The WFPC2 UV Survey: the BSS population in NGC 5824
N. Sanna, E. Dalessandro, F. R. Ferraro, B. Lanzoni, P. Miocchi, R. W. O'Connell
aa r X i v : . [ a s t r o - ph . GA ] N ov The WFPC2 UV Survey: the BSS population in NGC 5824 N. Sanna , , E. Dalessandro , F.R. Ferraro , B. Lanzoni , P. Miocchi , R. W. O’Connell INAF-Osservatorio Astrofisico di Arcetri, largo Fermi 5, 50127 Firenze, Italy E − mail : sanna @ arcetri.astro.it Dipartimento di Fisica e Astronomia, Universit`a degli Studi di Bologna, via Ranzani 1,I–40127 Bologna, Italy Department of Astronomy, University of Virginia, P. O. Box 400325, Charlottesville, VA22904, USA
ABSTRACT
We have used a combination of high-resolution
Hubble Space Telescope
WFPC2 and wide-field ground-based observations, in ultraviolet and opticalbands, to study the blue straggler star population of the massive outer-haloglobular cluster NGC 5824, over its entire radial extent. We have computed thecenter of the cluster and constructed the radial density profile, from detailedstar counts. The profile is well reproduced by a Wilson model with a small core( r c ≃ . ′′ ) and a concentration parameter c ≃ .
74. We also present the first agedetermination for this cluster. From the comparison with isochrones, we havefound t = 13 ± . r ∼ ′′ , and an upturn at largeradii. In the framework of the dynamical clock defined by Ferraro et al. (2012),this feature suggests that NGC 5824 is a cluster of intermediate dynamical age. Subject headings:
Globular clusters: individual (NGC 5824); stars: evolution -binaries: general - blue stragglers Based on observations with the NASA/ESA
HST (Prop. 11975), obtained at the Space Telescope ScienceInstitute, which is operated by AURA, Inc., under NASA contract NAS5-26555.
1. INTRODUCTION
The high stellar densities typical of globular clusters (GCs) make stellar interactionsvery likely events. For this reason it is expected that in GCs stellar evolution is stronglyaffected by the environment. Indeed GCs are efficient furnaces of exotic populations, like X-ray binaries, millisecond pulsars and blue straggler stars (BSSs; Paresce et al. 1992; Bailyn1995; Bellazzini et al. 1995; Ferraro et al. 1995, 2001; Ransom et al. 2005; Pooley & Hut2006). BSSs constitute the largest population among these ′′ exotica”, thus representing acrucial probe of the GC internal dynamics. In the optical color-magnitude diagram (CMD)of stellar systems BSSs appear bluer and brighter than the main sequence (MS) stars, thusmimicking a population of objects younger and more massive than the normal MS turnoffstars. Indeed, direct measurements (Shara et al. 1997, Gilliland et al. 1998, De Marco et al.2004) have shown that BSSs are two to three times more massive than the average stars inGCs ( h m i ∼ . M ⊙ ).Two main scenarios have been proposed to explain the formation of BSSs: mass transferin binary systems (McCrea 1964; Zinn & Searle 1976) and stellar collisions (Hills & Day1976). The two formation channels are believed to co-exist within the same cluster (seefor example the case of M30 Ferraro et al. 2009; Li et al. 2013; Dalessandro et al. 2013)with efficiencies that may vary as a function of the environment (Bailyn 1992, Ferraro etal 1995). Collisional BSSs more likely form in the core or in very dense GCs because thedensities are higher and so are the probabilities of direct collisions (e. g. Mapelli et al. 2004).Mass transfer is expected to be the dominant formation channel in the cluster’s outskirtsor in loose environments (but see Knigge et al. 2009). Chemical signatures of the differentBSS formation mechanisms have been found in 47 Tucanae (Ferraro et al. 2006a) and inM30 (Lovisi et al. 2013) where a fraction of stars shows anomalies in carbon and oxygenabundances, which are expected to be signatures of the mass-transfer process (Sarna & deGreve 1996).Because of their mass and the mechanisms involved in their formation, BSSs are power-ful tracers of the evolution of internal dynamics of clusters. In particular, as recently shownby Ferraro et al. (2012, F12), their radial distribution provides us with a measure of the dy-namical friction efficiency. In fact, based on the observed shape of several BSS distributions,F12 grouped GCs in three main families. Family I is composed of clusters where the radialBSS distribution is indistinguishable from that of the reference stars, suggesting that theseclusters have not undergone mass-segregation yet ( ω Centauri, Ferraro et al. 2006b; NGC2419, Dalessandro et al. 2008a; Palomar 14, Beccari et al. 2011). Most of the clusters showa bimodal BSS radial distribution: a peak in the center, a clear dip at intermediate radii,and an upturn in the external regions. These clusters are members of Family II. Family III 3 –is composed by clusters with a monotonically decreasing BSS distribution (M79, Lanzoniet al. 2007a; M75, Contreras Ramos et al. 2012; M80 and M30, F12). In these clusterseven the most distant BSSs already drifted toward the cluster’s center. Different familiescorrespond to different dynamical age of the clusters. As proposed by F12, Family I systemsare dynamically young, Family II clusters have intermediate dynamical ages and the FamilyIII ones are the oldest.Being more massive than the average cluster stars, BSSs are typically found in theinnermost regions of GCs, where stellar crowding and the dominant luminosity contributionof red giant branch (RGB) stars make the construction of complete samples of BSSs quitedifficult in optical bands. Conversely, it is fairy easy in the UV bands (Paresce et al. 1991).In particular, high-resolution UV observations with the
Hubble Space Telescope (HST) openthe possibility to survey UV bright populations, including horizontal branch (HB; Ferraroet al. 1998; D’Cruz et al. 2000; Dalessandro et al. 2011; see also O’Connell et al. 1997 andSchiavon et al. 2012), BSSs (Ferraro et al. 1993; Ferraro et al. 1997) and post-asymptoticgiant branch (see for example Brown et al. 2008) stars, even in the innermost regions.In the framework of an extensive UV survey of more than 30 Galactic GCs conductedwith the Wide Field Planetary Camera 2 (WFPC2) on board HST (Prop 11975, PI Ferraro),here we present a detailed multiwavelength photometric analysis of NGC 5824. This poorlystudied cluster is located at ∼
26 kpc from the Galactic center (Harris 1996, 2010 version)and, after NGC 2419, it is the most luminous outer-halo GC ( M V = − .
85, Harris 1996).Newberg et al. (2009) suggested that NGC 5824 might be associated to the Cetus PolarStream and it could have once been a dwarf galaxy core (Georgiev et al. 2009). Using the Ca II triplet measured for 17 RGB stars, Saviane et al. (2012) found a possible intrinsic ironspread of σ ([ F e/H ] = 0 .
