PROTON CAPTURE CHAINS IN GLOBULAR CLUSTER STARS. III. ABUNDANCES OF GIANTS IN THE SECOND-PARAMETER GLOBULAR CLUSTER NGC 70061

Transcription

1 THE ASTRONOMICAL JOURNAL, 115:15È1515, 1998 April ( The American Astronomical Society. All rights reserved. Printed in U.S.A. PROTON CAPTURE CHAINS IN GLOBULAR CLUSTER STARS. III. ABUNDANCES OF GIANTS IN THE SECOND-PARAMETER GLOBULAR CLUSTER NGC 761 ROBERT P. KRAFT,2 CHRISTOPHER SNEDEN,3 GRAEME H. SMITH,2 MATTHEW D. SHETRONE,4 AND JON FULBRIGHT2 Received 1997 July 24; revised 1997 December 18 ABSTRACT High-resolution spectra have been obtained of six red giants in the globular cluster NGC 76 using the HIRES instrument on the Keck I Telescope. The [iron-peak/fe] and [a-nuclei/fe] elemental ratios are similar to those found in other halo globular clusters. SigniÐcant, but modest, star-to-star abundance di erences are found for the elements O, Na, and Al, with the [O/Fe] abundance ratio being anticorrelated with both [Na/Fe] and [Al/Fe]. These anticorrelations indicate that O ] N processed material, within which Ne has also been converted to Na and Mg to Al by proton addition reactions, has been brought to the surface of the cluster giants via a deep-mixing mechanism that cycles material through the O ] N burning zone just outside the main hydrogen-burning shell. Despite the signiðcant [Al/Fe] di erences among the NGC 76 giants, there is very little accompanying di erence in [Mg/Fe]. This result, however, is not inconsistent with a deep-mixing origin for the [Al/Fe] enhancements; because Mg is initially much more abundant than Al, conversion of only a small fraction of Mg into Al can still produce a large percentage increase in the Al abundance. Plots of [Na/Fe] versus [O/Fe] and [Al/Fe] versus [Mg/Fe] indicate that the amount of protoncaptureèprocessed material that has been mixed into the envelopes of the NGC 76 giants is less extreme than within some giants of the globular clusters M13 and M15. In particular, unlike the NGC 76 giants, most giants in M13 and some in M15 exhibit signiðcantly reduced [O/Fe] and [Mg/Fe] ratios and very high [Al/Fe]. This di erence in the [Mg/Fe] behavior may be related to di erences in horizontal-branch morphology between M13 and NGC 76. Within some M13 giants, it is possible that deep mixing has accessed regions in which not only has Mg been converted to Al but signiðcant amounts of helium have also been produced from hydrogen burning. Deep mixing in such stars may not only signiðcantly reduce the surface [Mg/Fe] ratio but may also increase the [He/H] ratio by amounts that will a ect the zero-age horizontal-branch (HB) position that a giant will eventually occupy. This could explain the extension of the M13 HB to very blue colors. By contrast, giants in NGC 76 experience modest amounts of interior mixing that may be deep enough to detectably alter the surface [Al/Fe] ratio, but not the [Mg/Fe] and [He/H] ratios. These stars will have a relatively low [He/H] ratio and so will adopt relatively red positions when they evolve onto the horizontal branch; the HB of NGC 76 does indeed comprise a large percentage of stars redder than the RR Lyrae gap. In this scenario, deep mixing within cluster red giants may be one of the second parameters ÏÏ that a ects horizontal-branch morphology. Key words: globular clusters: individual (NGC 76) È stars: abundances È stars: evolution 1. INTRODUCTION The second-parameter problem ÏÏ arises from the observation that the morphology of the horizontal branch in globular cluster color-magnitude diagrams (CMDs) can vary among clusters having the same Fe-peak element metallicity (Sandage & Wildey 1967; van den Bergh 1967). The cluster pair M3 and M13 provides one of the best examples: they have Fe-peak element metallicities ÈÈÈÈÈÈÈÈÈÈÈÈÈÈÈ 1 Based on observations obtained with the Keck I Telescope of the W. M. Keck Observatory, which is operated by the California Association for Research in Astronomy (CARA), Inc., on behalf of the University of California and the California Institute of Technology. 2 UCO/Lick Observatory, Board of Studies in Astronomy and Astrophysics, University of California, Santa Cruz, CA 9564; graeme=ucolick.org, kraft=ucolick.org, jfulb=ucolick.org. 3 Department of Astronomy and McDonald Observatory, University of Texas at Austin, Austin, TX 78712; chris=verdi.as.utexas.edu. 4 European Southern Observatory, Alonso de Cordova 317, Casilla 191, Santiago 19, Chile; mshetron=eso.org. ([Fe/H] D [1.6)5 that di er by less than.1 dex (e.g., Kraft et al. 1993), and although the red giant branch (RGB) sequences in their CMDs are quite similar, their horizontalbranch (HB) morphologies are distinctly di erent. The horizontal branch of M3 is populated more or less uniformly from B[V D [.2 all the way to the red giant branch and contains nearly 2 RR Lyrae variables. By contrast, the horizontal branch of M13 has no stars redward of the RR Lyrae gap and, at most, two or three RR Lyrae stars; all nonvariable horizontal-branch stars in M13 are hotter than the RR Lyrae instability strip. Another factor must be at work besides Fe-peak element metallicity in determining horizontal-branch morphology; this unknown factor has come to be referred to as the second parameter.ïï Over the years, various explanations of this e ect have ÈÈÈÈÈÈÈÈÈÈÈÈÈÈÈ 5 We adopt the usual spectroscopic notations that [A/B] 4 log (N /N ) [ log (N /N ) and that log v(a) 4 log (N /N ) ] 12., for elements A B star A and B. A Also, B _ metallicity is arbitrarily deðned A H as the stellar [Fe/H] value. 15

