Chemical Abundances in Field Red Giants from High-Resolution H-Band Spectra using the APOGEE Spectral Linelist
Verne V. Smith, Katia Cunha, Matthew D. Shetrone, Szabolcs Meszaros, Carlos Allende Prieto, Dmitry Bizyaev, Ana Garcia Perez, Steven R. Majewski, Ricardo Schiavon, Jon Holtzman, Jennifer A. Johnson
aa r X i v : . [ a s t r o - ph . S R ] D ec Submitted to
The Astrophysical Journal
Chemical Abundances in Field Red Giants from High-ResolutionH-Band Spectra using the APOGEE Spectral Linelist
Verne V. Smith and Katia Cunha
National Optical Astronomy Observatories, Tucson, AZ 85719 USA andObservatorio Nacional, Sao Cristovao, Rio de Janeiro, Brazil [email protected]
Matthew D. Shetrone
Department of Astronomy and McDonald Observatory, University of Texas, Austin, TX78712 USA
Szabolcs Meszaros and Carlos Allende Prieto
Instituto d’Astrofisica de Canarias, 38205, La Laguna, Tenerife, Spain
Dmitry Bizyaev
Apache Point Observatory, Sunspot, NM 88349 USA andStrenberg Astronomical Institute, Moscow 119992, Russia
Ana Garcia Perez and Steven R. Majewski
Department of Astronomy, University of Virginia, Charlottesville, VA 22904 USA
Ricardo Schiavon
Astrophysics Research Institute, Liverpool John Moores University, Liverpool L3 5UX UK
Jon Holtzman
Department of Astronomy, New Mexico State University, Las Cruces, NM 88003 USA
Jennifer A. Johnson
Department of Astronomy, Ohio State University, Columbus, OH 43210
ABSTRACT α Boo and µ Leo), two M-giants ( β And and δ Oph), and onethermally-pulsing asymptotic giant branch (TP-AGB) star of spectral type MS(HD 199799). Measured chemical abundances include the cosmochemically im-portant isotopes C, C, N, and O, along with Mg, Al, Si, K, Ca, Ti, V,Cr, Mn, Fe, Co, Ni, and Cu. The K and M giants exhibit the abundance signa-ture of the first dredge-up of CN-cycle material, while the TP-AGB star showsclear evidence of the addition of C synthesized during He-burning thermalpulses and subsequent third dredge-up. A comparison of the abundances derivedhere with published values for these stars reveals consistent results to ∼
1. Introduction
The Apache Point Observatory Galactic Evolution Experiment (APOGEE) is one of4 experiments that are part of the Sloan Digital Sky Survey III (SDSS-III; Eisenstein etal. 2011). APOGEE is obtaining high-resolution (R ∼ ≥
100 per pixel), H-band ( λ µ m) spectra of evolved, late-type stars, with thegoal being to measure chemical abundances of 15 elements per star. Coupled with radialvelocities that are accurate to ∼
100 m-s − , APOGEE is creating the first high-precisionspectroscopic and radial-velocity survey of all Galactic stellar populations (bulge, bar, disks,halo) using a uniform set of stellar tracers and spectral diagnostics, with a plan to observeand derive these parameters for 100,000 red giants by the end of 2014.The physical and chemical parameters of the APOGEE stars are derived from a suite ofsoftware packages that, together, are called the APOGEE Stellar Parameters and Chemical 3 –Abundances Pipeline (ASPCAP); the details of this software and its analysis techniques willbe discussed by Allende-Prieto et al. (2012, in preparation). The first-generation output ofASPCAP, which will be based on 1-D stellar atmosphere analyses in LTE, will consist of thestellar parameters of effective temperature, T eff , surface gravity (noted as log g), and the1-D microturbulence parameter ( ξ ), along with the abundances of up to 16 elements (Fe, C, N, O, Mg, Al, Si, K, Ca, Ti, V, Cr, Mn, Co, Ni, or Cu) and, for many stars, the Cisotopic abundance.One of the central components of the ASPCAP machinery is a spectral linelist con-structed to produce the synthetic spectra that are used to compare to the observed ones;the APOGEE linelist will be presented and discussed in detail by Shetrone et al. (2012, inpreparation).The goal of this study is to analyze a small number nearby field red giants with well-defined stellar parameters and use the APOGEE linelist to derive detailed chemical abun-dances via standard “manual” abundance analysis techniques (Ramirez and Allende Prieto2011 is a recent example using such techniques). The derived stellar parameters and chemicalabundances can be compared to the same quantities derived via ASPCAP using the samespectra, linelist, and model atmosphere grid. This comparison will be an important test ofASPCAP. This work will also present new abundances in these stars based on an extensiveand up-to-date spectral linelist in the near-IR H-band.Abundances will be presented in this paper using two different types of nomenclature,with one abundance scale defined as A(X)=Log(N(X)/N(H)) + 12.0. The other scale com-pares the element ratio of N(X)/N(Q) in a target star to the Sun via [X/Q]=(A(X)
TargetStar – A(X) ⊙ ) – (A(Q) TargetStar – A(Q) ⊙ ).
