The C+N+O abundance of Omega Centauri giant stars: implications on the chemical enrichment scenario and the relative ages of different stellar populations
A. F. Marino, A. P. Milone, G. Piotto, S. Cassisi, F. D'Antona, J. Anderson, A. Aparicio, L. R. Bedin, A. Renzini, S. Villanova
TThe C+N+O abundance of ω Centauri giant stars: implicationson the chemical enrichment scenario and the relative ages ofdifferent stellar populations A. F. Marino , A. P. Milone , , G. Piotto , S. Cassisi , F. D’Antona , J. Anderson , A.Aparicio , , L. R. Bedin , A. Renzini , S. Villanova Max-Planck-Institut f¨ur Astrophysik, Karl-Schwarzschild-Str. 1, 85741 Garching bei M¨unchen, Germany [email protected] Instituto de Astrof´ısica de Canarias, E-38200 La Laguna, Tenerife, Canary Islands, Spain [email protected], [email protected] Departamento de Astrof´ısica, Universidad de La Laguna, E-38200 La Laguna, Tenerife, Canary Islands,Spain Dipartimento di Astronomia, Universit`a di Padova, Vicolo dell’Osservatorio 3, Padova, I-35122, Italy [email protected] INAF-Osservatorio Astronomico di Teramo, Via M. Maggini, 64100 Teramo, Italy [email protected] INAF-Osservatorio Astronomico di Roma, via Frascati 33, I-00040 Monteporzio, Italy [email protected] Space Telescope Science Institute, 3700 San Martin Drive, Baltimore, MD 21218, USA [email protected] INAF-Osservatorio Astronomico di Padova, Vicolo dellOsservatorio 5, 35122 Padova, Italy [email protected], [email protected] Departamento de Astronom´ıa, Universidad de Concepci´on, Casilla 160-C, Concepci´on, Chile a r X i v : . [ a s t r o - ph . S R ] N ov ABSTRACT
We present a chemical-composition analysis of 77 red-giant stars in OmegaCentauri. We have measured abundances for carbon and nitrogen, and combinedour results with abundances of O, Na, La, and Fe that we determined in ourprevious work. Our aim is to better understand the peculiar chemical-enrichmenthistory of this cluster, by studying how the total C+N+O content varies amongthe different-metallicity stellar groups, and among stars at different places alongthe Na-O anticorrelation. We find the (anti)correlations among the light elementsthat would be expected on theoretical ground for matter that has been nuclearlyprocessed via high-temperature proton captures. The overall [(C+N+O)/Fe]increases by ∼ ∼ − ∼ − Subject headings: globular clusters: individual (NGC 5139) —
1. Introduction
Omega Centauri ( ω Cen) is one of the most intriguing Globular Clusters (GCs) of theGalaxy. At odds with the majority of GCs, which are mono-metallic, it shows large star-to-star metallicity variations up to more than one dex (e.g. Norris, Freeman & Mighell 1996,Suntzeff & Kraft 1996). At the same time, it shares with the most mono-metallic GCsthe presence of a Na-O anticorrelation, which is present across almost the entire metallicityrange (Marino et al. 2011a, Johnson & Pilachowski 2010). These two facts suggest uniquecomplexity in the cluster’s star-formation history, and its chemical evolution.The complexity of ω Cen also manifests itself in its color-magnitude diagram (CMD)with the presence of multiple red-giant (RGBs), multiple sub-giant branches (SGBs, Lee [email protected] Based on data collected at the European Southern Observatory with the FLAMES/GIRAFFE spectro-graph. ω Cen RGB stars (Dupree et al. 2011).The correlation of the He line detection with [Fe/H], Al and Na supports the assumptionthat helium is enhanced in the bMS.Due to the observational scenario, more complex than in any other GC, it has often beensuggested (e.g. Bekki & Freeman 2003) that ω Cen may be the remnant of a now-dissolveddwarf galaxy, once similar of the Sagittarius dwarf (with its central globular cluster M54)now being thorn apart by the Galactic tidal field. In any case, a successful description of the ω Cen star-formation history should be able to explain both the Na-O anticorrelation at thevarious metallicities and the general rise in slow-process ( s -process) elements as a function ofmetallicity (Norris & Da Costa 1995; Johnson & Pilachowski 2010; Smith et al. 2000; Marinoet al. 2011a). The presence of the Na-O anticorrelation implies that ω Cen, similarly tomono-metallic GCs, has experienced enrichment from high-temperature H-burning processedmaterial. At the same time, an additional physical mechanism should be present to produce s -process elements. In the Sun, the s -process elements abundance is mainly due to twocomponents: the main component (attributed to low-mass AGB stars of 1.5-3M (cid:12) ; Busso,Gallino & Wasserburg 1999) and the weak component (attributed to massive stars; seeRaiteri et al. 1993, and references therein). At solar metallicity, this last component mainlyproduces the lighter s -process nuclei, at most up to Sr and Kr (Raiteri et al. 1993). Inthe case of ω Cen, if we assume that the s -elements are produced in the strong