2 PROTON CAPTURE CHAINS. III. 151 been o ered, including variations in age, He abundance, the [O/Fe] ratio, and mass loss during the Ðrst-ascent red giant branch phase of stellar evolution. Cluster age has often been discussed as a likely factor that controls the second parameter (Searle & Zinn 1978; Preston, Shectman, & Beers 1991; Lee, Demarque, & Zinn 1988, 1994; Lee 1991, 1992; Zinn 1993). Studies of the structure of red giant stars that are about to ignite the He core suggest that those with more massive, and therefore younger, main-sequence progenitors are more likely to take up residence on the redder part of the HB. If age di erences between globular clusters do indeed provide the correct explanation of the second parameter, then important consequences follow for the history of the Galaxy, since the clusters having excessively ÏÏ red HBs for their metallicity are often preferentially found in the outer Galactic halo. This would lend credence to the notion that the age of the halo increases with decreasing Galactocentric distance. Whether age alone is the second parameter has not been determined. The outer halo clusters Pal 12, Rup 16, Arp 2, and Ter 7 are almost certainly about 3È4 Gyr younger than typical Galactic globular clusters (Gratton & Ortolani 1988; Stetson et al. 1989; Buonanno et al. 199; Da Costa, Armandro, & Norris 1992). However, Richer et al. (1996) have recently challenged the assertion that age is a unique second parameter for other Galactic globular clusters. They have considered 35 clusters for which there are wellestablished deep CCD CMD arrays and have concluded that there exist few demonstrable age di erences exceeding 1 Gyr among clusters belonging to the halo, i.e., among clusters that have [Fe/H] \ [1 (the more metal-rich clusters of the thick disk probably are somewhat younger). The four anomalously young outer halo clusters can be identi- Ðed as having kinematics associated with the Magellanic Stream and, thus, are objects that may be in the process of accretion to the Galaxy. Richer et al. conclude that there is no compelling evidence for an age gradient among the original halo globular clusters, or for age being the unique second parameter. Stetson, VandenBerg, & Bolte (1996) reached a similar conclusion from a detailed intercomparison of color-magnitude diagrams of several clusters thought to evince the second-parameter phenomenon. If age is not responsible for the second-parameter e ect in every case, what else is plausible as an explanation? Highresolution spectroscopic studies have revealed an interesting distinction between the blue HB ÏÏ cluster M13 (Kraft et al. 1997) and the normal HB ÏÏ cluster M3 (Kraft et al. 1993, 1995). In both clusters, the spectra of Ðrst-ascent red giants indicate that material in their outer envelopes has undergone advanced proton-capture synthesis, in which sodium has been freshly synthesized at the expense of neon in regions where carbon and oxygen are being converted into nitrogen (Denisenkov & Denisenkova 199; Langer, Ho man, & Sneden 1993; Langer & Ho man 1995; Denissenkov et al. 1997a; Denissenkov, Weiss, & Wagenhuber 1997b). However, the synthesis progresses much further among M13, as opposed to M3, giants. Thus, in a sample of 1 of the brightest M3 giants, all but one star exhibited an [O/Fe] ratio above D., whereas in a complete sample of the 1 brightest M13 giants, all but three had [O/ Fe] \ [.2, and of these, three were as low as approximately [.8 (Kraft et al. 1993, 1995, 1997). The giants of M13 also exhibit striking correlations of nitrogen with sodium (Lehnert, Bell, & Cohen 1991) and magnesium with oxygen, as well as anticorrelations of aluminum with magnesium (Shetrone 1996a, 1996b; Kraft et al. 1997), the latter in particular implying that the envelope material of M13 giants has undergone advanced proton-capture synthesis in which Mg isotopes have been transmuted into 27Al (Langer & Ho man 1995; Cavallo, Sweigart, & Bell 1996, 1998; Langer, Ho man, & Zaidins 1997; Shetrone 1996b; Kraft et al. 1997). Although only three M3 giants have known Mg abundances and only one has a measured Al abundance (Cohen 1978), these abundances are consistent with a picture in which M3 giants, in contrast to M13 giants, have undergone considerably less proton-capture synthesis of Mg into Al, in concert with the behavior of O, N, Ne, and Na. Obviously, more observations of Mg and Al among M3 giants are needed. Why M13 and M3 di er in the distribution of these lightelement abundances is a subject of debate (cf. reviews by Suntze 1993; Briley et al. 1994; Kraft 1994). In one picture, M13 giants actually began with more of these extreme abundance anomalies than did M3 stars (the primordial ÏÏ scenario), in which case the proton-capture synthesis took place in two distinctively di erent nucleosynthetic environments prior to the formation of the stars we presently see. Alternatively, the hydrogen-burning shell, known to exist in all low-mass red giants, could be the site of the protoncapture synthesis, in which case the abundance di erences indicate that M13 giants mix their envelopes more iciently and perhaps more deeply into the shell than do giants of M3 (the evolutionary ÏÏ scenario). In the case of M13, there is direct observational evidence in favor of the second mechanism. Thus, oxygen becomes progressively more severely depleted as stars approach the red giant tip (Kraft et al. 1993), and in addition the Na abundance distribution, derived from a large sample of M13 giants, shifts sharply higher as Ðrst-ascent giants evolve above M B V [1.7 (log g B ]1) (Pilachowski et al. 1996). Consequently, we assume that the abundance anomalies are brought about by mixing of envelope material through the hydrogen-burning shell, and that this process proceeds more e ectively among M13 giants than among M3 giants. The possibility that stellar angular momentum may be the driving mechanism for this di erence in mixing (Sweigart & Mengel 1979; Sweigart 1997) is supported by the discovery (Peterson 1983; Peterson, Rood, & Crocker 1995) that some HB stars in M13 have rotational velocities up to v sin i D 4 km s~1 whereas, in M3, HB rotational velocities are all below the detection limit of 2 km s~1. Exactly which elements are subject to depletion or enhancement can be anticipated by examining the wellknown red giant interior models of Sweigart & Mengel (1979). It is assumed that some red giants have within their interiors the means for transporting material across the radiative zone known to exist between the hydrogenburning shell, where proton addition nucleosynthesis is occurring, and the base of the convective envelope, which does not quite reach down to the shell (Smith & Tout 1992; Denissenkov & Weiss 1996). Shallow mixing, into a precursor region well ahead of the hydrogen-burning shell, results in the production of 13C from proton captures on 12C and leads at the surface to an approach of the 12C/13C ratio to equilibrium. Progressively deeper mixing leads to the conversion of C to N, Ne to Na, and, eventually, O to N (see Denisenkov & Denisenkova 199; Langer & Ho man 1995). Recent models incorporating improved reaction rates (Cavallo et al. 1996; Langer et al. 1997) show that very deep mixing produces fresh 27Al from 24Mg, in accordance with

