specifically optimized to
handle under-sampled Point Spread Functions (PSFs; Buonanno & Iannicola 1989), as in
the case of the HST-WF chips.
To obtain a better coverage of the innermost regions of the cluster, we have also used a
set of public HST-WFPC2 and HST-ACS observations. The HST-WFPC2 data set has been
obtained through filters F439W (B) and F555W (V ) by Piotto et al. (2002), and because
of the different orientation of the camera, it is complementary to ours. Additional HST-ACS
data in filters F435W (B), F606W (V ), and F814W (I) have been retrieved from the ESO-
STECF Science Archive, and have been used to sample the central area not covered by the
WFPC2 observations. All the ACS images were properly corrected for geometric distortions
and effective flux (over the pixel area) following the prescriptions of Sirianni et al. (2005).
The photometric analysis was performed independently in the three drizzled images by using
the aperture photometry code SExtractor (Source-Extractor; Bertin & Arnouts 1996), and
adopting a fixed aperture radius of 2.5 pixels (0.125′′). The magnitude lists were finally
cross-correlated in order to obtain a combined catalog. The adopted combination of the
three HST data sets is sketched in Figure 1 and provided a good coverage of the cluster up
to r = 115′′.
2. The wide-field set - A complementary set of wide-field B and V images was secured by
using the Wide Field Imager (WFI) at the 2.2m ESO-MPI telescope during an observing run
in April 2000. Thanks to the exceptional imaging capabilities of WFI (each image consists of
a mosaic of 8 CCDs, for a global field of view of 34′×34′), these data cover the entire cluster
– 5 –
extension (see Figure 2, where the cluster is roughly centered on CCD #7). The raw WFI
images were corrected for bias and flat field, and the overscan regions were trimmed using
IRAF2 tools. The PSF fitting procedure was performed independently on each image using
DoPhot (Schechter, Mateo & Saha 1993). All the uncertain detections, usually caused by
photometric blends, stars near the CCD gaps or saturated stars, have been checked one by
one using ROMAFOT (Buonanno et al. 1983).
2.2. Astrometry and center of gravity
The HST+WFI catalog has been placed on the absolute astrometric system by adopting
the procedure already described in Ferraro et al. (2001, 2003). The new astrometric Guide
Star Catalog (GSC-II3) was used to search for astrometric standard stars in the WFI field of
view (FoV), and a cross-correlation tool specifically developed at the Bologna Observatory
(Montegriffo et al. 2003, private communication) has been employed to obtain an astrometric
solution for each of the 8 CCDs. Several hundred GSC-II reference stars were found in each
chip, thus allowing an accurate absolute positioning of the stars. Then, a few hundred stars
in common between the WFI and the HST FoVs have been used as secondary standards to
place the HST catalog on the same absolute astrometric system. At the end of the procedure
the global uncertainties in the astrometric solution are of the order of ∼ 0.′′2, both in right
ascension (α) and declination (δ).
Given the absolute positions of individual stars in the innermost regions of the cluster,
the center of gravity Cgrav has been determined by averaging coordinates α and δ of all
stars lying in the PC FoV following the iterative procedure described in Montegriffo et
al. (1995; see also Ferraro et al. 2003, 2004). In order to correct for spurious effects
due to incompleteness in the very inner regions of the cluster, we considered two samples
with different limiting magnitudes (m555 < 19.5 and m555 < 20), and we computed the
barycenter of stars for each sample. The two estimates agree within ∼ 1′′, giving Cgrav at
α(J2000) = 15h 18m 33.s53, δ(J2000) = +2o 4′ 57.′′06, with a 1σ uncertainty of 0.′′5 in both α
and δ, corresponding to about 10 pixels in the PC image. This value of Cgrav is located at
∼ 4′′ south-west (∆α = −4′′, ∆δ = −0.′′9) from that previously derived by Harris (1996) on
the basis of the surface brightness distribution.
2IRAF is distributed by the National Optical Astronomy Observatory, which is operated by the Associa-
tion of Universities for Research in Astronomy, Inc., under cooperative agreement with the National Science
Foundation.
3Available at http://www-gsss.stsci.edu/Catalogs/GSC/GSC2/GSC2.htm.
http://www-gsss.stsci.edu/Catalogs/GSC/GSC2/GSC2.htm
– 6 –
2.3. Photometric calibration and definition of the catalogs
The optical HST magnitudes (i.e., those obtained through the WFPC2 filters F439W
and F555W, and through ACS filters F435W, F606W, F814W), as well as the WFI B and
V magnitudes have been all calibrated on the catalog of Sandquist et al. (1996). The UV
magnitudes m160 and m255 have been calibrated to the Holtzman et al. (1995) zero-points
following Ferraro et al. (1997, 2001), while the U magnitude m336 has been calibrated to
Dolphin (2000).
In order to reduce spurious effects due to the low resolution of the ground-based obser-
vations in the most crowded regions of the cluster, we use only the HST data for the inner
115′′, this value being imposed by the FoV of the WFPC2 and ACS cameras (see Figure
1). In particular, we define as HST sample the ensemble of all the stars in the WFPC2
and ACS combined catalog having r ≤ 115′′ from the center, and as WFI sample all stars
detected with WFI at r > 115′′ (see Figure 2). The CMDs of the HST and WFI samples in
the (V, U − V ) and (V, B − V ) planes are shown in Figure 3.