2. Observational Data: The Fourier Transform Spectra of Bright Field RedGiants
The APOGEE high-resolution spectra cover the wavelength range from ∼ λ α Boo.The originally reduced spectra of selected stars were recorded as flux versus wavenumber(cm − ) over different dates and runs and covered a broad region in the H-band. The spectra 4 –analyzed here were resampled in wavelength space and restricted to a wavelength range of λ
3. The “Standard Star” Parameters3.1. The Choices
There are five stars analyzed here: the well-studied α Boo, a mildly metal-poor K1.5giant, µ Leo, the prototype metal-rich star (Spinrad & Taylor 1969), which is a K1 giant,two near-solar-metallicity M giants, β And (M0III) and δ Oph (M0III), and HD199799,an AGB star of spectral type MS with a mild enhancement of C and a slightly elevatedratio of C/O (relative to solar) due to third dredge-up (Smith & Lambert 1990). Figure 1illustrates a small portion of the APOGEE wavelength coverage for four of the red giantshere. This region from 15560 – 15590˚A includes representative species and how they varywith spectral type (or effective temperature) and metallicity. Lines from the main CNO-containing molecules, CO, OH, and CN fall within this region, as noted in the figure. Ofparticular interest is the metal-rich K-giant µ Leo, which exhibits quite strong CN lines;much of the variation in the spectrum of µ Leo compared to the other giants result from thehigh line-density of CN. This figure also points to the desirability of analyzing large numbersof red giants via spectrum synthesis techniques using a detailed linelist.With a significant fraction of its red-giant targets in the thin and thick disks, plus theinner bulge and bar, the five types of red giants studied here span a range of temperatureand metallicity that provides a good test for ASPCAP. The APOGEE survey targets redgiants with effective temperatures from T eff =3400K–5000K, having surface gravities fromlog g=3.0 to -0.5, and metallicities from [Fe/H]= -5.0 to +1.0. These nearby giants havewell-measured photometry and parallaxes, so temperatures, luminosities, and masses are allrather well-constrained. In addition, four of the five stars have previous abundance analysesin the literature and thus provide comparison tests for the abundances derived from theAPOGEE linelist. 5 – Deriving stellar chemical abundances from high-resolution spectra requires as basic inputthe effective temperature (T eff ), surface gravity (characterized as log g with g as cm-s − ),and overall metallicity as the fundamental parameters of the model atmosphere. In addition,high-resolution abundance analyses of cool stars that use static 1-D stellar atmospheres alsorequire the derivation of a non-thermal Doppler-like broadening term called microturbulence( ξ ), which is determined as part of the spectral analysis. The first step is the determinationof T eff and, for these relatively nearby red giants, is set by the near-infrared magnitudes ofJ and K. This particular choice is used because the APOGEE targets are selected from the2MASS catalog (Strutskie et al. 2006) and therefore all have (J–K) available. The aim is tofocus on utilizing near-IR data and then compare to published literature results.All of the red giants here have parallaxes measured by Hipparcos and thus have distancesknown to various levels of accuracy. The parallax is used to set the distance and, thus, M K .Bolometric corrections from Bessell, Castelli, & Plez (1998) are used in conjunction with(J–K) to determine M bol and consequently stellar luminosity. Basic data for the red giantsare presented in Table 1.The most important derived quantity affecting the overall shape of the spectrum isT eff , which here is based on an average of two T eff -(J–K) calibrations; one from GonzalezHernandez & Bonifacio (2009) and the other from Bessell et al. (1998). The GonzalezHernandez & Bonifacio calibration is defined in the 2MASS photometric system of (J–K S ),so the original Johonson (1965) (J–K) colors and K magitudes for these bright red giants(which have saturated magnitudes in the 2MASS catalog) were transformed to 2MASS valuesusing the prescription from Carpenter (2001). The differences between the two calibrations,for a given metallicity, are less than 30-50K for T eff ≤ S ) (where wehave transformed the Bessell et al. scale onto the 2MASS system). There is also a smallmetallicity effect on the (J–K S ) calibration which, for a given (J–K S ) can be ∼ eff calibrations of less than 30-50K over the temperature range of thetarget stars, as well as their modest range in metallicity, and well-determined magnitudeswith ∼ S ) (which translates to ∆T eff =30K), the derived values ofT eff have uncertainties of ±
50K or less.Given the stellar luminosity and effective temperature, these fundamental propertiesare then compared to stellar evolution models presented in Bertelli et al. (2008; 2009, withdata from these papers taken from the website http://stev.oapd.inaf.it/YZVAR ) to estimate 6 –stellar masses and, ultimately, surface gravities through:g/g ⊙ = (M/M ⊙ ) × (L ⊙ /L) × (T eff /T eff ⊙ ) .Determining the surface gravity from evolutionay tracks is an iterative process, becauseit requires knowledge of the stellar metallicity. This is illustrated in Figure 2 where the stellarluminosity and T eff are plotted along with evolutionary models from a range of masses, witheach panel having a different metallicity: Z=0.017 (top panel), 0.008 (middle panel), and0.004 (bottom panel), correpsonding to overall metallicities of [m/H]= +0.07, -0.26, and-0.56, respectively, when taking Z=0.0145 for the Sun, as suggested by Lodders (2010). Theillustrated evolutionary tracks follow stellar evolution up to the tip of the first ascent redgiant branch (which we label RGB). Initial model metalliities for the stars studied here arebased on previously published results, with β And, δ Oph, and HD199799 started with solarmetallicity, α Boo with [Fe/H]= -0.5, and µ Leo with [Fe/H]= +0.3. As there are numerous FeI lines in this spectral region, a sample of Fe I lines are used to determine the iron abundance.Given the limited range of metallicity spanned by the nearby red giants studied here, theFe abundance is a good proxy for the overall stellar metallicity. If the derived Fe abundancediffered by more than 0.1 dex from the assumed value, a new surface gravity was derivedand a new model generated in order to reanalyze the Fe I