12) dex in this cluster. Grillmair et al. (1995) found that the surfacedensity profile follows a power law over almost the entire extent of the cluster. Lutzgendorfet al. (2013) have recently suggested the presence of a 2000 M ⊙ black hole in this clusters.The paper is organized as follows. Section 2 describes the data sets and the photometricand astrometric analysis. In Section 3 we determine the center of gravity and the radialdensity profile of the system. In Section 4 we present the first age determination for thiscluster. In Sections 5 and 6 we discuss the BSS properties and present our conclusions. Thesummary of the paper is presented in Section 7. 4 –
2. OBSERVATIONS AND DATA ANALYSIS
As in our previous studies (e.g. Sanna et al. 2012 and references therein), we haveused a combination of high resolution and wide-field data to resolve the stars in the centralregions and to cover the entire radial extension of the cluster at the same time. We used thebest quality HST and ground-based data available for this cluster.The
HST data set consists of a series of images collected with three different pointingsof the WFPC2 in several bands, ranging from the UV to the optical. Pointing A includes11 optical and UV images (Prop 11975, PI Ferraro) obtained through the filters F255W,F336W and F555W with total exposure times t exp = 7200 s, t exp = 2700 s and t exp = 200 s,respectively. Pointing B consists of 12 optical images obtained through the filters F336W,F439W and F555W (Prop 5902, PI Fahlman) with total exposure times t exp = 1200 s, t exp =3200 s and t exp = 300 s, respectively. Pointing C includes three images obtained through thefilter F439W (Prop 8095, PI Ibata) with total exposure time t exp = 1200 s. Pointings B andC exactly overlap each other. The fields of view (FOVs) of the three pointings are shown inFigure 1. The center of the cluster is located in the Planetary Camera (pixel scale ∼ . ′′ pixel − ) in all pointings (see Figure 1).The WFI data set is composed of data obtained with the Wide Filed Imager (WFI)at the 2.2 m ESO/MPI telescope. Two images per filter through the B and V bands withtotal exposure times t exp = 3600 s and t exp = 500 s, respectively, were retrieved from theESO/STECF Science Archive. WFI consists of 8 CCDs with a pixel scale of ∼ . ′′ pixel − . The 33 ′ × ′ field of view allowed a complete sampling of the cluster. The clustercenter is located in chip ∼
15 bright and almost isolated stars in each frame. Typically the samestars were chosen in each filter. With the obtained PSF models we performed a first PSF-fit on each single image by using ALLSTAR. As second step, we built a median image foreach filter using IRAF tools . Then for each median image we built a list of stars detectedabove a given background threshold. The lists thus obtained have been combined for eachchip. For pointings A and B we built a master-list frame including only stars detected in atleast two median images. With this criterion, and for pointing A in particular, we secured IRAF is distributed by he National Optical Astronomy Observatory, which is operated by the Associationof Universities for Research in Astronomy, Inc., under cooperative agreement with the National ScienceFoundation.
DAOMATCH/DAOMASTER packages the photometry obtained for each single chip has been reported to a commonreference frame (Stetson 2000) that we have chosen to be a 50 ′ × ′ image obtained combiningdifferent ground-based data. Once all the chips were on the same reference frame, we built amaster list composed by all the stars detected in at least two frames. These stars have thenbeen forced to each single frame by using the ALLFRAME package. This choice allowedus to get full advantage of the dithering strategy adopted for these observations and tocompletely fill the gaps between different chips.All the catalogues were put on the absolute astrometric system using more than 10000stars in common with the Guide Stars Catalogue (GSC2.3). As a first step we obtainedthe astrometric solution for the entire WFI catalogue by using the procedure described inFerraro et al. (2001, 2003) and the cross-correlation tool CataXcorr (Montegriffo, privatecommunication). The HST (x, y) coordinates were first transformed to the HST WorldCoordinate System coordinates and then they were reported to the absolute astrometricsystem by using the stars in common with WFI. At the end of the procedure the estimatederror in the absolute positions, both in right ascension ( α ) and declination ( δ ), is of about0 . ′′ .All the WFPC2 magnitudes ( m , m , m and m ) were calibrated to the VEGA-MAG system by using the prescription by Holtzman et al. (1995) and the zero points fromthe WFPC2 data handbook. We used the equations by Dolphin (2009) to correct for chargetransfer efficiency. We converted the B and V magnitudes of the wide-field catalogue to the m and m VEGAMAG system, respectively, by means of the following color equations: m = B + 0 . × ( B − V ) + 8 . m = V − . × ( B − V ) + 4 . http://documents.stsci.edu/hst/wfpc2/documents/handbooks/dhb We found 35 objects in common (typ-ically galaxies and quasars) which were excluded from the following analysis. The opticalCMDs for the HST and WFI data sets are shown in Figure 3.