component,i.e. in a population of less massive asymptotic-giant-branch (AGB) stars (Busso et al. 1999,Ventura et al. 2009), we have to face a timescale discrepancy that makes hard to solve theentire puzzle (D’Antona et al. 2011). Indeed, these AGB have longer lifetimes than themore massive AGB stars, which are presumed responsible for the He enrichment and forNa-O anticorrelations. If, alternatively, the s -process elements are produced in the weakcomponent by massive stars that explode as Type II Supernovae, it remains to be seenwhether the models can produce elements as heavy as Ba and La in stellar environmentscharacterized by the chemical abundances observed in ω Cen.An important ingredient to understand the star-formation history of this poorly-understoodGC is the determination of the relative ages of its stellar populations. Numerous studiesbased on CMD analysis combined with metallicity distribution of turnoff (TO) and SGB 4 –stars have yielded conflicting results, suggesting age differences from less than 2 Gyr (Leeet al. 2005, Sollima et al. 2005, Calamida et al. 2009), up to 5 Gyrs (Villanova et al. 2007).Rapid-formation scenarios wherein the entire cluster could have formed within a few times10 years have recently been suggested by D’Antona et al. (2011) and Valcarce & Catelan(2011).Theoretical isochrones show that a variation of the total C+N+O abundance can havean effect on the GC ages obtained from CMD fitting. Therefore, the relative age-datingof the stellar populations hosted in the cluster would be severely affected by the pres-ence of C+N+O differences among the ω Cen sub-populations (Cassisi et al. 2008, Venturaet al. 2009, Pietrinferni et al. 2009, D’Antona et al. 2009).In an effort to shed light on these issues, in this paper we measure the overall C+N+Oabundance in 77 RGB stars of ω Cen, which span nearly its entire metallicity range. Thelayout of this paper is as follows: in Section 2 we describe the data analysis; results arepresented in Section 3, and their impact on the chemical enrichment scenario and the agemeasurements are discussed in Section 4 and Section 5, respectively; Section 6 is a summaryof our results.
2. Observations and data reduction
Our data-set consists in a sample of 77 RGB stars observed with the FLAMES/GIRAFFEHR4 setup (program: 082.D-0424A). We also have, from Marino et al. (2011a), additionalspectra for the same stars with different GIRAFFE setups, from which we have derivedabundances for Fe, Na, O, and n -capture element La (75% s -process in the solar system,Simmerer et al. 2004). In Marino et al. (2011a) we provide meaurements also for Ba abun-dances, but here we prefer to use only La as representative of n -capture elements becauseBa measurements are more uncertain, since the only analyzed Ba transition is a blend (seeMarino et al. 2011a). We refer the reader to this paper for a more detailed description of thesample and the data reduction. Here we analyse, for 77 stars of the Marino et al. (2011a)sample, chemical abundances for carbon and nitrogen.The HR4 setup covers the spectral range from ∼ ∼ R ∼ ∼ A ∆ − X Π) near 4314 and 4323 ˚A. The nitrogen abundance was derived from synthesis of the CN 5 –blue-system ( B Σ − X Σ) bandhead at ∼ eff ), gravities (log g ), metallicities, and microtur-bolences ( ξ t ) determined in Marino et al. (2011a). In computing the C abundance we usedthe previously determined O contents (Marino et al. 2011a), and for N, both observed C andO abundances needed to be employed. As an example, we show in Figure 1 the spectralsynthesis for the CH and CN bands, and the O line for star eff , log g , metallicity, and ξ t for several stars, and redetermining the abundances. The parameters were varied by∆T eff = ±
100 K, ∆log g = ± ξ t = ± ± ∼ ± ± ∼ ω Cen RGB stars were determined by Brown &Wallerstein (1993; 6 stars), Norris & Da Costa (1995, 40 stars), and Stanford, Da Costa, &Norris (2010; 33 stars). A proper comparison with these studies cannot be made becausethere are no stars in common, however we note that our C and N values span a rangesimilar to that of Brown & Wallerstein (1993) and Stanford et al. (2010). Possible systematicdifferences in C could not be excluded between our stars and those of Stanford et al., whilesystematically lower N abundances (of ∼ ±
3. Results
Figure 2 plots the abundances of C, N, and O (Tab. 1) measured in this paper andthe Na measurements from Marino et al. (2011a). We observe large star-to-star variationsin all of these elements, as has already observed in all the GCs studied to date. Carbonspans a range from [C/Fe] ∼ − ∼ +0.6. Most stars are C-depleted, but there is apopulation of stars with [C/Fe] > ∼ ∼ Upper panels : [N/Fe] and [O/Fe] vs. [C/Fe] (panels c, a), and [O/Fe] and [Na/Fe]vs. [N/Fe] (panels d, b). The dashed black line in panel (c) separates C-poor/N-rich starsfrom C-rich/N-poor.