lines until convergence betweenmodel abundance and derived abundance. For these relatively well-studied red giants, nomore than two iterations were required until convergence. Table 2 presents, for these redgiants, the derived parameters T eff , Log (L/L ⊙ ), mass, log g, microturbulent velocity ( ξ – seeSection 4.2), and overall metallicity, which is represented by [Fe/H] as discussed in Section4.2. Given the position of the stars in the Log L versus T eff plane shown in Figure 2, it canbe ascertained that 4 stars ( α Boo, µ Leo, δ Oph, and β And) are most likely first-ascentred giants of low to moderate mass. Based on its warmer temperature and somewhat lowerluminoisty, it may be more likely that µ Leo is in, or very near the core-He burning phase(clump giant) and has thus already ascended the RGB and experienced the core-He flash. Itis also possible that the other 3 stars already noted above may be post-He core-burning starsascending the AGB; however, lifetimes on the two separate evolutionary sequences wouldfavor them as being on the RGB. All fall above the “luminosity bump” (Fusi Pecci et al.1990) and would have thus experienced any extra-mixing that may occur during that phaseof stellar evolution. 7 –The coolest and most luminous red giant in this sample, HD199799, is an AGB star thatis experiencing the third dredge-up, as evidenced by the observation of Tc I in its spectrumand its elevated s-process abundances, such as [Y/Fe] ∼ +0.5 or [Nd/Fe] ∼ +0.4 (Smith &Lambert 1988; 1990).
4. The Abundance Analysis ∼ chris/moog.html).An extensive effort was made to produce as complete a spectral linelist as possible, sothe analysis technique used here, as in APOGEE in general, is to utilize spectrum synthesisin a quantitative comparison between synthetic and observed spectra. All abundances aredetermined for each spectral line, or group of closely spaced molecular lines, via minimizingthe residuals between observed and synthetic spectra as a function of the abundance inquestion. Due to both its relatively large cosmic abundance and large number of energy levels,iron has historically been used as a diagnostic species for certain stellar parameters, as wellas an overall metallicity indicator, so the results for Fe I lines selected from the APOGEElinelist (Shetrone et al. 2012) are discussed here. In particular, Fe I lines are well-suited forsetting the microturbulent velocity, ξ , that is required for 1-D model atmosphere abundanceanalyses.A large number of Fe I lines fall within the APOGEE spectral window (15100 – 16900˚A), although there are no Fe II lines strong enough to be detected in these red giants. Since 8 –all abundance determinations are done via spectrum synthesis, the Fe I lines were culledto include those lines which contain only contributions from Fe I (this was determined bypreliminary synthesis of the red giant spectra, then setting the Fe abundance to zero andensuring that the spectral feature in question vanished); since the Fe I lines are so numerous,this stringent selection still results in an adequate list of lines. Table 3 presents the Fe I linesused for the iron abundance determinations. As is true of most of the atomic lines in theH-band, the excitation energies are dominated by rather high-excitation lines, although forFe I there are two lower-excitation lines.The Fe I lines span a large enough range in line-strength that they can be used to setthe microturbulent velocity. This parameter is defined by the value of ξ that produces notrend in the Fe abundance as a function of line-strength. For a range of microturbulentvelocities Fe abundances were determined from each Fe I line and the adopted value of ξ wasone in which a linear regression between A(Fe) as a function of log(W λ / λ ) yielded a slopeof zero. This procedure is illustrated graphically in Figure 3, where the abundances derivedfor each individual line is plotted versus ξ (the continuous lines); the circle with errorbarsshows where the slope of A(Fe) with line-strength goes to zero.The microturbulent velocities set by the Fe I lines, as well as the Fe abundances arelisted in Table 2 as part of the derived set of stellar parameters. The individual line-by-lineabundances as determined for each Fe I line are also listed in Table 3. The Fe abundanceswere used to set the final overall metallicities of the ATLAS9 model atmospheres for theanalysis of additional elements via atomic or molecular lines. Abundances for C, N, O, and the minor carbon isotope C are derived fromcombinations of vibration-rotation (V–R) lines of CO (X Σ + ) and OH (X Π), along withelectronic transitions of CN (A Π – X Σ). Although the details of the ingredients that wentinto these molecular lines in the APOGEE linelist will be found in Shetrone et al. (2012 –in preparation), a few highlights are noted here. The adopted dissociation energies (D ) are11.092 eV for CO, 4.395 eV for OH, and 7.76 eV for CN. The gf-values are from Goorvitch(1994) for CO, Goldman et al. (1998) for OH, and the CN gf-values are taken from thebaseline linelist of Kurucz (1993), with updates to these values from the prescription givenin Melendez & Barbuy (1999).The procedure used for the CNO analysis, since all red giants studied here have C/O ≤
1, is to use CO to set the carbon abundance. With that carbon abundance, OH provides 9 –an O-abundance, which, if different from the initial value (taken to be the solar value scaledby the stellar value of [Fe/H]), is used to re-analyze the CO lines. This process is repeateduntil the C and O abundances yield consistent abundances from CO and OH. With thesevalues of C and O, the CN lines are used to derive the abundance of nitrogen. In general,the abundance of N has little to no effect on the CO and OH lines, however, the respectivefinal C, N, and O abundances all provide self-consistent results from CO, OH, and CN.All spectral features were synthesized and the molecular regions that were used toextract the abundances are listed in Table 4. In this case, if a range of wavelengths is givenit indicates that a spectral window containing numerous molecular lines was used (for COand OH), while for CN individual transitions were identified and synthesized. As with Table3, the abundances derived for each feature are also listed in Table 4. In the case of C, ashas been customery for minor isotopes, the abundance is listed as the ratio of C/ C, withA( C) being set by the mean abundance of carbon-12.