2.4. Density profile
We have determined the projected density profile over the entire cluster extension, from
Cgrav out to ∼ 1400
∼ 23.′3, by direct star counts, considering only stars brighter than
V = 20 (see Figure 3) in order to avoid incompleteness biases. The brightest RGB stars that
are strongly saturated in the ACS data set have been excluded from the analysis, but since
they are few in number, the effect on the resulting density profile is completely negligible.
Following the procedure already described in Ferraro et al. (1999a, 2004), we have divided
the entire HST+WFI sample in 27 concentric annuli, each centered on Cgrav and split in
an adequate number of sub-sectors. The number of stars lying within each sub-sector was
counted, and the star density was obtained by dividing these values by the corresponding
sub-sector areas. The stellar density in each annulus was then obtained as the average of
the sub-sector densities, and its standard deviation was estimated from the variance among
the sub-sectors.
The radial density profile thus derived is plotted in Figure 4, where we also show the best-
fit mono-mass King model and the corresponding values of the core radius and concentration:
rc = 27
′′ (with a typical error of ∼ ±2′′) and c = 1.68, respectively. These values confirm
that M5 has not yet experienced core collapse, and they are in good agreement with those
quoted by McLaughlin & van der Marel (2005, rc = 26.
′′3 and c = 1.71), and marginally
consistent with those listed by Harris (1996, rc = 25.
′′2 and c = 1.83), both derived from
– 7 –
the surface brightness profile. Our value of rc corresponds to ∼ 1 pc assuming the distance
modulus (m−M)0 = 14.37 (d ∼ 7.5 Kpc, Ferraro et al. 1999b).
3. DEFINITION OF THE SAMPLES
In order to study the BSS radial distribution and detect possible peculiarities, both the
BSS and a reference population must be properly defined. Since the HST and the WFI data
sets have been observed in different photometric bands, different selection boxes are needed
to separate the samples in the CMDs. The adopted strategy is described in the following
sections (see also Ferraro et al. 2004 for a detailed discussion of this issue).
3.1. The BSS selection
At UV wavelengths BSS are among the brightest objects in a GC, and RGB stars are
particularly faint. By combining these advantages with the high-resolution capability of HST,
the usual problems associated with photometric blends and crowding in the high density
central regions of GCs are minimized, and BSS can be most reliably recognized and separated
from the other populations in the UV CMDs. For these reasons our primary criterion for the
definition of the BSS sample is based on the position of stars in the (m255, m255 −U) plane.
In order to avoid incompleteness bias and the possible contamination from TO and sub-giant
branch stars, we have adopted a limiting magnitude m255 = 18.35, roughly corresponding
to 1 magnitude brighter than the cluster TO. This is also the limiting magnitude used by
W06, facilitating the comparison with their study. The resulting BSS selection box in the
UV CMD is shown in Figure 5. Once selected in the UV CMD, the bulk of the BSS lying in
the field in common with the optical-HST sample has been used to define the selection box
and the limiting magnitude in the (B, B − V ) plane. The latter turns out to be B ≃ 17.85,
and the adopted BSS selection box in the optical CMD is shown in Figure 6. The two stars
lying outside the selection box (namely BSS-19 and BSS-20 in Table 1) have been identified
as BSS from the (m255, m255−U) CMD. Indeed, they are typical examples of how the optical
magnitudes are prone to blend/crowding problems, while the BSS selection in UV bands is
much more secure and reliable. An additional BSS (BSS-47 in Table 1) lies near the edge
of the ACS FoV and has only V and I observations; thus it was selected in the (V, V − I)
plane (see Figure 7, where this BSS is shown together with the other 5 identified in the ACS
complementary sample).
With these criteria we have identified 60 BSS: 47 BSS in the HST sample (r ≤ 115′′)
– 8 –
and 13 in the WFI one. Their coordinates and magnitudes are listed in Table 1. Out of the
47 BSS identified in the HST sample, 41 are from the WFPC2 data set, and 6 from the ACS
catalog. As shown in Figure 1 their projected distribution is quite asymmetric with the N-E
sector seemingly underpopulated. The statistical significance of such an asymmetry appears
even higher if only the BSS outside the core are considered. However a quantitative discussion
of this topic is not warranted unless additional evidences supporting this anomalous spatial
distribution are collected. One of the inner BSS (BSS-29 in Table 1) lying at 21.′′76 from the
center, corresponds to the low-amplitude variable HST-V28 identified by Drissen & Shara
(1998)4. In the WFI sample (r > 115′′) we find 13 BSS, with a more symmetric spatial
distribution (see Figure 2). The most distant BSS (BSS-60 in Table 1, marked with an
empty triangle in Fig.6) lies at ∼ 24′ from the center, i.e., beyond the cluster tidal radius.
Hence, it might be an evaporating BSS previously belonging to the cluster. However, further
investigations are needed before firmly assessing this issue.
In order to perform a proper comparison with W06 study, we have transformed their
BSS catalog in our astrometric system, and we have found that 50 BSS of their bright sample
lie at r ≤ 115′′: 35 are from the HST sample, 13 from the Canada France Hawaii Telescope
(CFHT) data set, and 2 from the Cerro Tololo Inter-American Observatory (CTIO) sample;
in the outer regions (115′′ < r <∼ 425
′′) 9 BSS are identified, all from the CTIO data set.