Lower panels : [N/Fe] vs. [C/Fe] in metallicity bins defined in Marinoet al. (2011a). In each panel, the dashed black line is the same as in panel (c), the reddot-dashed line represents the ”best-fit” line of constant [C/Fe]+[N/Fe] (=0.80) for the mid-metallicity stars.discussed in Section 3.1. Such chemical patterns indicate the occurrence of high-temperatureH-burning involving the CNO and NeNa cycles, and are consistent with prediction in Ventura& D’Antona (2009).In the following sections we present the pattern of C, N, O abundances with [Fe/H] (Sec-tion 3.1) and with respect to the position of stars on the Na-O anticorrelation (Section 3.2)studied by Marino et al. (2011a). 8 –ID [C/Fe] [N/Fe] [O/Fe]246266 − − − − − − Before discussing in detail the C, N, O patterns, we recall that ω Cen exhibits a largespread in [Fe/H], ranging from ∼ − . ∼ − . ∼ − − − − ∼ − > − total [C+N+O] against [Fe/H], we find a correlationwith a Pearson coefficient of 0.71. This is shown in panel (d) for the abundance relative toiron and in panel (e) for the total abundance. The red points show the averages over bins in[Fe/H] defined above. The vertical error bars correspond to the error in the mean related tothe observed dispersion and the horizontal bars denote the boundaries of the [Fe/H] bins.The representative error bar shown in the lower right of panels (d) and (e) comes fromadding the individual C, N, and O errors in quadrature, and indicates that the error in thetotal [(C+N+O)/Fe] should be quite large ( ∼ − ≤ [Fe/H] < − ω Cen in thecontext of other clusters. M22 has stars with a range of metallicities (Marino et al. 2009),and these stars are shown as the green points, which happen to follow the ω Cen trendalmost perfectly. On the other hand, the mono-metallic clusters for which we have data(NGC 6397, NGC 6752, M4, and 47 Tuc) do not follow this trend and, in general, show nodiscernible trend of [(C+N+O)/Fe] with metallicity.
To investigate how the studied chemical abundances behave for stars occupying a dif-ferent location on the Na-O plane, we divided stars into two groups based on their locationalong the Na-O trend (Marino et al. 2011a). In normal mono-metallic clusters different gen-erations of stars can be segregated in this way, and the Na-O groupings tend to have differentC and N abundances, as can be seen in the case of M4 (Marino et al 2008, see their Fig10). The situation in ω Cen has an added level of complexity, since the large variation inmetallicity does not allow us to simply identify two (or more) populations via Na and O.Both Johnson & Pilachowski (2010) and Marino et al. (2011a) have shown that stars of both first and second generations are present across a large range of metallicities. Note that here 10 –we extend the same nomenclature used for mono-metallic GCs to ω Cen and name firstand second generation Na-poor/O-rich and Na-rich/O-poor stars respectively, whatever istheir [Fe/H]. Of course, in the case of ω Cen this designation should not be taken literally,as stars of both groups are present in a large range in metallicity. We will therefore firstfollow the chemical patterns for the elements affected by p -captures (C, N, O, Na), n -captureelements (La), and C+N+O for the first generation alone, and then will examine the secondgeneration.In Figure 4 we represent in different colors our selected first and second generationstars on the Na-O anticorrelation from Marino et al. (2011a). The Na-poor/O-rich firstgeneration (green open triangles) and the Na-rich/O-poor second generation stars (magentaopen squares), have been arbitrarily selected by the dashed line. Similarly to what done inFigure 2, stars have been represented in different panels depending on their [Fe/H] bin. Ineach panel, the grey crosses represent the entire sample. The so selected first and secondgeneration stars have been plotted in a O-La plane, as shown in Figure 5 (right panel). Itis worth noting that in first generation stars [O/Fe] abundances correlate with [La/Fe], withthe Pearson coefficient equal to 0.73. To quantify this correlation we have determined themean [O/Fe] abundances for three La intervals spanned by first generation stars, whose sizesare indicated by the horizontal dark-green lines in Figure 5. The mean [O/Fe] values forthe different Fe and La intervals are listed in Tab. 2, togheter with the rms and associatederrors. In first stellar generation, some hints for a correlation may be present