3 152 KRAFT ET AL. Vol. 115 TABLE 1 CLUSTER METALLICITY AND HB MORPHOLOGY Cluster [Fe/H]a (B[R)/(B]V ]R)b HB Type M3... [ Intermediate M13... [ Blue M1... [ Blue NGC [1.59 [.11 Red azinn &West1984. b Lee et al the observational results of Shetrone (1996b). In particular, the Langer et al. (1997) model shows that whenever there is observational evidence, as in M13, for severe deep mixing of the kind that seriously depletes O and signiðcantly enhances Al at the expense of Mg, the same deep mixing may bring fresh He to the surface at the expense of H, thereby increasing the He/H ratio in the envelope (Langer & Ho man 1995). These authors further suggested that such an increase might also serve to increase the e ective temperatures of their HB descendants and thus drive the second-parameter e ect. This idea was further elaborated by Sweigart (1997), who calculated models of red giants having di erent He/H ratios in their envelopes as these stars evolved through the He core Ñash. Using these models one can calculate, for any assumed precursor red giant massloss rate, the increase in e ective temperature of its HB descendant for any given increase in the helium content of the red giant envelope. For example, using the same massloss rate as would be appropriate for an unmixed star, one Ðnds that an increase of *Y \.12 would change the e ective temperature T of an HB descendant from 63 to 1, K. Thus, a star in the RR Lyrae domain would become a blue HB star. This result is consistent with earlier theoretical studies (Rood 1973; Dorman 1992) in which it was shown that HB morphology is very sensitive to He abundance. (In 6, we discuss the Sweigart models more extensively.) These results suggest that blue horizontal branches will be generated if the precursor red giants dredge up large amounts of CNO-processed material in which the He mass fraction has been increased. The diagnostics indicative of the e ect, as implied by the models, are severe depletion of O and, especially, conversion of Mg into Al. At least three other clusters with blue HBs are known to contain giants in which O is severely depleted and in which Al and Mg abundances are unusually high and low, respectively. These are NGC 6752 (Shetrone 1997), u Cen (Norris & Da Costa 1995a, 1995b; Zucker, Wallerstein, & Brown 1996), and M15 (Sneden et al. 1997). However, only NGC 6752 has a metallicity similar to M13 and M3 (Zinn & West 1984), whereas u Cen is on the average somewhat, and M15 very much more, metal-poor than M13 and M3, and thus the blueness ÏÏ of their HBs may reñect only the inñuence of the Ðrst parameter.ïï On the other hand, masses for very blue HB stars in M15 are difficult to account for (Moehler, Heber, & de Boer 1995) unless the He/H ratio were increased by an amount considered unacceptable by Moehler et al., but which would be plausible if the atmospheres of these stars had been enriched in He by amounts consistent with the extra mixing envisioned here. Furthermore, if some precursor red giants mix more than others (in response to di ering amounts of angular momentum?) the resulting differences in the He/H ratio would produce di erences in color on the HB. Thus it is possible that, in the same cluster, one might Ðnd that the He/H ratio on the average increases as one examines progressively bluer stars on the HB (Sweigart 1997). In summary, spectroscopic observations showing wide variations in O, Na, Mg, and Al abundances provide evidence of deep mixing within some cluster giants, especially those in M13, u Cen, NGC 6752, and M15, and lead to the suggestion that mixing-driven increases in envelope He abundance in these precursor red giants could be the cause of the very blue HBs seen at least in some clusters. Conversely, in a cluster showing little evidence of deep mixing among its giants, one might expect to Ðnd an HB that is considered too red for the cluster metallicity. In this paper, we test this suggestion by considering the archetypical red ÏÏ second-parameter cluster NGC 76. Whereas M13 shows a blue-dominated HB in comparison with M3, NGC 76 has a red-dominated HB, with most stars located redward of the RR Lyrae gap, this despite having a metallicity similar to that of M13 and M3. In Table 1, we summarize the HB morphology of M3, M1, M13, and NGC 76, using the morphological indicator (B[R)/ (B]V ]R) of Lee et al. (1994), in which B, V, and R refer to the number of blue HB stars, RR Lyrae variables, and red HB stars, respectively. We include M1 here because its metallicity is similar to that of M3, M13, and NGC 76 and it has a blue HB. It also has a signiðcant number of oxygen-depleted and sodium-enhanced stars, in which the degree of oxygen depletion is intermediate between that of M13 and M3 (Kraft et al. 1995). Unfortunately, nothing is known about the Mg and Al abundances in M1 giants. NGC 76 is not the only cluster to exhibit an excessively red HB for its metallicity. However, it is one of the most observationally accessible of such clusters. Lee et al. estimated HB morphological types for 83 Galactic globular clusters. Of this sample, only 1 clusters with metallicities [Fe/H] \ [1 have negative values of (B[R)/ (B]V ]R). Most of these lie in the outer Galactic halo, so that even the brightest giants have quite faint apparent magnitudes and are essentially unobservable at the required spectral resolution and signal-to-noise ratio (S/N), even with the Keck Telescopes. Two of the 1 also lie too far south to be accessible from Mauna Kea. NGC 76, by contrast, is reachable from Mauna Kea, both in terms of declination and the apparent magnitude of its red giant branch tip stars. Thus, of the eight candidate giants in the Sandage & Wildey (1967) photometric list that are brighter than V \ 16.5, we were able to observe six. In the following sections we describe the NGC 76 observations ( 2), the data analysis procedures ( 3), and the abundance results ( 4). A discussion of the abundance patterns found in NGC 76 and implications for the secondparameter e ect appear in 5. In Appendix A, we discuss the velocity dispersion and radial velocity of the cluster, both of which are somewhat unusual, and in Appendix B we examine the derivation of Al abundances in more detail. Caveats and concerns about the viability of the deep-mixing model as a means of enhancing 27Al at the expense of 24Mg (Shetrone 1996b) are discussed brieñy in OBSERVATIONS Our Ðrst observations of giants in NGC 76 were acquired during two nights in 1995 July using the HIRES