By cross correlating W06 bright sample with our catalog we have found 43 BSS in
common (see Table 1), 37 at r ≤ 115′′ and 6 outward. In particular, 33 BSS out of the 41
(i.e., 80% of the total) that we have identified in the WFPC2-HST sample5 are found in
both catalogs, while 3 of our BSS belong 5 to their faint BSS sample (namely, BSS-27, 34,
and 40, corresponding to their Core BSS 70, 79, and 76, respectively), 5 of our BSS have
been missed in W06 paper, and 2 objects in their sample are classified as HB stars in our
study. This is probably due to different selection criteria, and/or small differences in the
measured magnitudes, caused by the different data reduction procedures and photometric
analysis. For example, W06 identify the BSS on the basis of both the UV and the optical
observations, while we select the BSS only in the UV plane whenever possible. Out of the
other 15 BSS found at r ≤ 115′′ in the ground-based CFHT/CTIO sample of W06, 8 BSS
(Core BSS 38–45 in their Table 2) clearly are false identifications. They are arranged in a
very unlikely ring around a strongly saturated star, as can be seen in Figure 8, where the
position in the sky of the 8 spurious BSS are overplotted on the CFHT image. Though they
4The observations presented here do not have the time coverage needed to properly search for BSS
variability.
5Note that the WFPC2-HST observations used in W06 and in the present study are the same.
– 9 –
clearly are spurious identifications, they still define a clean sequence in the (B, B−I) CMD,
nicely mimicking the BSS magnitudes and colors. As already discussed in previous papers,
this once again demonstrates how automatic procedures for the search of peculiar objects
are prone to errors, especially when using ground-based observations to probe very crowded
stellar regions. We emphasize that all the candidate BSS listed in our Table 1 have been
visually inspected evaluating the quality and the precision of the PSF fitting. This procedure
significantly reduces the possibility of introducing spurious objects in the sample. Out of the
remaining 7 BSS, 4 objects (namely their Core BSS 32, 30, 37 and 28) are also confirmed
by our ACS observations (BSS-42, 43, 44, and 45 respectively), while 2 others (their Core
BSS 27 and Ground BSS 6) are not found in the ACS data set, and the remaining one (their
Ground BSS 7) is not included in our observation FoV. In turn, two BSS identified in our
ACS data set (BSS 46 and 47) are missed in their sample. Concerning the BSS lying at
115′′ < r < 450′′, 6 objects (out of 9 found in both samples) are in common between the two
catalogs (see Table 1), one (BSS-55) belongs to W06 faint sample (their Ground BSS 23),
while the remaining 2 do not coincide. Moreover, 4 additional BSS have been identified at
r > 450′′ in our study.
3.2. The reference population
Since the HB sequence is bright and well separable in the UV and optical CMDs, we
chose these stars as the primary representative population of normal cluster stars to be used
for the comparison with the BSS data set. As with the BSS, the HB sample was first defined
in the (m255, m255 −U) plane, and the corresponding selection box in (B, B − V ) has then
been determined by using the stars in common between the UV and the optical samples. The
resulting selection boxes in both diagrams are shown in Figures 5 and 6, and are designed
to include the bulk of HB stars6. Slightly different selection boxes would include or exclude
a few stars only without affecting the results.
We have used WFI observations to roughly estimate the impact of possible foreground
field stars contamination on the cluster population selection. As shown in the right-hand
panel of Figure 6, field stars appear to define an almost vertical sequence at 0.4 < B−V < 1
in the (B, B − V ) CMD. Hence, they do not affect the BSS selection box, but marginally
contaminate the reddest end of the HB. In particular, 5 objects have been found to lie within
the adopted HB box in the region at r > rt sample by our observations (∼ 194 arcmin
this corresponds to 0.026 spurious HB stars per arcmin2. On the basis of this, 11 field stars
6The large dispersion in the redder HB stars arises because RR Lyrae variables are included.
– 10 –
are expected to ”contaminate” the HB population over the sampled cluster region (r < rt).
4. THE BSS RADIAL DISTRIBUTION
The radial distribution of BSS in M5 has been studied following the same procedure
previously adopted for other clusters (see references in Ferraro 2006; Beccari et al. 2006).
First, we have compared the BSS cumulative radial distribution to that of HB stars. A
Kolmogorov-Smirnov test gives a ∼ 10−4 probability that they are extracted from the same
population (see Figure 9). BSS are more centrally concentrated than HB stars at ∼ 4σ level.
For a more quantitative analysis, the surveyed area has been divided into 8 concentric
annuli, with radii listed in Table 2. The number of BSS (NBSS) and HB stars (NHB), as
well as the fraction of sampled luminosity (Lsamp) have been measured in the 8 annuli and
the obtained values are listed in Table 2. Note that HB star counts listed in the table are
already decontaminated from field stars, according to the procedure described in Section 3.2
(1, 2, and 8 HB stars in the three outer annuli have been estimated to be field stars). The
listed values have been used to compute the specific frequency FHBBSS ≡ NBSS/NHB, and the
double normalized ratio (see Ferraro et al. 1993):
Rpop =
(Npop/N
(Lsamp/L
tot )
, (1)
with pop = BSS, HB.
In the present study luminosities have been calculated from the surface density profile
shown in Figure 4. The surface density has been transformed into luminosity by means of a
normalization factor obtained by assuming that the value obtained in the core (r ≤ 27′′) is
equal to the sum of the luminosities of all the stars with V ≤ 20 lying in this region. The
distance modulus quoted in Section 2.4 and a reddening E(B−V ) = 0.03 have been adopted
(Ferraro et al. 1999b). The fraction of area sampled by the observations in each annulus has
been carefully computed, and the sampled luminosity in each annulus has been corrected for
incomplete spatial coverage (in the case of annuli 3 and 8; see Figures 1 and 2).