also among[O/Fe] and [Fe/H], as shown in the middle panel of Figure 5. In this case the O rise is moredifficult to be claimed over observational errors and the mean [O/Fe] values in the selectedmetallicity bins (shown in dark-green in the middle panel of Figure 5) agree within a 3 σ values. [Fe/H] range [La/Fe] range < − . − . / − . − . / − . − . / − . − . / − . > − . < − . − . / . > . / Fe] 0.41 0.45 0.50 0.62 0.64 · · · ± · · · σ · · · Table 2: Oxygen mean content for first generation stars in the Fe bins of Marinoet al. (2011a), and for three intervals in [La/Fe].We suggest some caution with the O trend in first generation stars, as we cannot fullyexclude that it may be due to some unidentified systematics with iron. Johnson & Pila-chowski (2010) do not find a trend of [O/Fe] with [Fe/H], but claim instead to find weaktrends ( ∼ . α element ratios, such as [Ca/Fe] and [Si/Fe], which howeverare even weaker than our trend in [O/Fe]. Notice that such trend is detectable only among 11 –first generation stars, as in second generation stars oxygen has been depleted by p -capturereactions.As shown in the right panel of Figure 5, in first generation stars O increases in concertwith the total CNO. At odds with first generation stars, the second generation ones (rep-resented in magenta) do not appear to show any correlation with either La, Fe, and CNO,and occupy a spread region in all these abundance planes. Lanthanum abundance ratiosfollow a similar pattern for first and second generation stars as shown if Figure 6. Thisimplies that the material out of which second generation stars have formed was not exposedto neutron sources (no additional n -captures besides those experienced by first generationstars) but only to p -captures. If AGB stars were responsible for the p -capture processing,this must have taken place in stars experiencing negligible third dredge-up, otherwise thematerial would have been further enriched in s -process elements.In Figure 7 we summarize our results for C, N, O, and Na abundance ratios, that havebeen plotted as a function of [(C+N+O)/Fe], [Fe/H] and [La/Fe], for the first and secondgeneration stars. As well as O, carbon increases as a function of the total CNO and [La/Fe],suggesting that C and O evolve in a similar way (though their errors are correlated).At odds with C and O, neither N nor Na show any evidence for correlation with thetotal CNO abundance in first generation stars. However, the minimum values for these twoelements may increase slightly with both Fe and La. N and Na display a much strongertrend with either the overall CNO, iron, and lanthanum for second generation stars, i.e. Nand Na are higher for stars with higher CNO/Fe/La.All the observed abundance pattern suggests that C and O from one side, and N andNa form the other, have undergone a similar processing in the evolution of the cluster. Inaddition, the first and second generation, selected on the basis of their Na and O content,appear to show well defined individual chemical trends along the overall observed range inmetallicity.Finally, the CNO increase among first generation stars is driven by the rise of O (andpartly also of C), rather than N. As C and O increase with metallicity, more N could beproduced by CNO cycling, and so in O-poor/Na-rich second generation stars both [N/Fe]and [Na/Fe] increase with metallicity. Figure 8 shows CNO vs. Fe separately for first andsecond generation stars, demonstrating the increase of CNO with Fe is common to bothgenerations. The small offset ( < ∼ . p -captures), but we cannotfully exclude a slight systematic underestimate of N among second generation stars. 12 –
4. The formation and chemical enrichment of ω Cen
The interpretation of the photometric and spectroscopic evidences on the stellar popu-lations of ω Cen in terms of its formation and evolution is posing formidable difficulties. Inthis Section we try to extract as much as possible from what the data themselves apparentlydemand.Although there is photometric evidence for at least six different stellar populations in ω Cen (e.g., Bellini et al. 2010), we have distinguished a first and a second generation asmade of those stars that do not or do show evidence for additional p -capture processing athigh temperature, respectively. Thus, the first generation may represent a less complex caseto interpret in terms of chemical evolution, with the second generation possibly arising fromthe AGB ejecta of the former