4 No. 4, 1998 PROTON CAPTURE CHAINS. III. 153 TABLE 2 OBSERVATIONAL DATA Exposure S/N Slit Star V a B[V a Date (UT) (minutes) (approximate) (arcsec2) V Jul ].86 I Jul ].86 II Jul ].86 II Aug ] 1.15 V Aug ].86 II Aug ] 1.15 a V and B[V are taken from Sandage & Wildey For long-period variable stars V19 and V54, the average values from that study are quoted here. B and V magnitudes vary within a total range of D.8 mag in V19 and D.6 mag in V54. echelle spectrograph (Vogt 1992; Vogt et al. 1994) of the Keck I Telescope. Three NGC 76 giants were observed on this run. The red-optimized collimator of HIRES was employed. A KV48 Ðlter was placed in the light path at the slit to block violet second-order light produced by the grating cross-disperser. The spectrograph slit for the 1995 observations was 7A ] A.86 in size, which yields a resolving power of R \ 45,. The detector used was a thinned back-sideèilluminated Tektronix 248 ] 248 pixel device binned 1 ] 1, each pixel being of dimension 24 km. The detector did not cover the full format of the HIRES echelle; consequently, it was positioned to ensure, Ðrst, that the [O I] lines at 63 and 6364 A fell on the chip and, second, that some lines of the interesting elements Na, Mg, Al were also obtained, e.g., Na I jj5682, 5688, Mg I j5711, and Al I jj6696, The orders recorded a range in wavelength from 54 to 76 A. During the 1995 run the NGC 76 giants were observed at air masses less than 1. A second set of observations were made over the course of three half-nights in 1996 August; spectra for a further three giants were obtained. During this run a wider slit (7A ] 1A.15) was used for observations of two stars, but otherwise the spectrograph setup was similar to that of the 1995 run, except for a slightly di erent location of the CCD relative to the HIRES format. The 1996 observations were typically continued to larger air masses than those made in 1995, with some exposures being initiated at air masses greater than 2.. V NGC 76 Sandage & Wildey 1967 probable members program stars B-V FIG. 1.È(V, B[V ) color-magnitude diagram of evolved stars in the globular cluster NGC 76 from the data of Sandage & Wildey (1967). This diagram shows only those stars deemed to be probable cluster members by Sandage & Wildey. Stars we observed with the Keck I Telescope are shown as larger Ðlled circles. NGC 76 is so distant that even its RGB stars are quite faint; therefore, even with the Keck I Telescope multiple exposures had to be obtained for each star. Accumulation of cosmic-ray events on the CCD detector limited individual exposures to less than an hour. Typically, three to six separate exposures of 18 s duration were obtained for each star. We also observed the rapidly rotating hot star 28 Vul (B5 IV), in order to record the telluric absorption lines that would be removed from the cluster star spectra in the reduction phase. Calibration exposures acquired on each night also included a ThAr arc lamp for use in obtaining an initial wavelength calibration, and a quartz lamp (through an NG3 Ðlter) used for Ñat-Ðeld corrections. The stars chosen for observation were selected from the color-magnitude study of Sandage & Wildey (1967); their star designations are used in this paper. Data for the observed stars are given in Table 2. In this and other tables, we list the NGC 76 stars in descending order of B[V. The V and B[V values are from the photographic photometry of Sandage & Wildey; also given is the date of observation, the total exposure time for each star, the approximate S/N near 635 A in the reduced spectra, and the spectrograph slit size. The observed giants are among the reddest and brightest in NGC 76. This is illustrated in Figure 1, which shows a (V, B[V ) color-magnitude diagram for evolved stars in the cluster based on the photometry of Sandage & Wildey (1967). Several stars in our sample are known to be long-period variables, and it is quite possible that others are variable as well. 3. DATA REDUCTION AND ANALYSIS Initial processing of the data was carried out using the IRAF software package. The CCD frames were corrected for both bias and Ñat-Ðeld e ects, and the individual orders were extracted. Further analysis of the data was performed using the SPECTRE code (Fitzpatrick & Sneden 1987) and involved continuum placement and normalization, cosmicray removal, a wavelength calibration using absorption lines within each order of a stellar spectrum, and removal of telluric absorption features. We measured equivalent widths (EWs) for all lines of interest by either of two techniquesè direct integration of the Ñux across an observed line proðle, or by adopting the equivalent width of a Gaussian proðle that was Ðtted to a line. These EWs are listed in Table 3. In order to calculate basic parameters for the NGC 76 red giants, values of the cluster distance modulus and reddening have to be adopted. We used the values in the tabulation by Djorgovski (1993) [(m [ M) \ 17.96, E(B[V ) \.5, A \.16], which in turn is based on the V photometric study by Buonanno et al. (1991). With these

5 TABLE 3 NGC 76 EQUIVALENT WIDTHS Line EP (A ) (ev) log gf V19 I-1 II-13 II-46 V54 II-18 [O I]: [ [ Na I: [ [ Al I: [ [ Si I: [ [ [ [ Ca I: [ [ [ [ [ [ Sc II: [ [ Ti I: [ [ V I: [ [ [ [ [ [ Fe I: [ [ [ [ [ [ [ [ [ [ [ [ [ [ [ [ [ [ Fe II: [ [ [ [ [ Ni I: [ [ [ [ Ba II: [ [ [ Eu II: ]