With the Fe and C, N, and O abundances in hand, additional elements are derived viaatomic spectral lines in the wavelength interval 15100 – 16900˚A. The APOGEE linelist wassearched for suitable species, all neutral over the T eff range of the red giants here, and some 12other elements were deemed detectable and able to be analyzed quantitatively: Mg I, Al I, SiI, K I, Ca I, Ti I, V I, Cr I, Mn I, Co I, Ni I, and Cu I. The additional atomic lines are listedin Table 5. As with the molecular lines, if a precise wavelength is shown, the transitionin question is represented by a single line, however if the wavelength is approximated byan integer value it indicates multiple transitions due to hyperfine (hfs) splitting or isotopicsplitting. In the case of isotopic components, solar-system isotopic fractions are assumed and,in general, this assumption has no significant effect on the derived total atomic abundances.Figure 4 illustrates the fitting procedure for an atomic line from Mn I in δ Oph, with theunderlying hfs components illustrated, as well as nearby blending features, with severalsynthetic spectra plotted which have differing Mn abundances. As with the previous Tables3 and 4, individual abundances are provided for each line or feature. As noted previously,the details of the generation of the APOGEE linelist will be found in Shetrone et al. (2012).Table 6 provides a summary table of all abundances listed as their mean values andstandard deviations from the lines or spectral intervals shown in Tables 3, 4, and 5. In caseswhere only one line or region was observed, no standard deviation is given. The abundanceof C is given as the isotopic ratio of C/ C. 10 –
Most of the abundances presented here are derived from more than one atomic or molec-ular line or features of lines, thus internal consistency in the average abundances are definedby the scatter from the individual features. This scatter is characterized by the standarddeviations of Table 6. Because multiple features arise from differing line strengths (weakversus strong), excitation potentials, or ionization fractions, these respective standard devi-ations provide some estimate as to how well the 1-D LTE analysis can recover the detailsof the underlying stellar spectrum from the basic model atmosphere parameters T eff , logg, metallicity, and derived value of ξ . The typical standard deviations from Table 6 are ∼± µ Leo, which exhibits exceptionally strongCN lines, that are numerous throughout this H-band window. The background CN absorp-tion, creates blending which lowers the overall signature of other abundances and increasesthe uncertainty in these other abundances; this is reflected in the generally larger standarddeviations found in the µ Leo abundances.Besides internal consistency from the various spectral lines of different species, it is usefulto explore their sensitivities to stellar parameters in order to characterize the magnitudes ofpossible systematic effects due to uncertainties in the four primary parameters of effectivetemeprature, surface gravity, metallicity, and 1-D derived microturbulent velocity. For thisdiscussion T eff is written as T, log g as G, model atmosphere metallicity ([m/H]) as m, andmicroturbulent velocity as ξ . Because the abundance for a given element, A=log[N(A)/N(H)]+ 12. = A(T,G,m, ξ ), the incremental change in this abundance due to small variations inthe parameters is given bydA=( ∂A/∂T )dT + ( ∂A/∂G )dG + ( ∂A/∂ξ )d ξ + ( ∂A/∂m )dm.The scatter in the abundance caused by small changes in the stellar parameters is thendefined as σ = h dA i and dA can be written as(dA) =[( ∂A/∂T )dT] + [( ∂A/∂G )dG] + [( ∂A/∂ξ )d ξ ] + [( ∂A/∂m )dm] +( ∂A/∂T )( ∂A/∂G )[dTdG + dGdT]+( ∂A/∂T )( ∂A/∂ξ )[dTd ξ + d ξ dT]+( ∂A/∂T )( ∂A/∂m )[dTdm + dmdT]+( ∂A/∂G )( ∂A/∂ξ )[dGd ξ + d ξ dG] 11 –+( ∂A/∂G )( ∂A/∂m )[dGdm + dmdG]+( ∂A/∂ξ )( ∂A/∂m )[d ξ dm + dmd ξ ].The partial derivative terms are calculated by changing each model parameter anddetermining the resulting change in logarithmic abundance, A. Since the perturbation instellar parameters is relatively small, the changes are approximated by linear trends. Asthis paper is a limited study of a few stars as a test of the APOGEE linelist, a test ofabundance sensitivity to stellar parameters is carried out for a representative model takenwith T eff =4000K, log g=1.3, [m/H]=+0.0, with ξ =2.0 km-s − . Predicted equivalent widthswere generated for the lines used in the abundance analysis and the stellar parameters werethen pertrubed by dT=+50K, dG=+0.2 dex, dm=+0.1 dex, and d ξ =+0.2 km-s − , with newabundances derived for each separate perturbation, giving the change in abundance for eachparameter change. These coefficients, which reflect the change in abundance, are listed inTable 7 and provide an overall view of how sensitive the derived abundances are to changesin stellar parameters that are representative of realistic errors for these particular red giants.Columns 6 and 7 in Table 7 provide the values of dA two cases: one using only thefirst four terms which contains the change in abundance