The resulting radial trend of RHB is essentially constant with a value close to unity
over the surveyed area (see Figure 10). This is just what expected on the basis of the
stellar evolution theory, which predicts that the fraction of stars in any post-main sequence
evolutionary stage is strictly proportional to the fraction of the sampled luminosity (Renzini
& Fusi Pecci 1988). Conversely, BSS follow a completely different radial distribution. As
shown in Figure 10 the specific frequency RBSS is highly peaked at the cluster center (a
– 11 –
factor of ∼ 3 higher than RHB in the innermost bin), decreases to a minimum
7 at r ≃ 10 rc,
and rises again outward. The same behavior is clearly visible also in Figure 11, where the
population ratio NBSS/NHB is plotted as a function of r/rc.
Note that the region between 800′′ and rt ≃ 1290
′′ (and thus also BSS-59, that lies at
r ≃ 995.′′5) has not been considered in the analysis, since our observations provide a poor
sampling of this annulus: only 35% of its area, corresponding to ∼ 0.4% of the total sampled
light, is covered by the WFI pointing. However, for sake of completeness, we have plotted in
Figure 12 the corresponding value of FHBBSS even for this annulus (empty circle in the upper
panel): as can be seen, there is a hint for a flattening of the BSS radial distribution in the
cluster outskirts.
4.1. Dynamical simulations
Following the same approach as Mapelli et al. (2004, 2006), we now exploit dynamical
simulations to derive some clues about the BSS formation mechanisms from their observed
radial distribution. We use the Monte-Carlo simulation code originally developed by Sig-
urdsson & Phinney (1995) and upgraded in Mapelli et al. (2004, 2006). In any simulation
run we follow the dynamical evolution of N BSS within a background cluster, taking into
account the effects of both dynamical friction and distant encounters. We identify as COL-
BSS those objects having initial positions ri
∼ rc, and as MT-BSS stars initially lying at
ri ≫ rc (this because stellar collisions are most probable in the central high-density regions
of the cluster, while primordial binaries most likely evolve in isolation in the periphery).
Within these two radial ranges, all initial positions are randomly generated following the
probability distribution appropriate for a King model. The BSS initial velocities are ran-
domly extracted from the cluster velocity distribution illustrated in Sigurdsson & Phinney
(1995), and an additional natal kick is assigned to the COL-BSS in order to account for the
recoil induced by the encounters. Each BSS has characteristic mass M and lifetime tlast. We
follow their dynamical evolution in the cluster (fixed) gravitational potential for a time ti
(i = 1, N), where each ti is a randomly chosen fraction of tlast. At the end of the simulation
we register the final positions of BSS, and we compare their radial distribution with the
observed one. We repeat the procedure until a reasonable agreement between the simulated
and the observed distributions is reached; then, we infer the percentage of collisional and
mass-transfer BSS from the distribution of the adopted initial positions in the simulation.
For a detailed discussion of the ranges of values appropriate for these quantities and
7Note that no BSS have been found between 3.′5 and 5′.
– 12 –
their effects on the final results we refer to Mapelli et al. (2006). Here we only list the
assumptions made in the present study:
– the background cluster is approximated with a multi-mass King model, determined as
the best fit to the observed profile8. The cluster central velocity dispersion is set to
σ = 6.5 km s−1 (Dubath et al. 1997), and, assuming 0.5M⊙ as the average mass of the
cluster stars, the central stellar density is nc = 2× 10
4 pc−3 (Pryor & Meylan 1993);
– the COL-BSS are distributed with initial positions ri ≤ rc and are given a natal kick
velocity of 1× σ;
– initial positions ranging between 5 rc and rt (with the tidal radius rt ≃ 48 rc) have been
considered for MT-BSS in different runs;
– BSS masses have been fixed to M = 1.2M⊙ (Ferraro et al. 2006a), and their charac-
teristic lifetime to tlast = 2 Gyr;
– in each simulation run we have followed the evolution of N = 10, 000 BSS.
The simulated radial distribution that best reproduces the observed one (with a reduced
χ2 ≃ 0.6) is shown in Figure 11 (solid line) and is obtained by assuming that ∼ 80% of the
BSS population was formed in the core through stellar collisions, while only ∼ 20% is made
of MT-BSS. A higher fraction ( >∼ 40%) of MT-BSS does not correctly reproduce the steep
decrease of the distribution and seriously overpredict the number of BSS at r ∼ 10 rc, where
no BSS at all are found, but it nicely matches the observed upturning point at r ≃ 13 rc (see
the dashed line in Figure 11). On the other hand, a population of only COL-BSS is unable to
properly reproduce the external upturn of the distribution (see the dotted line in Figure 11),
and 100% of MT-BSS is also totally excluded. Assuming heavier BSS (up to M = 1.5M⊙) or
different lifetimes tlast (between 1 and 4 Gyr) does not significantly change these conclusions,
since both these parameters mainly affect the external part of the simulated BSS distribution.
Thus, an appreciable effect can be seen only in the case of a relevant upturn, and negligible
variations are found in the best-fit case and when assuming 100% COL-BSS. The effect starts
to be relevant in the simulations with 40% or more MT-BSS, which are however inconsistent
with the observations at intermediate radii (see above).
By using the simulations and the dynamical friction timescale (from, e.g., Mapelli et al.
2006), we have also computed the radius of avoidance of M5. This is defined as the char-
acteristic radial distance within which all MT-BSS are expected to have already sunk to
8By adopting the same mass groups as those of Mapelli et al. (2006), the resulting value of the King
dimensionless central potential is W0 = 9.7
– 13 –
the cluster core, because of mass segregation processes. Assuming 12 Gyr for the age of
M5 (Sandquist et al. 1996) and 1.2M⊙ for the BSS mass, we find that ravoid ≃ 10 rc. This
nicely corresponds to the position of the minimum in the observed BSS radial distribution,
in agreement with the findings of Mapelli et al. (2004, 2006).