6 PROTON CAPTURE CHAINS. III. 155 data we converted the V and B[V photometry (Table 2) into photometric H-R diagram quantities (B[V ) and M, which are listed in columns (2) and (3) of Table 4. Also given V in column (4) of this table are estimates of M, computed with the aid of G. WortheyÏs (1994, private communication) bol bolometric corrections. The photometric quantities were then used to estimate initial model atmosphere parameters T and log g through two [(B[V ), M] ] (T, log g) mappings based on M3 and M13: (1) V Suntze Ïs (1981) transformation grid (his Fig. 15), and (2) our own previous spectroscopic studies (Kraft et al. 1992, 1993, 1997). The T, log g, and an initial assumed NGC 76 metallicity of [M/H] \[1.5 were input to the MARCS code (Gustafsson et al. 1975) to produce trial model atmospheres. Final model parameters were determined iteratively as in our past work by using the current version of the line analysis code MOOG (Sneden 1973) with the trial models, the measured EWs, and atomic parameters of the lines (identical to those of Kraft et al. 1997; see that paper for references and information on transition probabilities) to compute abundances for the lines given in Table 3. The model atmosphere parameter estimates were changed and new models computed until the derived abundances showed no signiðcant trend with excitation potential for Fe I lines, no signiðcant trend with EW for lines of any species, and good agreement between the average of Fe I and Fe II lines. In columns (5)È(8) of Table 4, the Ðnal values of T, v, and t log g are listed, and in columns (9)È(1) the Fe I and Fe II abundances derived with these models are given. The mean abundance di erence between these two species is only.7 dex; the average of these two abundances, given in the last column of the table, serves as our fundamental metallicity estimate for each star. In column (7) of Table 4 we list log g(pred), an estimate based on the value of M and an assumed stellar mass of.85 M. The predicted V ÏÏ and observed ÏÏ surface gravities are _ in good agreement, with S* log gt \].15 ^.6 (p \.14). If we omit II-46, the star with the largest deviation, then S* log gt \ ].9 ^.2 (p \.4). If we had assumed all stars to have su ered mass loss of.1 M by this evolutionary stage, the latter di erence would have _ been reduced to only ].4 dex. With the Ðnal model atmospheres, we used the EW matches to determine Ðnal abundances for most species. Exceptions were (as in our previous studies) (1) oxygen, for which we used synthetic spectrum Ðts to the observed [O I] lines at 63 and 6364 A ; (2) sodium, for which we combined synthetic spectrum calculations for the jj5682, 5688 lines with EW matches for the jj6154, 616 lines; (3) magnesium, for which we used synthetic spectrum Ðts to the line at 5711 A ; and (4) barium, for which we used observed/ synthetic EW matches but took proper account of the isotopic and hyperðne substructure in the jj5853, 6141, 6496 lines. All derived abundances for each star are given in the top section of Table 5, and the cluster mean abundances for each element are in the bottom part of this table. TABLE 4 FINAL MODEL PARAMETERS AND [Fe/H] VALUES log g [Fe/H] T v STAR (B[V ) M M (K) Spectroscopic Predicted (km t s~1) Fe I Fe II Average (1) (2) (3) V (4) bol (5) (6) (7) (8) (9) (1) (11) V : [2.62: [3.7: [1.62 [1.52 [1.57 I [23 [ [1.55 [15 [1.5 II [1.89 [ [1.55 [1.54 [1.54 II [1.74 [ [1.6 [1.53 [1.56 V : [2.27: [2.9: [1.65 [1.56 [1.6 II [1.8 [ [1.56 [1.5 [1.53 NOTE.ÈColons apply to stars known to be variable. TABLE 5 ABUNDANCES OF NGC 76 GIANTS Star Fe O Na Mg Al Si Ca Sc Ti V Ni Ba Eu Abundances for Individual Stars V19... [1.57 ].7 ].38 ].34 ].74 ].25 ].16 [.13 ].24 [.8 [.14 ]8 ].34 I-1... [1.5 [.15 ]2 ] ].71 ].17 ].23 ].18 ].3 [.1 ].6 ].23 ] II [1.54 [.1 ].38 ]4 ].75 ].29 ].31 [.8 ].31 ].6 ].1 ].81 ]4 II [1.56 ].14 ].29 ].35 ].63 ].26 ].35 [.18 ].18 ].1 [.7 ].22 ].35 V54... [1.6 ].1 ].22 ].3 ¹ ]4 ].32 ].13 [.22 ].22 [.12 ]. ].32 ].3 II [1.53 ].15 ]. ].33 ].53 ].28 ].21 [.8 ].6 [.8 [.8 ].7 ].33 Cluster Mean Abundances Mean... [1.55 ].5 ].28 ].36 ].63a ].26 ].23 [.9 ].22 [.4 [.4 ].36 ].36 s.d.b pc NOTE.ÈThe Fe abundances are written as [Fe/H] values; all other abundances are written as [el/fe] values. a Assuming that the upper limit to [Al/Fe] for V54 is correct. With V54 omitted, the mean is ].67. b Standard deviation of the mean. c Standard deviation of a single measurement.

7 156 KRAFT ET AL. Vol. 115 Abundance (a) [Fe/H] (mean) (b) [Sc/Fe] [V/Fe] [Ni/Fe] (c) [Si/Fe] [Ca/Fe] [Ti/Fe] (d) (e) (f) (g) (h) (i) [O/Fe] [Na/Fe] [Mg/Fe] [Al/Fe] [Ba/Fe] [Eu/Fe] T FIG. 2.ÈAbundances of all elements studied in this work plotted against T for the six NGC 76 giants observed with Keck I. Note that the ordinate for (a) is the [Fe/H] metallicity, while the ordinates of the remaining panels are the relative [el/fe] ratios. In (b), the abundances of Fe-peak elements Sc, V, and Ni are plotted together, since these elemental abundances are expected to be invariant from star to star. Likewise, in (c) the abundances of Si, Ca, and Ti, three a-capture elements likely to be constant in all stars, are plotted together. 4. ABUNDANCES 4.1. Results A panorama ÏÏ of our results is shown in Figure 2, which is a plot of the abundances of all elements measured versus T. A number of elementsèmg, Si, Ca, Ti, V, Ni, Fe, and EuÈhave an abundance scatter (p in Table 5) that is consistent with being entirely due to observational uncertainty. Whereas the cluster M13 exhibits Mg variations among its upper red giant branch stars, this appears not to be the case for our NGC 76 sample. In addition, whereas M15 exhibits Eu di erences among its red giants, the data provide no evidence of such an inhomogeneity for NGC 76. The mean [Fe/H] metallicity ÏÏ of [1.55 dex for our six-star sample is measured with a small internal uncertainty of ^.1 (p \.3), comparable to the internal error in metallicity found in previous papers of this series (e.g., Sneden et al. 1991, 1992, 1997; Kraft et al. 1993, 1995, 1997). There are no previous high-resolution studies of NGC 76 with which to compare our cluster metallicity, so we may remark here only that this mean value is, as expected, close to the metallicities of M3, M1, and M13 that we have determined in previous studies. These are listed in Table 6, column (2). The row marked Mean ÏÏ in Table 6 is the mean of [Fe/H] determinations for M3, M1, and M13, and * ÏÏ represents the di erence between [Fe/H] for NGC 76 and this mean. We list also in Table 6 values of [Fe/H] for these four clusters derived using a variety of low-resolution spectroscopic and photometric techniques employed over the past 2 years. Not all of these are of equal weight, and some are rather interdependent. Recent estimates have tended to reduce the metallicity of NGC 76 by.1è.2 dex (cols. [9] and [1] of Table 6). For example, in the Claria et al. (1994) recalibration of DDO photometry, [Fe/H] is decreased by.26 dex relative to the original Hesser, Hartwick, & McClure (1977) calibration. However, what counts in our discussion is the abundance of NGC 76 relative to clusters such as