due to changes in each primaryparameter (dA’ in column 6, which is valid if all parameters are independent of each other),while column 7 includes the covariant terms. Within the analysis technique applied here,it was found that the microturbulent velocity did not measurably depend on the derivedeffective temperature or gravity, and the dT-d ξ and dG-d ξ cross terms are not included. Ascan be seen by the comparison of dA’ and dA, the covariant terms contribute significantly,in many species, to the uncertainty. The largest covariant terms are typically those betweenT eff and log g, caused by the slope of the RGB in the T eff –log(L/L ⊙ ) plane (since L ∼ T /g). The abundances derived here rely on the APOGEE H-band spectral window from λ λ α Boo, µ Leo, β And, and HD 199799. 12 – α Boo
Being the nearest red giant, α Boo is often used as the standard star of choice for abun-dance comparisons. Ramirez & Allende Prieto (2012) have produced a recent and detailedhigh-resolution optical spectral analysis, revieiwng both fundamental stellar parameters forthis star, as well as deriving an extensive set of elemental abundances. They used the ob-served spectral energy distribution and compared to theoretical ones from the Kurucz gridof no-overshoot model atmospheres with α -element enhanced compositions ([ α /Fe]=+0.4;Castelli & Kurucz 2003) in order to set the T eff for α Boo. The flux from λ ∼ µ mwas used with the resultant best-fit T eff =4286 ± eff =4275 ± α Boo, as the effective temperature and parallax are known to highprecision. The mass was set by comparison to Yonsei-Yale isochrones (Yi et al. 2001; Kimet al. 2002) and was found to be M=1.08 ± ⊙ , which yields a surface gravity of logg=1.66 ± eff =4275K value from the (J–K) color is adopted, whilethe Ramirez & Allende Prieto log g is used. These stellar parameters were used to generatea model atmosphere as described in Section 4 with the ATLAS9 code in the same manneras is being done for the APOGEE targets (Meszaros et al. 2012).The microturbulence, ξ , derived by Ramiriz & Allende Prieto was set by Fe I lines andwas found to be ξ =1.74 km-s − . The Fe I lines used from the APOGEE linelist yielded avalue of ξ =1.85 ± − : very close to that derived from optical Fe I lines.A direct comparison of the derived abundances here with those from Ramirez & AllendePrieto (2012) using completely differsent sets of lines leads to the following mean differenceand standard deviation for 13 elements in common of ∆A(x)(This paper – Ramirez & AllendePrieto)=+0.03 ± µ Leo
This K-giant is often cited as the “prototype” for metal-rich stars, having been notedas such by Spinrad & Taylor (1969), and some of its spectral regions contain numerous 13 –and strong CN lines. Two older abundance studies in the literature (Gratton & Sneden1990; Smith & Ruck 2000) are used to compare to the abundances derived here, as well asa more recent analysis by Fulbright, McWilliam, & Rich (2007–FMR). As with the α Boocomparison in Section 4.5.1, both of these published analyses are based on high-resolutionspectra in the visual. Both studies used T eff =4540K for µ Leo, which is similar to the value4550K adopted here. The surface gravities were slightly different, with log g= 2.3 ± ± ξ =1.8 km-s − found from the H-band Fe I lines is slightly higher than1.2 km-s found by both Gratton & Sneden (1990) and Smith & Ruck (2000), but with allstudies having uncertainties of 0.3-0.5 km-s − , this offset cannot be claimed to be significant.A comparison of abundances between here and Gratton & Sneden (1990) for 12 elementsfinds a mean ∆A(x)(This paper - GS)=-0.05 ± ± µ Leo by FMR was part of their study of bulgered giants. Their derived stellar parameters were quite similar to those found here, withT eff (FMR)=4520K, (∆T(this study - FMR)= +30K), Log g = 2.33 (∆log g(This study -FMR)= -0.23), and ξ =1.50 km-s − (∆ ξ (This study - FMR)= +0.3 km-s − ). There are 7elements in common (Fe, O, Mg, Al, Si, Ca and Ti) and the mean difference is ∆A(x)(Thisstudy - FMR)=-0.10 ± µ Leo studies points torather small offsets in stellar parameters ( ∼
30K in T eff , 0.2 dex in log g, 0.3-0.5 km s − ξ , and 0.1 dex in [m/H]) with the conclusion that H-band spectroscopy can be comparedreliably to optical analyses with reasonable accuracies, even at this higher metallicity. β And
This M-giant was analyzed previously by Smith & Lambert (1985) who derived T eff =3800K(from its (V–K) color), log g=1.6, and ξ =2.1 km-s − . The derived effective temperature hereof 3825K is very close to Smith & Lambert, as is the microturbulence of 2.2 km-s − . The sur-face gravity from Smith & Lambert is larger by 0.7 dex, but their value relied on the distance(as it does in this paper), for which they used a luminosity calibration for the Ca II K-linereversal that was tied to a pre-Hipparcos Hyades distance that was too small. Making theestimated log g corrections to the Smith & Lambert (1985) abundances (based on their pub- 14 –lished sensitivities of abundance with log g) for C, N, O, Ti, Fe, and Ni we find a meandifference and standard deviation of ∆A(x) (This paper - Smith & Lambert)=+0.02 ± eff =3850 ± ξ =1.96 km-s − , giving differences of only 25K in T eff , 0.0 inlog g, and +0.24 km-s − in ξ . Their analysis used the spectral region around λ ± ± ± ± HD 199799 is classified as spectral type M2S, which means that it displayed enhancedbands of ZrO in classification spectra. Smith & Lambert (1988) detected the radioactives-process element Tc I in its spectrum, identifying this star as an intrinsic TP-AGB starundergoing third dredge-up. A quantitative abundance analysis was then carried out bySmith & Lambert (1990) and they derived similar stellar parameters to those derived here:the same T eff =3400K, with a slightly lower value of log g=+0.3, compared to +0.5 here.The microturbulent velocity determined here from the H-band spectra is ξ =2.4 km-s − ,which is 0.7 km-s − larger than Smith & Lambert (1990), who used Fe I lines near λ ξ , the abundances from Smith &Lambert (1990) can be compared to those here for Fe, C, N, O, Ti, Cr, Co, and Ni, withthe mean difference being ∆A(x)( Us - SL90)= -0.08 ± C/ C=28, while a value of 27 isobtained here. The derived C/O ratios are the same in the two studies, with C/ O=0.68.The comparisons presented here are between studies that have possibly used somewhatdifferent families of model atmospheres (e.g. MARCS or ATLAS), different spectral regions,and linelists derived from heterogeneous sources. Given the variety of data and analysistechniques, the derived abundances, in all cases, compare well at approximately the 0.10to 0.15 dex level. A glance at Table 7 reveals that such differences fall within this rangeof uncertainty for most elements given modest changes in fundamental model atmosphereparameters. Since the derived stellar parameters themselves rely to some degree on the typeof analysis, the generally good agreement between the derived abundances from the variouspublished sources, along with the expected uncertainties in this study, as presented in Table 15 –7, is encouraging. The gist of the comparisons, since many of the studies are based on opticalspectra, is that H-band spectrscopy can be used to define a precise internal abundance scalethat can also extended to abundances derived from other, more common spectral windows,such as in the optical.
5. Discussion
The abundances from Table 6 and an estimation of their respective uncertainties fromTable 7, along with the stellar parameters from Table 2 represent, to some extent, what AS-PCAP will produce from the APOGEE spectra. The following discussion highlights theseabundances in light of internal red giant stellar evolution, as well as Galactic chemical evo-lution and stellar populations. These discussion topics anticipate what sorts of informationwill be gleaned from the APOGEE targets.
Based on their luminosities and effective temperatures (Figure 2), four of the red giantsstudied here are either first-ascent giants or, perhaps for µ Leo, a clump red giant. HD199799 is in a more advanced stage of stellar evolution where it is undergoing thermal pulsesand third dredge-up on the AGB. The combination of nucleosynthesis and convective mixingin evolved red giants should result in surface abundances of C, C, and N that havebeen altered from their initial (main sequence) values due to CN-cycle H-burning and firstdredge-up.The result of CN-cycle mixing is seen mostly easily in the C/ C ratios which aremuch higher in main-sequence stars, with the Sun having a value of 89. All four of theRGB or clump giants have quite low values, ranging from C/ C ∼ ∼ ⊙ have values of C/ C ∼ C/ C values versus the estimated stellar mass. Also plottedare representative results for open clusters from Gilroy (1989) and Mikolaitis et al. (2012),and the globular cluster M71 from Briley et al. (1995). The cluster red giants have recently 16 –evolved from the cluster turn-offs, which have a well-determined masses. For the lower-massstars included in these cluster studies, evolution along the RGB is rapid ( ∼ yrs) whencompared to main sequence lifetimes ( ∼ -10 yrs) so the red giant mass is similar to theturn-off mass. This can be demonstrated using the evolutionary tracks from Bertelli et al.(2009), where an isochrone with a turn-off mass of 4.0M ⊙ has a red giant tip mass of 4.2M ⊙ ,or a ratio of RGB to TO masses of 1.05. This ratio of masses becomes even closer to 1.0 forlower masses, which include all of the points in Figure 5. The points from Gilroy representthe mean values of C/ C for 19 open clusters with a total of 55 red giants in her sample.The mean carbon isotopic ratios for two other open clusters (with a total of 10 red giantsfrom Cr261 or NGC6253) from Mikolaitis et al. (2012) are also shown (open blue squares).A representative globular cluster, M71, which has a lower main-sequence mass (M ∼ ⊙ )from Briley et al. (1995) is plotted as the open blue triangle and represents the mean oftwo of the CN-weak stars in this cluster. The CN-strong giants that Briley et al. alsostudied are not plotted, as they exhibit slightly lower isotopic ratios but almost certainlyrepresent 2nd generation