5. SUMMARY AND DISCUSSION
In this paper we have used a combination of HST UV and optical images of the cluster
center and wide-field ground-based observations covering the entire cluster extension to de-
rive the main structural parameters and to study the BSS population of the galactic globular
cluster M5.
The accurate determination of the cluster center of gravity from the high-resolution
data gives α(J2000) = 15h 18m 33.s53, δ(J2000) = +2o 4′ 57.′′06, with a 1σ uncertainty of 0.′′5
in both α and δ. The cluster density profile, determined from direct star counts, is well fit
by a King model with core radius rc = 27
′′ and concentration c = 1.68, thus suggesting that
M5 has not yet suffered the core collapse.
The BSS population of M5 amounts to a total of 59 objects, with a quite asymmetric
projected distribution (see Figure 1) and a high degree of segregation in the cluster center.
With respect to the sampled luminosity and to HB stars, the BSS radial distribution is
bimodal: highly peaked at r <∼ rc, decreasing to a minimum at r ≃ 10 rc, and rising again
outward (see Figures 10 and 11).
The comparison with results of W06 has revealed that 43 (out of 59) bright BSS iden-
tified by these authors at r <∼ 450
′′ are in common with our sample. Moreover, 4 additional
stars classified as faint BSS in their study are in common with our BSS sample at r <∼ 450
Considering that we find 56 BSS within the same radial distance from the center, this corre-
sponds to 84% matching of our catalogue. The discrepancies are explained by different data
reduction procedures, photometric analysis, and adopted selection criteria, other than the
spurious identification of 8 BSS by W06, due a strongly saturated star in their sample. The
central peak of the RBSS distribution in our study is slightly higher (but compatible within
the error bar) compared to that of W06, and we extend the analysis to larger distance from
the center (out to r > 800′′), thus unveiling the external upturn and the possible flattening
of the BSS distribution in the cluster outskirts.
Moreover, we have compared the BSS radial distribution of M5 with that observed
in other GCs studied in a similar way. In Figure 12 we plot the specific frequency FHBBSS
as a function of (r/rc) for M5, M3, 47 Tuc, and NGC 6752. Such a comparison shows
– 14 –
that the BSS radial distributions in these clusters are only qualitatively similar, with a
high concentration at the center and an upturn outward. However, significant quantitative
differences are apparent: (1) the FHBBSS peak value, (2) the steepness of the decreasing branch of
the distribution, (3) the radial position of the minimum (marked by arrows in the figure), and
(4) the extension of the “zone of avoidance,” i.e., the intermediate region poorly populated
by BSS. In particular M5 shows the smallest FHBBSS peak value: it turns out to be ∼ 0.24,
versus a typical value >∼ 0.4 in all the other cases. It also shows the mildest decreasing slope:
at r ≈ 2 rc the specific frequency in M5 is about a half of the peak value, while it decreases
by a factor of 4 in all the other clusters. Conversely, it is interesting to note that the value
reached by FHBBSS in the external regions is ∼ 50-60% of the central peak in all the studied
clusters. Another difference between M5 and the other systems concerns the ratio between
the radius of avoidance and the tidal radius: ravoid ≃ 0.2 rt for M5, while ravoid
∼ 0.13 rt for
47 Tuc, M3, and NGC 6752 (see Tables 1 and 2 in Mapelli et al. 2006).
The dynamical simulations discussed in Section 4.1 suggest that the majority of BSS in
M5 are collisional, with a content of MT-BSS ranging between 20% and 40% of the overall
population. This fraction seems to be smaller than that (40-50%) derived for M3, 47 Tuc
and NGC 6752 by Mapelli et al. (2006), in qualitative agreement with the smaller value of
ravoid/rt estimated for M5, which indicates that the fraction of cluster currently depopulated
of BSS is larger in this system than in the other cases. More in general, the results shown
in Figure 11 exclude a pure collisional BSS content for M5.
Our study has also revealed the presence of a candidate BSS at ∼ 24′ from the center,
i.e., beyond the cluster tidal radius (see Figures 2 and 6 and BSS-59 in Table 1). If confirmed,
this could represent a very interesting case of a BSS previously belonging to M5 and then
evaporating from the cluster (a BSS kicked off from the core the because of dynamical
interactions?).
This research was supported by Agenzia Spaziale Italiana under contract ASI-INAF
I/023/05/0, by the Istituto Nazionale di Astrofisica under contract PRIN/INAF 2006, and
by the Ministero dell’Istruzione, dell’Università e della Ricerca. RTR is partially funded by
NASA through grant number GO-10524 from the Space Telescope Science Institute. We
thank the referee E. Sandquist for the careful reading of the manuscript and the useful
comments and suggestions that significantly improved the presentation of the paper.