M3, M1, and M13. As Claria et al. (1994) point out, their recalibration moves all clusters on the average more metal-poor by.16 dex relative to Zinn & West (1984), and by.35 dex relative to Hesser et al. (1977). If we ignore M1, for which the reddening is poorly known (cf. Hurley, Richer, & Fahlman 1989; Arribas, Caputo, & Martinez-Roger 1991; Kraft et al. 1995), the reduction in [Fe/H] below the mean of M3 and M13 is.11 dex, using the Claria et al. (1994) calibration. We conclude that even if the metallicity of NGC 76 were that measured by Claria et al., its ranking relative to M3 and M13 would not be substantially changed. Finally, we formed the average value of * (Table 6) from all determinations of [Fe/H], excluding TABLE 6 METALLICITY ESTIMATES FOR NGC 76, M3, M1, AND M13 Cluster [Fe/H] [Fe/H] [Fe/H] [Fe/H] [Fe/H] [Fe/H] [Fe/H] [Fe/H] [Fe/H] (1) (2) (3) (4) (5) (6) (7) (8) (9) (1) Individual Clusters NGC [1.55 [15 [1.59 [1.6 [1.55 [15 [15 [1.77 [1.86 M3... [17 [1.67 [1.66 [1.5 [1.8 [17 [1.6 [1.64 [1.74 M1... [1.52 [1.26 [1.6 [ [ [1.56 [1.55 M13... [1.58 [1.6 [1.65 [1.6 [1.6 [17 [1.6 [1.6 [1.76 Mean Values Mean... [1.52 [1.51 [1.64 [1.57 [1.7 [19 [1.6 [1.6 [1.68 p *... [.3 ].6 ].5 [.3 ].15 ].4 ].15 [.17 [.18 REFERENCES.ÈCol. (2): this series, Kraft et al. 1993, 1995, 1997; col. (3): Searle & Zinn 1978; col. (4): Zinn & West 1984; col. (5): Hesser et al. 1977; col. (6): Frogel, Cohen, & Persson 1983, quoted in Zinn & West 1984; col. (7): Cohen & Frogel 1982; col. (8): Friel et al. 1982; col. (9): Buonanno et al. 1991, based on (B[V ) ; col. (1): Claria et al ,g

8 No. 4, 1998 PROTON CAPTURE CHAINS. III. 157 the present result, and obtained S*T \. ^.4 (p \.12). The a-nuclei elements Si, Ca, and Ti have [a/fe] enhancements of.1è dex relative to solar and, within the errors, appear not to be variable from star to star. This tends to also be the case for other halo globular clusters of similar metallicity to NGC 76, such as M3, M13, and M1, although the precise [a/fe] may vary by D.1È.2 dex from cluster to cluster. One point that we may infer from the a-element/fe enhancements is that NGC 76 formed within the Galactic halo before signiðcant numbers of Type I low-mass supernovae began to contribute Fe to the chemical enrichment of the Galaxy. The Fe-peak elements Sc, V, and Ni have abundances relative to Fe within ^.2 dex of the solar ratio, as is seen in other globular clusters such as M15, M92, M3, and M13 (Sneden et al. 1991, 1997; Kraft et al. 1992, 1993, 1997). In summary, we see no indications among the a-element and Fe-peak element abundances of the NGC 76 red giants to suggest that this cluster formed at an unusual location or an unusual time within the evolution of the Galactic halo. This is in distinct contrast to the unusually low [a/fe] ratios found recently in the relatively young ÏÏ globular cluster Rup 16 (Brown, Wallerstein, & Zucker 1997). In several globular clusters, e.g., M13, M15 (Sneden et al. 1997), and u Cen (Norris & Da Costa 1995b), [Mg/Fe] ratios vary from approximately ] to B. and are anticorrelated with [Al/Fe] ratios. This does not appear to be the case among giants in NGC 76, where S[Mg/Fe]T \].36 with very little scatter. We return to the signiðcance of this result further below. The elements exhibiting the largest star-to-star scatter are O, Na, Al, Sc, and Ba, but the evidence for Ba or Sc abundance di erences among our sample is not strong. The scatter in Eu could merely reñect observational errors incurred by using only one absorption line for the abundance analysis. For Sc, the large p of.14 is due almost entirely to one star, I-1, with an [Sc/Fe] abundance that is more than.25 dex higher than for the other Ðve stars. The star-to-star scatter in [O/Fe], [Na/Fe], and [Al/Fe], although small, is more likely to be real. This is illustrated by Figure 3, which shows plots of [Na/Fe] and [Al/Fe] versus [O/Fe]. These plots show that the NGC 76 sample can be roughly divided into two groups: (1) stars I-1, V19, and II-13 have [O/Fe] \.9, [Na/Fe] [.35, and [Al/Fe] [.7, whereas (2) stars II-46, II-18, and V54 have [O/Fe] [.9, [Na/Fe] \.35, and [Al/Fe] \.7. Dotted horizontal and vertical lines are added to Figure 3 to illustrate these abundance divisions. There appears to be an anticorrelation between [O/Fe] and both [Na/Fe] and [Al/Fe] among the two groups of stars. These anticorrelations suggest that the scatter in O, Na, and Al is real. This is strengthened by the fact that anticorrelations have been found between these same elements in a number of halo globular clusters (see, e.g., Norris & Da Costa 1995a, 1995b and Zucker et al for u Cen; and the summaries of Kraft et al. 1993, 1995, 1997; Sneden et al. 1992, 1997; Minniti et al. 1996; and Shetrone 1997 for other clusters) Sources of Abundance Errors Color Excess The color excess E(B[V ) appears to be somewhat uncertain and ranges from the generally assumed value of.5 [Na/Fe] [Al/Fe] [O/Fe] FIG. 3.ÈProton-capture element abundance correlations. Top: [Na/Fe] vs. [O/Fe] is plotted for the NGC 76 red giants. Two dotted straight lines are shown in the Ðgure, a vertical line at [O/Fe] \].9 and a horizontal line at [Na/Fe] \].35. These lines divide the diagram into four quadrants, only two of which are occupied by NGC 76 giants. This indicates that there is an O-Na anticorrelation among the NGC 76 giants, even though the scatter among both [O/Fe] and [Na/Fe] is relatively small (Table 5). Bottom: [Al/Fe] vs. [O/Fe] is plotted, along with a vertical line at [O/Fe] \].9 and a horizontal line at [Al/Fe] \].65. This division into quadrants makes apparent the O-Al anticorrelation. Ðrst estimated by Searle & Zinn (1978; cf. Djorgovski 1993) to the value.1 adopted by Claria et al. (1994). The interstellar reddening map of Burstein & Heiles (1982) suggests a value near.9. Adoption of E(B[V ) \.9 instead of.5 would have required an increase in T by 6 K, which we could not have discerned, within the errors, in the diagnostic plot of log v(fe I) versus excitation potential. The Fe abundance would typically have been raised by.3 and log g(obs) by.15 (Table 7) to keep pace with the required ionization equilibrium of Fe I and Fe II. This modiðcation would serve to bring log g(pred) and log g(obs) into slightly closer agreement (Table 4). Corresponding changes in the key [el/fe] ratios would generally be negligible (Table 7).