stars in M71 that were formed from large fractions of the ejectaof 1st generation intermediate-mass red giants and thus are not representative of single-starred giant evolution.The curves in the top panel of Figure 5 represent stellar evolutionary model resultsfrom Charbonnel & Lagarde (2010); the solid black line illustrates their predictions forstandard RGB dredge-up (no “extra-mixing” or rotation included). Clearly, standard RGBfirst dredge-up does not fit the observed trend of C/ C for red giants with M ≤ ⊙ .This poor comparison between theory and observation has led to work in trying to identifyother types of “extra mixing” mechanisms that would lead to lower carbon-isotopic ratios inthe lower-mass red giants. Two such hypothesized mechanisms are thermohaline mixing orthermohaline mixing coupled with rotational mixing, as discussed recently by Charbonnel& Lagarde (2010). The dashed curves present results from Charbonnel & Lagarde also,with the short dashes for thermohaline mixing only and the long dashes for thermohalinemixing plus rotational mixing for models with initial equatorial rotational velocities of 110km-s − . The stellar models that incorporate these particular types of extra-mixing providefair agreement with the observations, although these are not the only types of processes thathave been studied and are used here to merely illustrate the effect of extra mixing. Thetrends found in the clusters also overlap well with the field stars analyzed here, except forHD 199799, which is in a different evolutionary phase than the other red giants. This starhas enriched its surface in C that was synthesized by He-burning as a TP-AGB star. Inthe simplest case of pure C dredge-up, the surface C/ C ratio would increase, which iswhere HD 199799 lies in the top panel of Figure 5 relative to the other red giants. Such aneffect has been noted by both Lambert et al. (1986) and Smith & Lambert (1990), where 17 –these studies find increasing values of C/ C as TP-AGB stars evolve from M to MS to Sto C stars (with increasing values of C/ O).The bottom panel of Figure 5 plots the abundance by number ratio of N(C)/N(N) versusstellar mass, where the carbon abundance is now the total of C + C: for reference the solarratio is 3.80. Nitrogen is also the total abundance, however in the case of the observations, N is ignored as its abundance is very low (and not detected in these red giants) whencompared to N. There is an observed decrease in C/N with increasing mass as predictedby both standard RGB dredge-up and dredge-up plus thermohaline mixing (Charbonnel &Lagarde 2010), which are shown as the solid and dashed curves, respectively, and are fromthe same models as in the top panel. In the case of the C/N ratios, the total carbon andnitrogen abundances are not as strongly influenced by non-standard mixing as the carbonisotopic ratios.Figure 6 combines the C/N ratios versus C/ C ratios for both the observed redgiants and the models from Charbonnel & Lagarde (2010). When displayed this way, thedifferent behavior between C/N versus C/ C is striking and these two values are purelyspectroscopically derived quantities. The trend of C/N with C/ seems to provide asensitive indicator of red giant mixing processes, as well as yields a mass estimate for starson the RGB. Again, HD 199799 stands apart from the other red giants as it has internallyenhanced its C abundance by about a factor of two since it was a RGB star, based on its C/ C and mass, relative to other RGB stars in the top panel of Figure 5. The dashed-dotted line connects to the region in this diagram where HD 199799 would fall if its surface C abundance were reduced by a factor of two, which is a simplified sketch of how TP-AGBstars might evolve across such a diagram from the RGB.The assumption of only CN-cycle mixing taking place in the red giants is tested in Figure7, where the total carbon abundance ( C + C) is plotted versus the nitrogen abundance(here N). The solar values of C and N are also shown, along with the dashed line illustratingscaled solar-abundances for both C and N. The continuous blue curves represent decreasingcarbon and increasing nitrogen abundances from the scaled solar lines, with the constraintthat the number of C + N nuclei is conserved, as in the CN-cycle (with no leakage into theON-cycle). Each curve is scaled in its initial C and N abundances by the stellar value of[Fe/H] for α Boo, β And, δ Oph, and µ Leo. CN-cycle mixing alone is sufficient to representthe derived abundances of C and N with initial abundances having [C + N/Fe] ∼ C enrichment, HD 199799 willbe offset from its CN curve, so no curve is plotted for it. The dashed-dotted line simplyshows the shift in A(C) if HD 199799 has doubled its RGB C abundance. Having nearlythe same metallicity as β And, HD 199799 might be expected to have been near the CN 18 –curve for β And before it began its evolution as a TP-AGB star.The APOGEE H-band wavelength region presents good opportunities for probing manynuclear and mixing processes that occur over the various phases of red giant evolution.