one, as generally entertained. We first explore a scenario in which the progenitor system to ω Cen has evolved as asingle entity, in which the ISM was chemically homogeneous at each time. Thus, one canenvisage the first generation to be the result of five or six successive bursts of star forma-tion, with stars in each successive episodes being progressively enriched by nucleosynthesisproducts from core-collapse supernovae from the previous bursts. By construction, the firstgeneration is made of stars which formed out of an ISM that was not yet polluted by p -capture products by AGB stars. This requirement sets an upper limit of 30-40 Myr betweenthe first and the last burst, as later AGB stars would come at play.Now, before the second generation starts to form most residual gas must be ejectedfrom the system, thus allowing second generation stars being formed out of almost pureAGB ejecta. As emerging from all previous figures, second generation stars span the fullmetallicity range of ω Cen, which suggests that each episode of first generation formationhas produced its own specific contribution to the second generation. Thus, after the five orsix initial bursts of star formation, and the ejection of the residual gas, a similar series ofbursts should have taken place out of gas from AGB ejecta being accumulated at the bottomof the potential well, via a series of cooling flows.Also for this second generation series of bursts one can set a time limit to its duration,which comes from the necessity of avoiding the contamination by C/O > × yrs after the very first burst. Therefore,in this scenario all star formation episodes were confined to within ∼ × yrs. Of course,such a scenario with its complex series of events may require yet additional complications 13 –to explain the presence of the Na-O anticorrelation among stars within different individualstar formation episode of the second generation (i.e., within each [Fe/H] group). Even morecomplications may be required if the Na-O anticorrelation is due to partial mixing of AGBejecta with pristine material as often invoked (e.g., D’Ercole et al. 2011, and referencestherein), though the origin of the diluting material remains to be understood. Five or sixsuccessive dilutions with pristine material would ask for an extremely contrived scenario.Alternatively, common-envelope ejecta from intermediate-mass binary stars may offer a lessimplausible option (Vanbeveren, Mennekens & De Greve 2011) but it remains to be seen ifbinaries can provide enough diluting material.Moreover, as well known the original first stellar generation must have been much moremassive than at present, if enough second generation stars formed from its AGB ejecta (e.g.,Bekki & Norris 2006; Renzini 2008). This led to the hypothesis of a dwarf galaxy precursor ofwhich ω Cen would be the stripped, remnant nucleus (e.g., Bekki & Freeman 2003; D’Ercoleet al. 2010). In summary, in this scenario formation proceeds through a rapid series of starformation episodes (bursts), separated by just few Myrs in time, each followed by its owncooling flow made of AGB ejecta to form the second generation stars in a series of secondarybursts. Tidal stripping will then remove most of the original stellar mass of the system,leaving ω Cen as the bare nucleus of the original dwarf galaxy.Alternatively, the several star formation episodes, rather than being sequential in timemay have been separated in space, with each lump of matter having evolved (quasi)independentlyfrom the others, before merging together to make the cluster we observe today. In the orig-inal dwarf galaxy several high-density sites of star formation were active, each forming aproto globular cluster of first generation stars, with each of them later feeding the formationof its own second generation stars via a cooling flow of AGB-ejecta material. In turn, thedensest part of the individual proto-clusters would merge together by dynamical friction andcoalesce, thus forming a massive nucleus, to become ω Cen after all the rest of the dwarfgalaxy is tidally stripped.Difficult to say whether the former, monolithic scenario is more (or less) contrived thanthis merger one in which five or six proto-clusters evolve separately before coalescing together.Yet, the stellar population content of this most massive globular cluster in the Galaxy is socomplex that we suspect no simple, straightforward model can be found for its formation.Eventually we must admit that we are still far from understanding this most puzzling cluster. 14 –