9 158 KRAFT ET AL. Vol. 115 TABLE 7 CHANGES RESULTING FROM AN E(B[V )INCREASE FROM.5 TO.9 Parameter Change T (K)... ]6 log g... ].15 [Fe/H]... ].3 [O/Fe]... ].4 [Na/Fe].... [Mg/Fe]... [.2 [Al/Fe]... [.3 [Si/Fe]... [.3 [Ca/Fe].... [Sc/Fe]... ].2 [Ti/Fe]... ].5 [V/Fe]... ].7 [Ni/Fe]... ] Uncertainties in [el/fe] Ratios In Table 8, we tabulate changes in derived [el/fe] ratios induced by modiðcations in input parameters (e.g., T, log g) corresponding to plausible systematic errors (^1 K in T, ^.25 in log g, and ^.25 km s~1 in v ). These are t based on numerical experiments associated with the analysis of a star (II-13) having a T (42 K) lying in the middle of our six-star sample. The derived value of [Fe/H] is essentially unchanged by plausible variations in T, since the variations in Fe I and Fe II have opposite signs; [el/fe] ratios for key elements such as oxygen, sodium, magnesium, and aluminum do not exceed.1 dex, and are often much smaller. The size and sign of these variations correspond to those found earlier among giants of similar T, log g, and [Fe/H] in M13 (cf. Kraft et al. 1997; Shetrone 1996b) using similar observations and analysis procedures. Values of [O/Fe] were obtained in all cases by spectrum synthesis of the [O I] jj63, 6364 lines, in order to allow for the slight blending of the Sc II line at A. Except for V54, the EWs of the j6364 line were never less than 1 ma, and the oxygen abundance di erence between the two [O I] lines never exceeded.1 dex. In the case of V54, the generally poorer S/N did not permit us to make a reliable measurement of j6364; lower weight has been assigned to [O/Fe] for this star. [Al/Fe] ratios were derived from the EWs of the jj6696, 6699 pair in three of the NGC 76 giants, but in the other three stars, only the line j6699 was available. Global di erences in the Al abundance of giants in NGC 76 versus M13 are of interest as a diagnostic of the second-parameter e ect ( 5), and we therefore discuss the derivation of Al abundances in more detail in Appendix B. TABLE 8 ABUNDANCE UNCERTAINTIES (BASED ON STAR II-13) d[species] dt d(log g) dv t d[fe/h]: Fe I... ].11 ([.8) ].2 ([.1) [.8 (].14) Fe II... [.1 (].12) ].12 ([.12) [.2 (].4) Average... ].1 (].2) ].7 ([.6) [.5 (].9) d[o/fe]... ].3 ([.3) ].1 ([.3) ].5 ([.1) d[na/fe]... ].9 ([.12) [.1 (].9) ].4 ([.7) d[al/fe]... ].6 ([.9) [.9 (].7) ].3 ([.5) d[mg/fe]... ].5 ([.7) [.9 (].9) ].1 ([.4) d[eu/fe]... [.2 (.) ].3 ([.4) ].4 ([.6) 5. DISCUSSION Among the halo clusters ([Fe/H] \ [1) that have been studied at high spectral resolution with large giant-star samples (M92, M15, u Cen, M3, M13, M1, NGC 76, NGC 6752, and M5), 6 NGC 76 seems to be singular in terms of the small range of O-Na-Al abundances detected among its red giant stars. One possibility that can only be assessed by further observation is whether this is a chance result of the small number of giants observed. Perhaps these six giants do not completely sample the real abundance distribution in NGC 76. For the purposes of the following discussion we assume that this is not the case, although this uncertainty serves as a caveat to any conclusions that we draw. As noted in 1, di erences in the abundances of O, Na, and Al among red giants within globular clusters such as M3 and M13 are commonly interpreted within a deep mixing ÏÏ paradigm. However, in assessing the role of deep mixing on the blueward extension of the HB, the critical issue is whether the mixing is deep enough to increase the He content of the stellar envelope (Langer & Ho man 1995; Sweigart 1997). Langer et al. (1997) argue that a signiðcant increase in He content may accompany any large depletion of O (and Ne) and corresponding enhancement of N (and Na). However, if the mixing is deep enough to increase the envelope abundance of Al at the expense of Mg, then an increase in the He abundance is assured. How much fresh He is produced in a given giant, and therefore how far to the blue will be its HB descendant, is monitored by the degree to which Mg is depleted and Al is enhanced in the giant. We argue in the following that NGC 76 giants su er little Mg depletion and Al enhancement relative to M13 giants. The deep-mixing model encounters some difficulties with respect to the interior structure of low-mass, low-metallicity red giants. We discuss these brieñy in 6, along with an alternative scenario TheOand Na Abundances Interpretation of the NGC 76 oxygen abundances is made difficult by the fact that there are no stars in our sample that have [O/Fe] \].3È, the usual baseline oxygen abundance found in other globular clusters. This leaves open two rather di erent possibilities. 1. If we assume that the initial oxygen abundance was [O/Fe] \], then we would conclude that all of the stars in our sample have experienced some degree of interior mixing that has lowered the [O/Fe] ratio by.25è.55 dex. Although all stars have sustained some deep mixing, the amount of this mixing has been variable from star to star. 2. If instead we assume that stars II-46 and II-18 ([O/Fe] D ].15) reñect the initial oxygen abundance of ÈÈÈÈÈÈÈÈÈÈÈÈÈÈÈ 6 Spectra of seven M22 giants (S[Fe/H]T D [1.65) have been studied by Brown, Wallerstein, & Oke (199) and Brown & Wallerstein (1992), who found a noticeable variation in [Fe/H], a small range in [O/Fe] (with most stars near ].3), nonconstancy of C ] N ] O, and a correlation of N with Na, all of which point to the existence of signiðcant primordial variations among stars in this cluster. In another large-scale, but photometric, investigation, Anthony-Twarog, Twarog, & Craig (1995) found [Fe/H] constant, [Ca/Fe] variable, and a signiðcant variation in E(B[V ) from star to star. We do not attempt to discuss the complicated, and somewhat controversial, abundance patterns of M22 in this paper.