REFERENCES
Albrow, M. D., et al. 2001, ApJ, 559, 1060
– 15 –
Bailyn, C. D. 1994, AJ, 107, 1073
Beccari, G., Ferraro, F. R., Lanzoni, B., & Bellazzini, M., 2006, ApJ, 652, L121
Bellazzini, M., Fusi Pecci, F., Messineo, M., Monaco, L., & Rood, R. T. 2002, AJ, 123, 1509
Bertin, E., & Arnouts, S. 1996, A&AS, 117, 393
Buonanno, R., Buscema, G., Corsi, C. E., Ferraro, I., & Iannicola, G. 1983, A&A, 126, 278
Buonanno, R., Iannicola, G. 1989, PASP, 101, 294
Dolphin, A. E. 2000, PASP, 112, 1383
Drissen, L., & Shara, M. M. 1998, AJ, 115, 725
Dubath P., Meylan G., Mayor M., 1997, A&A, 324, 505
Ferraro, F. R., Fusi Pecci, F., Cacciari, C., Corsi, C., Buonanno, R., Fahlman, G. G., &
Richer, H. B. 1993, AJ, 106, 2324
Ferraro, F. R., Paltrinieri, B., Fusi Pecci, F., Cacciari, C., Dorman, B., Rood, R. T., Buo-
nanno, R., Corsi, C. E., Burgarella, D., & Laget, M., 1997, A&A, 324, 915
Ferraro, F. R., Paltrinieri, B., Rood, R. T., Dorman, B. 1999a, ApJ 522, 983
Ferraro F. R., Messineo M., Fusi Pecci F., De Palo M. A., Straniero O., Chieffi A., Limongi
M. 1999b, AJ, 118, 1738
Ferraro, F. R., D’Amico, N., Possenti, A., Mignani, R. P., & Paltrinieri, B. 2001, ApJ, 561,
Ferraro, F. R., Sills, A., Rood, R. T., Paltrinieri, B., & Buonanno, R. 2003, ApJ, 588, 464
Ferraro, F. R., Beccari, G., Rood, R. T., Bellazzini, M., Sills, A., & Sabbi, E. 2004, ApJ,
603, 127
Ferraro, F. R., 2006, in Resolved Stellar Populations, ASP Conference Series, 2005, D. Valls-
Gabaud & M. Chaves Eds., astro-ph/0601217
Ferraro, F. R., et al. 2006a, ApJ, 647, L53
Ferraro, F. R., Sollima, A., Rood, R. T., Origlia, L., Pancino, E., & Bellazzini, M. 2006b,
ApJ, 638, 433
http://arxiv.org/abs/astro-ph/0601217
– 16 –
Fusi Pecci, F., Ferraro, F. R., Corsi, C. E., Cacciari, C., Buonanno, R. 1992, AJ, 104, 1831
Harris, W.E. 1996, AJ, 112, 1487
Holtzman, J. A., Burrows, C. J., Casertano, S., Hester, J. J., Trauger, J. T., Watson, A. M.,
& Worthey, G. 1995, PASP, 107, 1065
Lombardi, J. C. Jr., Rasio, F. A., Shapiro, S. L. 1995, ApJ, 445, L117
Mapelli, M., Sigurdsson, S., Colpi, M., Ferraro, F. R., Possenti, A., Rood, R. T., Sills, A.,
& Beccari, G. 2004, ApJ, 605, L29
Mapelli, M., Sigurdsson, S., Ferraro, F. R., Colpi, M., Possenti, A., & Lanzoni, B. 2006,
MNRAS, 373, 361
McLaughlin, D. E., & van der Marel, R. P. 2005, ApJS, 161, 304
Montegriffo, P., Ferraro, F. R., Fusi Pecci, F., & Origlia, L. 1995, MNRAS, 276, 739
Piotto, G., et al. 2002, A&A, 391, 945
Pryor C., & Meylan G., 1993, Structure and Dynamics of Globular Clusters. Proceedings of a
Workshop held in Berkeley, California, July 15-17, 1992, to Honor the 65th Birthday of
Ivan King. Editors, S.G. Djorgovski and G. Meylan; Publisher, Astronomical Society
of the Pacific, Vol. 50, 357
Renzini, A., & Fusi Pecci, F. 1988, ARA&A, 26, 199
Sabbi, E., Ferraro, F. R., Sills, A., Rood, R. T., 2004, ApJ 617, 1296
Sandquist, E. L., Bolte, M., Stetson, P. B.; Hesser, J. E. 1996, ApJ, 470, 910
Sarna, M. J., & de Greve, J. P. 1996, QJRAS, 37, 11
Schechter, P. L., Mateo, M., & Saha, A. 1993, PASP, 105, 1342
Shara, M. M., Saffer, R. A., & Livio, M. 1997, ApJ, 489, L59
Sigurdsson S., Phinney, E. S., 1995, ApJS, 99, 609
Sirianni, M., et al. 2005, PASP, 117, 1049
Warren, S. R., Sandquist, E. L., & Bolte, M., 2006, ApJ 648, 1026 (W06)
Zaggia, S. R., Piotto, G., & Capaccioli, M., 1997, A&A, 327, 1004
– 17 –
Fig. 1.— Map of the HST sample. The heavy solid line delimits the HST-WFPC2 FoV
of our UV observations (Prop. 6607), the dashed line bounds the FoV of the optical HST-
WFPC2 observations by Piotto et al. (2002), and the dotted line marks the edge of the
complementary ACS data set. The derived center of gravity Cgrav is marked with a cross.
BSS (heavy dots) and the concentric annuli used to study their radial distribution (cfr. Table
2) are also shown. The inner and outer annuli correspond to r = rc = 27
′′ and r = 115′′,
respectively.
This preprint was prepared with the AAS LATEX macros v5.2.
– 18 –
Fig. 2.— Map of the WFI sample. All BSS detected in the WFI sample are marked as heavy
dots, and the concentric annuli used to study their radial distribution are shown as solid lines,
with the inner and outer annuli corresponding to r = 115′′ and r = 800′′, respectively (cfr.
Table 2). The circle corresponding to the tidal radius (rt ≃ 21.