The puzzle becomes even more difficult to compose when considering the specific chem-ical patterns exhibited by the stars in the cluster. Here we limit the discussion to the CNOelements.In Figure 8 [(C+N+O)/Fe] appears to increase with [Fe/H] among first generation stars,a trend that is dominated by the increase of [O/Fe] since oxygen is the most abundant of thethree elements. This is also indicated by the tight correlation between [(C+N+O)/Fe] and[O/Fe] seen in the right panel of Figure 6. The possible [O/Fe] increase with [Fe/H] amongfirst generation stars, is contrary to the decreasing trend shown in all other environments.In fact, [O/Fe] decreases in lockstep with [Fe/H] in the Galactic thin disk, thick disk, bulgeand halo (e.g., Bensby, Feltzing & Lundstr¨om 2004; Zoccali et al. 2006; Wheeler, Sneden& Truran 1989), a trend that is generally interpreted as due to a prompt enrichment of α -element by core collapse supernovae, followed by a slower iron enrichment by Type Iasupernovae (e.g., Matteucci & Greggio 1986). Moreover, values of [O/Fe] in excess of 0.5are common among first generation stars, whereas they are extremely rare in any otherenvironment. Therefore, also the normal population in ω Cen, the one not affected byproton captures, shows a unique chemical evolution history, not paralleled in any otherknown environment.Thus, this classical scheme of chemical evolution may not apply to ω Cen’s first gener-ation, as we have that oxygen keeps increasing faster than iron, a clear evidence contrary toType Ia supernovae playing a role in its enrichment. If so, one has to resort only on corecollapse supernovae in trying to explain the [(C+N+O)/Fe] trend with [Fe/H]. This is notan easy task.The existence of a broad range of metallicities among the first generation stars arguesfor enrichment and star formation having proceeded together in the progenitor object of ω Cen, though quite possibly as a series of successive bursts , as discussed in the previoussubsection and indicated by the distinct photometric sequences present in this cluster (e.g.,Bellini et al. 2010) paralleled by a multimodal distribution of [Fe/H] (e.g., Marino et al.2011a).In the scenario in which star formation proceeds with a series of bursts, the first succes-sive episode would have experienced an enrichment due only to the most massive supernovaefrom the first episode, and the following ones an enrichment due to a broader and broaderrange of supernova masses. Thus a trend in [O/Fe] with [Fe/H] could be explained if the ratioof oxygen-to-iron yields from core collapse supernovae were a decreasing function of stellarmass, i.e., if low mass supernovae were to make more oxygen relative to iron, compared to 15 –high mass supernovae.Theoretical yields of oxygen and iron as a function of the initial mass of the supernovaprogenitor formally do not support this possibility (see e.g. Woosley & Weaver 1995; Limongi& Chieffi 2006; Kobayashi et al. 2006), but are sufficiently uncertain to leave this option a priori relatively viable. The main uncertainty comes from the iron yield as a function ofstellar mass being poorly constrained by supernova models, mainly due to the difficulty inlocating the mass cut between ejecta and compact remnant that typically lies just within theiron layer of the pre-supernova structure. Nevertheless, this scheme would require lower massstars to dominate the oxygen production, whereas all supernova models agree in predictingan increase of the oxygen yield with stellar mass, at least up to ∼ M (cid:12) , and in most modelseven beyond.To explain the anomalous trend of [O/Fe] with [Fe/H] one may be tempted to appeal tothe last resort of desperate situations, and invoke different IMFs for each successive episodeof star formation, a very arbitrary and contrived scenario indeed. Moreover, note that thistrend of increasing [O/Fe] with increase [Fe/H] does not extend to the highest metallicitygroup in ω Cen, corresponding to the anomalous
RGB-a and MS-a sequences in the CMD(cf. old ref, Bellini et al. 2010), but here we have ascribed this whole metallicity group tothe p -capture processed second generation.In summary, even limiting ourselves to the normal , non p -capture processed, firstgeneration, we have to admit our failure in providing an explanation for the observed[(C+N+O)/Fe] and [O/Fe] trends with [O/Fe].