10 No. 4, 1998 PROTON CAPTURE CHAINS. III. 159 NGC 76, then we would conclude that these two stars have sustained no deep mixing, while three other stars in our sample (I-1, V19, and II-13) have incurred moderate amounts of interior mixing that has reduced the [O/Fe] ratio by D.1È.3 dex. What arguments can be o ered to support these two possibilities? If favor of (1), we note that all of the stars in our sample are very bright. By analogy with M13 and M1, one tends to Ðnd a higher ratio of extremely oxygen-depleted giants to oxygen-normal giants in the upper 1 mag of the RGB. This suggests that our sample would be expected to contain some of the most severely mixed stars in NGC 76. Thus, assuming that they have all experienced some amount of mixing seems reasonable if NGC 76 behaves like other clusters. Further evidence for the Ðrst possibility comes from Figure 4. This shows a plot of [Na/Fe] versus [O/Fe] for the clusters M92, M15, M3, M1, and M13, as well as our results for NGC 76. In the former Ðve clusters a much wider range of O and Na abundances are observed than for NGC 76. Since the highest [O/Fe] ratio in the clusters shown is approximately ], the data points in Figure 4 are assumed to show how [Na/Fe] varies as a function of [O/Fe] for material with an initial [O/Fe] \] that is subjected to variable amounts of Ne-Na cycle nucleosynthesis within the O ] N processing shell of Population II red giants with metallicities [2.5 \ [Fe/H] \ [1.5. The data points for NGC 76 fall right in among the points for the other clusters and, thus, are compatible with what would be expected from the conversion of O into Na for material with an initial [O/Fe] \]. However, even if the initial value of [O/Fe] was as high as ], the degree to which O depletion and Na enhancement occurs among giants in NGC 76 remains quite modest in comparison with giants in M1 and especially M13, and is rather similar to what is found among giants in M TheMg and Al Abundances A notable feature of the NGC 76 giants is that the average [Al/Fe] ratio is moderately high (].63 dex; see Table 5) and shows a signiðcant but still relatively small dispersion. This is in contrast with giants of M13, where [Al/Fe] ranges over 1.3 orders of magnitude, from [.1 to ]1.2 (Kraft et al. 1997). A wide range of [Al/Fe] values is also found in the more metal-poor cluster M15 (Sneden et al. 1997), but in this case the lowest value of [Al/Fe] is near ].2. These results suggest that the products of Mg-to-Al conversion within the hydrogen-burning shells of NGC 76 giants have indeed been conveyed by deep mixing into the atmospheres of these stars. This inference is supported by the anticorrelation of O and Al exhibited in Figure 3. The behavior of Al is further illustrated by Figure 5, which shows [Al/Fe] versus [Mg/Fe] for the giants in NGC 76, as well as M15, M92, and M13 (Kraft et al. 1997; Sneden et al. 1997). This Ðgure seems to indicate that NGC 76 giants could initially have begun with [Al/Fe] ratios B. and achieved [Al/Fe] ratios near ].63 as a result of a modest degree of mixing (i.e., modest by comparison with giants in M13 and M15). Our sample of giants does not extend to faint enough luminosities to permit determination of the baseline abundance of Al, but if it is as low as has been found previously in M13 and M15 (or even somewhat higher), then we can conclude that the degree of deep mixing in NGC 76 giants is considerably less than in the giants of M13 and M15. An apparent difficulty with this scenario is that the Mg abundances observed among the NGC 76 giants do not anticorrelate with Al, as seen in clusters like M15, M13, and u Cen, for which deep mixing is also invoked to explain the abundance pattern seen on the RGB. However, this is readily explained if the initial [Mg/Fe] and [Al/Fe] ratios were ] and., respectively, and a modest degree of deep mixing increased [Al/Fe] to ].65. Because the initial Mg abundance in this case would be 4 times larger than [Al/Fe] NGC 76 M15 [Na/Fe] - NGC 76 M15 M1 M3 M92 M [O/Fe] M92 M [Mg/Fe] FIG. 4.È[Na/Fe] vs. [O/Fe] for red giants in the globular clusters M92, M15, M3, M13, M1, and NGC 76. The points for the NGC 76 giants fall within the region occupied by giants in the other clusters, but show a dispersion much smaller than exists among giants in M13. FIG. 5.È[Al/Fe] vs. [Mg/Fe] for red giants in the globular clusters M92, M15, M13, and NGC 76. Once again the dispersion among the points for NGC 76 giants is small compared with both M13 and M15 giants.