′5) is also shown as dashed-
dotted line. The BSS lying beyond rt might represent a BSS previously belonging to M5
and now evaporating from the cluster.
– 19 –
– 20 –
Fig. 3.— Optical CMDs of the WFPC2-HST and the WFI samples. The hatched regions
indicate the magnitude limit (V ≤ 20) adopted for selecting the stars used to construct the
cluster surface density profile.
– 21 –
Fig. 4.— Observed surface density profile (dots and error bars) and best-fit King model (solid
line). The radial profile is in units of number of stars per square arcseconds. The dotted line
indicates the adopted level of the background, and the model characteristic parameters (core
radius rc, concentration c, dimensionless central potential W0) are marked in the figure. The
lower panel shows the residuals between the observations and the fitted profile at each radial
coordinate.
– 22 –
Fig. 5.— CMD of the ultraviolet HST sample. The adopted magnitude limit and selection
box used for the definition of the BSS population are shown. The resulting fiducial BSS are
marked with empty circles. The open square corresponds to the variable BSS identified by
Drissen & Shara (1998). The box adopted for the selection of HB stars is also shown.
– 23 –
Fig. 6.— CMD of the optical HST-WFPC2 and WFI samples. The adopted BSS and HB
selection boxes are shown, and all the BSS identified in these samples are marked with the
empty circles. The two BSS not included in the box in the left-hand panel lie well within
the selection box in the UV plane and are therefore considered as fiducial BSS. The empty
triangle in the right-hand panel corresponds to the BSS identified beyond the cluster tidal
radius, at r ≃ 24′.
– 24 –
Fig. 7.— CMD of the ACS complementary sample. The BSS selection box is shown, and
the resulting fiducial BSS are marked with empty circles.
– 25 –
Fig. 8.— Left-hand panel: position of the 8 false BSS (marked with white circles) as derived
from Table 2 of W06, overplotted to the CFHT image (units are the same as in their Figure
1). As can be seen, a heavily saturated star is responsible for the false identification. Right-
hand panel: location of the 8 false BSS (empty circles) in the (B, B − I) plane, as derived
from Table 2 of W06 (cfr. to their Fig. 2).
– 26 –
Fig. 9.— Cumulative radial distribution of BSS (solid line) and HB stars (dashed line) as
a function of the projected distance from the cluster center for the combined HST+WFI
sample. The two distributions differ at ∼ 4σ level.
– 27 –
Fig. 10.— Radial distribution of the BSS and HB double normalized ratios, as defined in
equation (1), plotted as a function of the radial coordinate expressed in units of the core
radius. RHB (with the size of the rectangles corresponding to the error bars computed
as described in Sabbi et al. 2004) is almost constant around unity over the entire cluster
extension, as expected for any normal, non-segregated cluster population. Instead, the radial
trend of RBSS (dots with error bars) is completely different: highly peaked in the center (a
factor of ∼ 3 higher than RHB), decreasing at intermediate radii, and rising again outward.
– 28 –
Fig. 11.— Observed radial distribution of the specific frequency NBSS/NHB (filled circles
with error bars), as a function of r/rc. The simulated distribution that best reproduces the
observed one is shown as a solid line and is obtained by assuming 80% of COL-BSS and 20%
of MT-BSS. The simulated distributions obtained by assuming 40% of MT-BSS (dashed line)
and 100% COL-BSS (dotted line) are also shown.
– 29 –
Fig. 12.— Radial distribution of the population ratio NBSS/NHB for M5, M3, 47 Tuc, and
NGC 6752, plotted as a function of the radial distance from the cluster center, normalized to
the core radius rc (from Mapelli et al. 2006, rc ≃ 30
′′, 21′′, 28′′ for M3, 47 Tuc, and NGC 6752,
respectively). The arrows indicate the position of the minimum of the distribution in each
case. The outermost point shown for M5 (empty circle) corresponds to BSS-58, lying at
r ≃ 995′′. This star has not been considered in the quantitative study of the BSS radial
distribution since only a negligible fraction of the annuls between 800′′ and rt is sampled by
our observations.