3. CENTER OF GRAVITY AND RADIAL DENSITY PROFILE
We determined the center of gracity ( C grav ) of NGC5824 from resolved bright starstaking advantage of the high-resolution of the WFPC2 data. We followed the iterativeprocedure described by Montegriffo et al. (1995), as already done in previous studies (seefor example Sanna et al. 2012). We used the center reported by Harris (1996) as first guessof our iterative procedure. We averaged the positions α and δ of the stars contained withincircles of three different radii (8 ′′ , 10 ′′ and 12 ′′ ), until convergence was reached. We wereof course limited in this analysis by the size and the shape of the Planetary Camera FOV.In order to avoid any possible spurious effect due to incompleteness of the catalogue, weconsidered three samples with different limiting magnitudes ( m = 19 . , . , . ∼ . ′′ and their average was therefore assumed as C grav : α ( J .
0) = 15 h m . s , δ ( J .
0) = − ◦ ′ . ′′ . This new determination is located ∼ . ′′ South-West (∆ α ≃ − . ′′ , ∆ δ ≃ − . ′′ ) from the Harris center.Starting from C grav , we divided the entire data set in two sub-samples. From the HSTdata set we selected only stars with a distance r < ′′ from C grav (see Figure 1): this isnamed the inner sample , composed of 22242 stars. Even if the region at r < ′′ is notentirely sampled by the WFPC2 FOV, we conservatively preferred not to complement itwith ground-based data. Hence, from the WFI data set we considered all stars at r > ′′ (see Figure 2), thus defining the outer sample consisting of 48828 objects.We determined the projected density profile of NGC 5824 by measuring the star countsin concentric annuli covering the entire cluster extension, from C grav to r ∼ ′′ . In orderto limit the strong contamination from field stars in the most external regions, only fiducialRGB and HB stars (see Figure 3) have been taken into account. Starting from this firstselection and because of the high crowding affecting the very central regions ( r ∼ . ′′ ),we further limited our selections to stars with m < .
0. We divided the entire FOVin 19 annuli centered on C grav and each annulus was divided into two, three or four sub- http://ned.ipac.caltech.edu/ r > ′′ ) have been used to estimate the contributionof background stars.We have tried to reproduce the observed density profile by using both King (King 1966)and Wilson (Wilson 1975) models. These are widely used to describe stellar systems likeGCs that are thought to have reached a state of (quasi-)equilibrium. The projected densityprofile in both cases is characterized by a constant value in the innermost portion and adecreasing behavior outward, with the Wilson model showing a more extended outer region(for more details see, e.g., Miocchi et al. 2013). We find that the observed density profilecannot be reproduced by a standard mono-mass King model. Similar conclusions have beenreached by Grillmair et al. (1995) using photographic photometry and star counts. Figure4 shows the best fit obtained by using a King model: as can be seen, in the external regions( r > ′′ ) the observed star density profile shows a radial decrease significantly steeper thanpredicted by the model. The profile is instead nicely reproduced by a single mass Wilsonmodel with a concentration c ≃ .
74 and core radius (i.e., the radius at which the centralsurface density equals its central value) r c ≃ . ′′ (see Figure 5). The nominal limiting radius(at which the model density drops to zero) is located at r ∼ ′′ . These values are onlymarginally consistent with those ( c = 2 . r c = 3 . ′′ ) quoted by McLaughlin & van derMarel (2005) for the Wilson model fit to fit the surface brightness (instead of the surfacedensity) profile.
4. METALLICITY SPREAD AND AGE DETERMINATION
As quoted in Section 1, Saviane et al. (2012) recently suggested that NGC 5824 showsa possible metallicity spread of the order of 0.1 dex. The position and the morphology ofthe RGB is a sensible function of the metallicity (see, e.g., Valenti et al. 2004). Thus, ametallicity spread is expected to produce a measurable dispersion in the color distributionof the RGB in the ( m , m − m ) CMD and we can use our data to put constrains onthe possible metallicity spread.For a quantitative estimate, we followed the procedure described in previous papers(e.g., Ferraro et al. 1991, 1992) to determine the intrinsic width (IW) of the RGB. In orderto minimize at most any possible bias that can artificially broaden the color distribution of 8 –the RGB, we selected the most vertical portion of the RGB (18 < m < .