5. CNO and age of ω Cen populations
The C+N+O abundance has a strong impact on the location of isochrones at the TO andSGB level, and therefore on age determinations. In order to estimate the effects of the largevariation of the overall CNO among ω Cen’s stellar populations on the age measurement,we have calculated stellar models by including a variety of assumptions on the C+N+Oabundance and spanning a metallicity range suitable for the ω Cen stellar populations.The adopted physical inputs and numerical assumptions are the same as in Pietrinferniet al. (2009), which presents evolutionary computations accounting for a CNO enhancementof a factor of ∼ α − enhanced mixture. A fully consistent set of α -enhanced isochrones with no CNO enhancement for various iron abundances (Pietrinferniet al. 2006), allows us to compare, in a fully homogeneous theoretical framework, CNO-enhanced and not-CNO-enhanced isochrones. 16 –As expected, we find that, for a fixed TO and SGB brightness, the CNO-enhancedisochrones provide younger ages than isochrones corresponding to a canonical α -enhancedmixture (Pietrinferni et al. 2006): as a rule-of-thumb we found the ∂ age /∂ [CNO] ∼ − . − ,regardless of the iron content. By accounting for the observed [CNO/Fe] abundance in ω Cen,we have estimated the difference between the age one would measure for each populationby adopting appropriate CNO-enhanced isochrones and that obtained by using ‘normally’ α − enhanced isochrones, i.e. neglecting the measure CNO enhancement. The overall effect isto decrease the age estimate for the higher metallicity stars. The age decrease ranges from ∼ ∼ − ∼ ∼ − (cid:15) (C+N+O) ∆age(Gyr) − < [Fe/H] < − ± ± − ± − < [Fe/H] < − ± ± − ± − < [Fe/H] < − ± ± − ± − < [Fe/H] < − ± ± − ± − < [Fe/H] < − ± ± − ± ω Cen [Fe/H] groups.This result confirms that the measurement of relative ages of stellar populations in GCsrequires accurate measurements of the overall C+N+O abundance. Indeed, it is well knownthat the double SGBs detected in NGC 1851, M22, 47 Tucanae (Milone et al. 2008, 2011,Anderson et al. 2009, Piotto 2009), can be interpreted in quite different ways, depending onthe CNO content of the single populations. In particular, the faint SGB can be associatedeither with a stellar population significantly older than the stars in the bright SGB (up to1 Gyr), or to a second generation, almost coeval with the brighter SGB (age differencesof 100-200 Myr), but with enhanced CNO (Cassisi et al. 2008, Ventura et al. 2009). Thisscenario has received further support by recent findings in M22 by Marino et al. (2009,2011b): this GC hosts two stellar groups with different s -element and CNO abundances,and, in addition, the two groups of CNO-rich and CNO-poor stars have also different ironabundances, in close analogy with ω Cen.In NGC 1851, results on possible internal variations of C+N+O are contradicting (Yonget al. 2009, Villanova et al. 2010). However, theoretical models, properly accounting forphotometric signatures of the chemical peculiarities observed in the cluster sub-populationsas provided by Sbordone et al. (2011), suggest that in NGC 1851 only a C+N+O-enhancedsecond generation can satisfy all the observational constraints given by CMDs in variousphotometric bands (Milone et al. 2008, Han et al. 2009). 17 –Many studies have attempted to reconstruct the evolutionary history of ω Cen, by de-termining relative ages among the different stellar sub-groups, but none has been able toaccount for variations in the overall C+N+O abundance. Considering the complexity of itsSGB morphology, and of the metallicity distribution, and different assumptions regarding Heenhancement among ω Cen populations, the relative ages for the various stellar populationsas measured in different studies are quite different, ranging from a small or null age disper-sion, as proposed by Sollima et al. (2005), to an age spread of the order of 2-3 Gyr (Hilkeret al. 2004, Stanford et al. 2006) to 4 or more Gyr (Hughes & Wallerstein 2000, Hilker &Richtler 2000, Villanova et al. 2007).It is beyond the scope of this paper to evaluate in detail how C+N+O abundances changethe age-dating obtained in the multiple cited works. However, we can identify some generaltrends. By using the same α -enhanced isochrones adopted here, Villanova et al. (2007) founda large age difference ( ∼ ω Cen populations. Specifically they identify fourgroups of stars: (1) an old metal-poor group ([Fe/H] ∼ − ∼ − ∼ − ∼− ∼ ∼ ∼ − ∼ ω Cen has experienced aprolonged star-formation of ∼ ≥ ω Cen up to ∼ ω Cen should beconfined to within ∼ ∼ − ∼ ∼ ω Cen stellar populations formed in a short time period.As mentioned above, the goal of this paper is not do a detailed analysis of previousage determinations. As discussed before there are contradicting results in recent literature.Our aim is simply to show that there is significant variation in the C+N+O content among ω Cen’s populations that may easily change these datations. In addition to that caveat, itis worth noting that He enhancement is likely present in the O-poor, metal-rich stars. Thisoccurrence should be taken into account when measuring the relative ages via isochronesfitting among the various sub-populations, since He affects the SGB shape (but not theluminosity level).