– 30 –
Table 1. The BSS population of M5
Name RA DEC m255 U B V I W06
[degree] [degree]
BSS-1 229.6354506 2.0841090 16.52 16.15 15.88 15.71 - CR2
BSS-2 229.6388102 2.0849660 17.95 17.38 17.40 17.04 - CR4
BSS-3 229.6383433 2.0842640 18.21 17.63 17.64 17.32 - CR3
BSS-4 229.6416234 2.0851791 17.59 17.22 17.05 16.90 - CR5
BSS-5 229.6416518 2.0836794 16.28 15.99 15.79 15.70 - CR1
BSS-6 229.6381953 2.0810119 17.36 16.99 16.81 16.65 - CR21
BSS-7 229.6403657 2.0824062 17.40 17.07 16.97 16.76 - CR12
BSS-8 229.6412279 2.0823768 17.91 17.47 17.41 17.15 - CR13
BSS-9 229.6376256 2.0793288 17.84 17.12 16.99 16.77 - CR23
BSS-10 229.6401139 2.0794858 17.57 16.98 16.87 16.62 - CR22
BSS-11 229.6396566 2.0784944 17.51 17.20 17.12 16.92 - CR24
BSS-12 229.6432834 2.0797197 18.12 17.64 17.78 17.54 - -
BSS-13 229.6384406 2.0776614 17.36 16.88 16.88 16.59 - CR25
BSS-14 229.6274500 2.0864896 18.07 17.63 17.64 17.33 - CR8
BSS-15 229.6204246 2.0879629 18.33 17.61 17.75 17.36 - CR11
BSS-16 229.6209379 2.0917858 17.80 17.28 17.26 16.98 - CR18
BSS-17 229.6264834 2.0960870 16.32 16.22 16.20 16.13 - CR20
BSS-18 229.6368731 2.0896002 16.56 16.30 16.11 16.01 - CR14
BSS-19 229.6367309 2.0917639 18.27 17.35 17.58 17.07 - CR17
BSS-20 229.6345837 2.0906438 17.88 16.81 16.96 16.43 - CR16
BSS-21 229.6382677 2.0934706 18.25 17.58 17.71 17.35 - CR19
BSS-22 229.6340227 2.0853879 17.67 17.32 17.22 17.03 - CR7
BSS-23 229.6332685 2.0875294 17.69 17.34 17.21 17.08 - CR10
BSS-24 229.6366685 2.0807168 18.23 17.78 17.67 17.37 - -
BSS-25 229.6393544 2.0762832 18.11 17.79 17.72 17.50 - -
BSS-26 229.6378381 2.0779999 17.86 17.52 17.43 17.27 - -
BSS-27 229.6349851 2.0807202 18.17 17.51 17.74 17.30 - CR70
BSS-28 229.6397645 2.0736403 18.19 17.60 17.69 17.28 - CR33
BSS-29 229.6370495 2.0770798 16.83 16.56 16.57 17.75 - CR26
BSS-30 229.6358816 2.0747883 18.25 17.81 17.79 17.51 - CR31
BSS-31 229.6361653 2.0720147 18.29 17.77 17.81 17.47 - CR36
BSS-32 229.6339822 2.0723032 16.73 16.10 16.16 15.95 - CR35
BSS-33 229.6281392 2.0756490 17.74 17.41 17.22 17.09 - CR29
BSS-34 229.6241278 2.0750261 18.21 17.50 17.65 17.27 - CR79
BSS-35 229.6332759 2.0603761 17.48 17.17 16.95 16.86 - CR48
BSS-36 229.6270877 2.0662947 17.33 17.18 17.06 16.95 - CR47
BSS-37 229.6244175 2.0693612 16.89 16.41 16.51 15.71 - CR46
BSS-38 229.6180419 2.0724090 17.37 17.23 17.12 17.00 - CR34
– 31 –
Table 1—Continued
BSS-39 229.6311963 2.0857800 18.31 17.33 17.40 16.76 - -
BSS-40 229.6297499 2.0664961 18.16 17.58 - 17.27 - CR76
BSS-41 229.6443367 2.0872809 - - 17.50 17.23 - CR9
BSS-42 229.6448646 2.0738335 - - 16.53 16.06 15.95 CR32
BSS-43 229.6460645 2.0748695 - - 16.64 16.44 16.66 CR30
BSS-44 229.6481631 2.0718829 - - 16.72 16.61 16.87 CR37
BSS-45 229.6433942 2.0760163 - - 17.03 16.79 16.91 CR28
BSS-46 229.6439884 2.0775670 - - 17.44 16.99 16.81 -
BSS-47 229.6180420 2.0598328 - - - 17.18 17.12 -
BSS-48 229.6092873 2.1680914 - - 16.85 16.68 - OR2
BSS-49 229.6723094 2.0882827 - - 16.94 16.64 - OR9
BSS-50 229.6006551 2.0814678 - - 17.00 16.74 - OR10
BSS-51 229.6669956 1.9781808 - - 17.20 16.74 - OR1
BSS-52 229.5949935 2.0469325 - - 17.69 17.46 - OR4
BSS-53 229.6706625 2.0695464 - - 17.82 17.50 - -
BSS-54 229.6667908 2.1149550 - - 17.82 17.72 - -
BSS-55 229.7370667 2.0323392 - - 17.80 17.42 - OR23
BSS-56 229.5476990 2.0112610 - - 16.88 16.60 - OR5
BSS-57 229.6711255 1.9415566 - - 16.98 16.64 - -
BSS-58 229.4381714 2.0302088 - - 17.75 17.33 - -
BSS-59 229.7408412 2.3399166 - - 17.49 17.08 - -
BSS-60 229.3218200 2.3271022 - - 16.34 16.09 - -
Note. — The first 41 BSS have been identified in the WFPC2 sample; BSS-42–
46 are from the complementary ACS observations; BSS-47–59 are from the WFI
data-set. BSS-59 lies beyond the cluster tidal radius, at ∼ 24′ from the center. The
last column list the corresponding BSS in W06 sample, with ”CR” indicating their
”Core BSS” and ”OR” their ”Outer Region BSS”.
– 32 –
Table 2. Number counts of BSS and HB
stars
′′ re
′′ NBSS NHB L
samp/L
0 27 22 94 0.14
27 50 15 94 0.16
50 115 10 135 0.26
115 150 3 46 0.09
150 210 2 52 0.10
210 300 0 45† 0.10
300 450 4 42† 0.09
450 800 2 38† 0.06
Note. — † The NHB values listed here
are those corrected for field contamination
(i.e., 1, 2 and 8 stars have been subtracted
to the observed number counts in these
three external annuli, respectively).
	INTRODUCTION
	OBSERVATIONS AND DATA ANALYSIS
	The data sets
	Astrometry and center of gravity
	Photometric calibration and definition of the catalogs
	Density profile
	DEFINITION OF THE SAMPLES
	The BSS selection
	The reference population
	THE BSS RADIAL DISTRIBUTION
	Dynamical simulations
	SUMMARY AND DISCUSSION