5) from allstars at r < ′′ observed with the Wide Field 3 chip of the WFPC2. We then computedthe distribution of the residuals in the ( m − m ) color with respect to the adoptedRGB mean ridge line (see the histogram in Figure 6). This has been compared with thedistribution of the intrinsic photometric errors (corresponding to IW=0, solid line), and withsuch a distribution convolved with a metallicity spread of 0.1 dex (dashed line; note thataccording with theoretical isochrones, δ ( m − m ) /δ [Fe/H]=0.0125). Based on a χ -test,the solid line turns out to better reproduce the observed distribution, meaning that the colordistribution of the selected RGB stars is consistent with no metallicity spread (IW=0) orwith a spread smaller than 0.1 dex.The collected data set also offers the opportunity to determine for the first time theage of NGC 5824. This is of great importance in the context of the formation scenarios ofour Galaxy. In fact, studies focused on the age-metallicity relation of Galactic GCs (e.g.,Dotter et al. 2010, 2011) show that two different formation histories can be distinguished: arapid chemical enrichment for the GCs in the inner regions of the Galaxy (at a galactocentricdistance R GC < ∼
20 and ∼
40 kpc from the Galactic center,both because they are few and because the available photometric data are typically notaccurate enough. In this context, the case of NGC 5824, which lies at R GC ∼
26 kpc (Harris1996) is quite interesting. In order to determine the age of this cluster, we have comparedthe optical HST CMD with isochrones from the Girardi’s database (Bressan et al. 2012).For a more accurate determination, we have excluded the first 30 ′′ , where the photometricquality is lower. We have adopted a distance modulus ( m − M ) = 17 .
53 and a reddening E ( B − V ) = 0 .
14 (Ferraro et al. 1999). The assumed metallicity is [Fe/H]= − .
91 (Harris1996). We have superimposed to the CMD four isochrones with different ages, rangingbetween 12.0 Gyr and 13.5 Gyr, stepped by 0.5 Gyr (see Figure 7). From this comparisonwe have determined t = 13 ± . R GC < R GC >
5. THE BSS POPULATION
As already done in previous studies (e.g. Lanzoni et al. 2007b; Dalessandro et al. 2008b;Sanna et al. 2012, and references therein), in order to perform a meaningful study of BSSs 9 –in terms of both specific frequency and radial distribution we need to select a population of”normal” cluster stars, like HB or RGB stars, which are expected to be distributed as thecluster light and, for this reason, represent a reference population.
As quoted in the Introduction, UV-CMDs are the best planes to study the hottest stellarpopulations in GCs. For this reason we selected the BSSs in the ( m , m − m ) CMD.In order to minimize possible contamination from the MS turnoff and sub-giant branch, welimited the sample to m < . < ( m − m ) < .