6. Conclusion
In this paper, we have presented C, N, O abundances for 77 ω Cen RGB stars in themetallicity range ∼ − < [Fe/H] < − • a correlation between the total CNO and the iron abundance, with the most metal-rich population being enhanced by ∼ • O and C grows with [Fe/H] and [La/Fe] in the O-rich/Na-poor stars. • stars selected on the basis of their position on the Na-O plane, show defined chemicalpatterns in their light elements C, N, O, and Na as a function CNO, Fe, and La, thatallow us to distinguish between a first and a second generation of stars, each possiblyresulting from a series of separate bursts of star formation. • [La/Fe] correlates tightly with [Fe/H], following almost precisely the same trend re-gardless of the O/Na abundances.In an attempt to make sense of this observational trends we have explored two (specu-lative) scenarios for the formation and evolution of this most puzzling object. In one optionthe system has evolved monolithically , i.e., remaining chemically homogeneous at each time.Thus, within less than ∼ −
40 Myrs a series of bursts of star formation each enriched iniron and α elements by core-collapse supernovae from the previous burst(s) would have been 19 –followed each by a secondary burst of star formation originated from material ejected byAGB stars, heavily processed by p -captures. The formation of the whole stellar populationinhabiting the cluster today would have taken less than ∼ × yrs. Alternatively, ratherthan sequential in time, such primary and secondary bursts of star formation would havetaken place separately in space, before their products could merge at the bottom of thepotential well. In both cases, secondary bursts would be fed via cooling flows made of AGBejects, leading to more centrally concentrated second generations , which indeed appears tobe so in this (and other) clusters (Sollima et al. 2005, Bellini et al. 2009, Johnson & Pila-chowski 2010). We admit that both scenarios require a series of very contrived and ad hoc assumptions, yet the extraordinary complexity of ω Cen may not admit simple solutions.The CNO abundance affects the determination of the relative ages of cluster subpopu-lations via isochrone fitting of the turnoff SGB region. In the light of our results, we havediscussed this issue in the case of ω Cen by comparing isochrones with standard and withenhanced CNO, with the latter ones giving younger ages for the same turnoff luminosity.Although the determination of relative ages of ω Cen sub-populations is beyond the aims ofour study, we argue that a trend in CNO/Fe can help reducing large age spread among thevarious sub-populations, as found by some studies in the literature.We thank the anonymous referee whose suggestions have significantly improved thiswork. APM, GP, SC and AA are founded by the Ministry of Science and Technology of theKingdom of Spain (grant AYA 2010-16717). APM and AA are also founded by the Institutode Astrofsica de Canarias (grant P3-94).
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This preprint was prepared with the AAS L A TEX macros v5.2.
23 –Fig. 3.—
Upper panels : [C/Fe], [N/Fe], and [O/Fe] vs. [Fe/H].
Middle panels :[(C+N+O)/Fe] and log (cid:15) (C+N+O) vs. [Fe/H]. Stars with upper limits for O or N abun-dances are represented as grey crosses. Red points represent the average C+N+O contentin the metallicity intervals spanned by the horizontal red bars. Vertical bars are the errorsassociated with the mean values.
Lower panels : Average [(C+N+O)/Fe] vs. iron for theGCs quoted in the inset compared with the ω Cen CNO mean values (red points) of pan-els d and e . Data for the mono-metallicity GCs are from Ivans et al. (1999) for M4, andCarretta et al. (2005) for 47 Tuc, NGC 6752, NGC 6397. Data for M22 are from Marinoet al. (2011b). 24 –Fig. 4.— Our adopted division of stars in O-rich/Na-poor ( first generation ) and O-poor/Na-rich ( second generation ) has been represented in the Na-O plane for each metallicity bin, asquoted in the insets. The dashed line separates the selected first generation stars representedby green triangles, from the selected second generation stars represented by magenta squares.Gray crosses in each panel represent the entire sample analysed by Marino et al. (2011a).Fig. 5.— O abundance ratios as a function of [La/Fe], [Fe/H] and [(C+N+O)/Fe]. Symbolsare as in Figure 4. The dark green vertical error bars represent the error associated withthe mean [O/Fe] abundance in different intervals in [La/Fe] and [Fe/H] delimited by thehorizontal line. 25 –Fig. 6.— [La/Fe] as a function of [Fe/H]. Symbols are as in Figure 4.Fig. 7.— C, N, O, and Na abundance ratios as a function of [(C+N+O)/Fe], [Fe/H] and[La/Fe]. Symbols are as in Figure 4. 26 –Fig. 8.— CNO relative to Fe and log (cid:15) (CNO) as a function of [Fe/H]. Symbols are as inFigure 4. Crosses mark stars for which we have only an upper limit for N or O.Fig. 9.— Age difference for each ωω