1. The adopted selection box is shown in Figure 9. Within these limits, weidentified 37 BSSs from the UV.Unfortunately, the UV data do not cover the entire extension of the cluster. Hence, asdone in previous papers (see Lanzoni et al. 2007b) we ”translated” the UV selection box inthe optical plane in the following way: we identified the position of the UV-selected BSSs inthe ( m , m − m ) HST CMD, then we defined the boundaries of the optical selectionbox as to include the bulk of the UV-selected BSSs. The resulting optical selection box isshown in Figure 10. From the portion of the inner sample not covered by pointing A wethus selected 3 stars, for a total of 40 BSSs in the inner sample .The same box has been used to selected BSSs in the outer sample . In this case, welimited the analysis to r < ′′ , the distance at which the cluster density becomes smallerthan the background density (see Figure 5). Even if this value is several times smaller thanthe estimated limiting radius (see Section 3), we conservatively preferred to adopt this limitto minimize the impact of contamination for Galactic field stars. In the outer sample wethen selected 23 BSSs. As representative of the normal cluster stars we considered both the HB and the RGBpopulations. Since NGC 5824 has an extended HB, the UV diagram is the best plane to selectthese stars. Following the same procedure adopted for BSSs, we defined the HB selectionbox in the UV plane (Figure 9) and we then converted this selection in the optical plane.NGC 5824 hosts 26 confirmed RR Lyrae stars (Samus et al. 2009), 19 of which are locatedin the FOV covered by our data (five in the inner sample and 14 in the outer sample ). Thesehave been included in the HB selection (see triangles in Figures 9 and 10). Considering the 10 –entire cluster, from C grav to r = 500 ′′ , we identified 819 HB stars, 557 in the inner sample and 262 in the complementary outer sample .To select the RGB stars we used the optical CMDs, where these objects are bright andthe branch well defined. Following the RGB mean ridge line, we limited our selection to m < . inner sample and 808 in the outer sample . As evident from the CMD shown in the right panel of Figure 3, the WFI data set isstrongly contaminated by field stars. For this reason we carefully estimated the expectednumber of field stars in each selection box. In order to statistically quantify the Galactic fieldcontamination we used the CMD obtained for r > ′′ , where field stars define two verticalsequences roughly located at 0 . < ( m − m ) < . < ( m − m ) < . ρ BSS ∼ . − , ρ HB ∼ . − , ρ RGB ∼ . − . Star counts afterdecontamination are: N BSS = 60, N HB = 759, N RGB = 2004.Following Dalessandro et al. (2013), we divided the FOV into five concentric annuli cen-tered on C grav and, for each of them, we randomly subtracted a number of stars constrainedby the computed field star densities. Figure 12 shows the statistically decontaminated cu-mulative radial distribution for the three selected populations. As evident, BSSs are morecentrally concentrated than RGB and HB stars. The Kolmogorov-Smirnov test gives a prob-ability of ∼ .
08 and ∼ .
18 that the BSS distribution is extracted from the same parentdistribution as the RGB and HB stars, respectively.For a more quantitative analysis, we computed the population ratios N BSS /N HB , N BSS /N RGB and N HB /N RGB in five concentric annuli centered in C grav . We adopted Poissonian errors forthe populations and their propagation for the population ratios. The star counts for eachannulus are listed in Table 1. Note that, due to the shape of the WFPC2 FOV, the annuli10 ′′ < r < ′′ and 30 ′′ < r < ′′ are not entirely covered by our data (see Figure 1). Theradial distributions of the population ratios are shown in Figure 13. The distribution of both N BSS /N HB and N BSS /N RGB is clearly bimodal, with a peak in the center, a minimum locatedat r min ≃ ′′ = 5 r c and a rising branch in the outer regions. In contrast, the N HB /N RGB ratio (bottom panel) has a flat behavior across the entire extension of the cluster, as expectedfor normal populations following the distribution of the light. 11 –We computed also the double normalized ratios for the three populations, defined as(Ferraro et al. 1993, 1997): R pop = N annpop /N totpop L annsamp /L totsamp . (1)We adopted Poissonian errors for the populations and the luminosity and their propagationfor the double normalized ratios.The luminosity in each annulus has been obtained by integrating the best-fit Wilsonprofile and appropriately taking into account the effective area covered by the observations.The computed luminosity ratios are reported in Table 1. Figure 14 shows the results. Asexpected by the stellar evolutionary theory (Renzini & Fusi Pecci 1988), the HB and RGBstars follow the distribution of the light: the radial distribution is constant (grey rectangles)at a value close to unity. Again, the BSS distribution is found to be bimodal.
6. DISCUSSION
Following the scenario proposed by F12, the shape of the BSS radial distributions canbe interpreted in terms of the dynamical age of stellar systems. In particular, the bimodalBSS radial distribution and the location of the minimum ( r min ≃ r c ) found for NGC 5824indicates that this cluster is a member of the intermediate dynamical-age systems ( FamilyII , in the F12 classification). The BSS distribution is very similar to that found in M55 (seeFigure 2 of F12), where the minimum is located at r min ∼ r c . The rising trend toward thecenter is very steep in both these clusters and the minimum is well defined. This means thatdynamical friction has already been effective in segregating BSSs toward the cluster center,but it has poorly affected the outskirts.As shown in F12, there is a strong relation between the core relaxation time ( t rc , oneof the classical theoretical indicator of the cluster dynamical age) and the position of theminimum of the BSS radial distribution measured in units of the cluster core radius ( r min /r c ).This allowed F12 to define the empirical dynamical clock : the position of r min /r c can be usedas a sort of clock time-hand to measure the dynamical age of clusters. In this framework, a flatBSS radial distribution, where r min cannot be defined, indicates clusters with a relaxationtime of the order of the age of the Universe ( Family I ), while a monotonic BSS radialdistribution with only a central peak indicates dynamically old clusters, where the action ofdynamical friction has operated out to the most remote regions of the cluster, segregatingthe entire BSS population in the center (
Family III ). Following F12, we have computed t rc for NGC 5824 by using equation (10) of Djorgovsky (1993), adopting the cluster structuralparameters obtained in Section 3. We adopted the reddening and distance modulus quoted 12 –by Ferraro et al. (1999), the central luminosity density listed by Harris (1996) and the totalcluster mass estimated by McLaughlin & van der Marel (2005). In Figure 15 we show theposition of NGC 5824 (black solid circle) in the “dynamical clock plane” ( t rc /t H as a functionof r min /r c ; see Figure 4 in F12), where t H = 13 . Family II , where the action of dynamicalfriction has already started to segregate BSSs (and binary systems of similar total mass)toward the cluster center. In this scenario, the most remote BSSs are thought to be stillevolving in isolation in the outer cluster regions.
7. SUMMARY
In this paper we have used a combination of HST UV and optical images to sample thecluster center, and wide-field ground-based optical observations covering the entire clusterextension to derive the main structural parameters of the globular cluster NGC 5824 and tostudy its BSS population.From the high-resolution data we derived the cluster center of gravity lying at α ( J .
0) =15 h m . s , δ ( J .
0) = − ◦ ′ . ′′ . We determined the radial density profile from starcounts, finding that it is best fitted by a Wilson model with core radius r c ≃ . ′′ andconcentration c ≃ . t = 13 ± . inner sample and 20 in the outer sample ) has beenidentified. The comparison between the radial distribution of BSSs and normal cluster stars(HB and RGB), as well as of the double normalized ratios shows that the BSS distributionis bimodal: peaked in the center, with a clear-cut dip at intermediate-small radii ( r min ≃ ′′ ≃ r c ), and with an upturn in the external regions. This suggests that NGC 5824 is anintermediate dynamical-age cluster. 13 – ACKNOWLEDGMENTS
This research is part of a Project COSMIC-LAB founded by The European ResearchCouncil (under contract ERC-2010-AdG-267675). Research at the University of Virginia wassupported in part by NASA grant GO-11975 from the Space Telescope Science Institute toR. T. Rood and R. W. O’Connell.
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This preprint was prepared with the AAS L A TEX macros v5.2. r ′′ i r ′′ e N BSS N fieldBSS N HB N fieldHB N RGB N fieldRGB L annsamp /L totsamp C grav .The inner sample (see Sect. 3) consists of all the stars (black dots) included in the FOV ofpointings B+C and within a distance r = 75 ′′ from C grav (solid circle). 18 –Fig. 2.— Map of the WFI data set. The outer sample (see Sect. 3) consists of all thestars observed at r > ′′ (most internal circle). The radial distribution of the various stellarpopulations (BSSs, RGB and HB stars) has been studied only for r < ′′ (black dots withinthe large solid circle), where the cluster density becomes comparable to that of field stars.To estimate the Galactic field contamination, we have used the stars observed at r > ′′ (dashed circle). Only stars with m <
22 are shown for the sake of clarity. 19 –Fig. 3.— Optical CMDs of NGC 5824 for the inner sample (left panel) and the outer sample (right panel). Only the stars included in the grey boxes and with m <
20 have been usedto determine the density profile. 20 –Fig. 4.— Observed surface density profile (empty squares) in units of number of stars persquare arcseconds. The dotted line indicates the adopted level of the Galactic background,computed as the average of the three outermost points. Solid circles show the background-subtracted density profile. The best-fit King model is plotted as a solid line and its residualswith respect to the observations are shown in the lower panel. The labels quote the valuesof the model concentration and core radius. 21 –Fig. 5.— As in Figure 4, but for the best-fit Wilson model. 22 –Fig. 6.— Observed color width of the RGB for the stars detected in the Wide Field 3 chip ofthe WFPC2, along the RGB at 18 ≤ m ≤ . r < ′′ . The observed distributionof the ( m − m ) color residuals with respect to the RGB mean ridge line is shown as anhistogram. The solid line represents the distribution of the internal photometric errors (aGaussian with σ = 0 . σ = 0 .
02, simulating a metallicity spread of δ [Fe/H]=0.1dex. 23 –Fig. 7.— Optical CMD of the inner sample zoomed in the MS-turnoff region, for 30 ′′ ≤ r ≤ ′′ . Superimposed are isochrones from the Girardi’s database computed at the clustermetallicity ( Z = 0 . − .
91) and at different ages (see labels). The isochronethat best fits the observed CMD is the 13 Gyr one (solid line). The inset shows an enlargedview of the CMD. 24 –Fig. 8.— Age-metallicity relation for Galactic GCs at R GC < R GC > inner sample . The adopted BSS and HB selection boxes areshown. Solid circles highlight the selected BSSs, while empty triangles mark the known RRLyrae stars. 26 –Fig. 10.— Optical CMDs of NGC 5824 zoomed in the BSS region. The left panel shows theCMD of the inner sample not covered by the UV data; the CMD of the outer sample for r ≤ ′′ is shown in the right-hand panel. The adopted BSS and HB selection boxes areshown with solid lines. Solid circles mark the BSSs, while empty triangles mark the knownRR Lyrae stars. 27 –Fig. 11.— CMD of the outer sample at r > ′′ , used to estimate the contaminationof Galactic field stars to the BSS, HB and RGB population selections (the correspondingselection boxes are marked with solid lines). 28 –Fig. 12.— Cumulative radial distribution of the statistically decontaminated populations ofBSSs (solid line), HB stars (dashed line) and RGB stars (dotted line) as a function of theprojected distance from C grav . 29 –Fig. 13.— Radial distribution of the population ratios N BSS /N HB , N BSS /N RGB , N HB /N RGB (top, middle, and bottom panels, respectively) as a function of the radial distance from thecluster center, normalized to the core radius. 30 –Fig. 14.— Radial distribution of the double normalized ratios of BSSs (black dots) and thereferences stars (grey rectangles: HB stars in the top panel, RGB stars in the bottom one). 31 –Fig. 15.— Core relaxation time ( t rc ) normalized to the age of the Universe ( t H = 13 . r min /r c (from Figure 4 of F12). The dynamically young systems ( Family I )are plotted as lower limit arrows at r min = 0 .
1. The grey triangles mark the dynamically oldclusters (
Family III ), while the grey circles mark the intermediate dynamical-age clusters(