Other than the C and N isotopes, the vast majority of APOGEE red giant targets willnot have altered their surface abundances of the other studied elements, so the remainingelements are dominated by chemical evolution within the various stellar populations of theMilky Way (or its accreted systems). The derived abundances of each of the other elementsare discussed in the light of observed trends of these elements with A(Fe) from other studiesin the literature. The number of published abundances is now enormous; however in thisfirst exploratory study using the APOGEE linelist, a comparison with a small number ofselected literature sources is sufficient. For comparison purposes in this paper a stellarsample representative of the Milky Way stellar populations was constructed using Reddy etal. (2003; thin disk with some thick disk and having 58 stars), Reddy et al. (2006; thickdisk with 176 stars), Johnson (2002; halo with 23 stars), and Fulbright (2002; halo with 178stars).The investigation of elements derived from the APOGEE linelist is via a number offigures, beginning with Figure 8, which presents values of O/Fe, Mg/Fe, and Al/Fe versusA(Fe); the shorthand notation here of x/Fe denotes the logarithmic number ratio of elementx to iron, log[N(x)/N(Fe)], or A(x) - A(Fe). This nomenclature is chosen over the [x/Fe]in the plots as it presents a direct logarithmic number ratio with the solar abundance ratiobeing noted in each panel, thus no reference to the underlying solar abundance is needed, butthe relative scales are equivalent to [x/Fe]. The top panel of Figure 8 shows the behavior ofO/Fe and, as has been well-known for decades, values of O/Fe are elevated in the thick diskand halo stars. Among the field red giants here, the slightly metal-poor objects α Boo and β And both have O/Fe enhanced by ∼ +0.2 – +0.3. The near-solar metallicity M-giant δ Ophhas a near-solar ratio of O/Fe, while the metal-rich giant µ Leo also exhibits a near-solarvalue O/Fe.The behavior of Mg/Fe is shown in the middle panel of Figure 8 and, like oxygen,Mg/Fe ratios are typically elevated in the thick disk and halo populations. The four redgiants closest to solar in A(Fe) all exhibit basically solar Mg/Fe ratios, with α Boo having aslight enhancement of ∼ +0.1 dex, which overlaps the thick disk populations.Aluminum abundances are plotted in the bottom panel of Figure 8 and the behavior of 19 –Al/Fe versus A(Fe) is not so well-defined, with relatively large scatter. Mu Leo is found to bemodestly enhanced in Al/Fe by about +0.10–+0.15 dex, as is α Boo, with an enhancementof ∼ +0.2.The abundances for the rest of the elements studied here are illustrated in Figure 9 (forSi, K, and Ca), Figure 10 (for Ti, V, and Cr), Figure 11 (for Mn, Co, and Ni), and Figure12 (for Cu). In all elements, the general trends found in the general Galactic populationsare recovered in the small sample of field red giants included here. The one exception is Tiin δ Oph and µ Leo, both of which are found to exhibit enhanced values of Ti/Fe of +0.2dex. There is no obvious explanation for this, however the continuing analysis of APOGEEspectra with this linelist wuill reveal if this possible effect appears in significant numbersof other red giants. The results presented and discussed here indicate that H-band spectraanalyzed in LTE with the APOGEE linelist and 1D model atmospheres can provide accurateresults for chemical abundances.
6. Conclusions
Archival high-resolution FTS spectra from λ eff from 3400K to 4540K andlog g from 0.5 to 2.1, covering a significant part of the effective temperature - surface gravityrange being observed in the APOGEE survey.Chemical abundances are determined for 16 elements, along with the minor carbonisotope, C. Comparisons of the abundances derived here with those from previously pub-lished studies of four of the target stars here, mostly from visual-wavelength high-resolutionspectra, find abundance differences of ∼
7. Acknowledgments
Facilities:
NOAO, Mayall 4m (FTS).
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This preprint was prepared with the AAS L A TEX macros v5.2.
23 – . O H O HCNC O ( - ) CN F e I F e I O HCN F e I F e I C O .
51 00 .
51 00 . Fig. 1.— Sample spectra of red giants in the region near the C O (3–0) bandhead. Thesestars span a range in temperature, from T eff = 4550 K ( µ Leo) down to T eff = 3825 K ( β And) and in metallicity, with α Boo having [Fe/H] =-0.52, while µ Leo has [Fe/H]=+0.31dex. These panels illustrate both the quality of the spectra and the variation of molecularand atomic line absorption with primarily T eff . 24 – Fig. 2.— The positions of the field red giants in the Log (L/L ⊙ ) - T eff plane, which is amodified HR-diagram. The continous curves represent stellar evolution tracks for differentmass stars with each panel representing a different heavy-element mass fraction (Z); this isequivalent to overall stellar metallicity with the Sun having Z=0.0145 (Lodders 2010). Themodel tracks are used to set log g for each red giant studied here. 25 –Fig. 3.— The individual Fe I abundances as a function of microturbulent velocity for µ Leo(top panel) and β And (bottom panel). Each curve represents a single Fe I line. Strongerlines are more sensitive to changes in ξ than weaker lines. The value of ξ that yields nosignificant abundance differences as a function of reduced line-strength (log W λ / λ ) is takenas the characteristic microturbulent velocity for the star. The derived microturbulent velocityas displayed here also corresponds to the smallest scatter in A(Fe). 26 – Fig. 4.— A comparison of the observed spectrum near the Mn I line at 15262˚Ain δ Oph (blueopen circles) with synthetic spectra which were computed for Mn abundances separated inintervals of 0.1 dex. This Mn I transition consists of a number of hfs components of differingwavelengths, as shown, with the vertical length of each component proportional to the gf-value of that component transition. Note that Mn I is blended with a CN feature containingboth C N and C N lines, thus the value of C/ C ratio must be included. Both thehfs structure and CN blending illutrates the need for the APOGEE analysis to be based onspectrum synthesis with an accurate linelist. 27 –
Fig. 5.— The behavior of the nuclei C, C, and N with stellar mass on the RGB ascharacterized by C/ C and C/N (where C= C+ C and N= N). The relative abundancesof these nuclei are altered by H-burning via the CN-cycle and convective mixing on the RGBdredges up the processed abundances. The TP-AGB star HD 199799 is identified, as it hasincreased significantly its C abundance from He-burning and third dredge-up and has thusincreased its C/ C relative to when it was on the RGB. 28 –
Fig. 6.— Observed and model predicted behavior of N(C)/N(N) versus C/ C; the modelcurves show so-called standard dredge-up (solid line curve) and a thermohaline mixing modelas a function of stellar mass from Charbonel & Lagarde (2010; dashed line curve). Standarddredge-up does not predict the observed behavior of N(C)/N(N) as a function of C/ C andsome type of additional mixing process is needed. The TP-AGB star HD 199799 is identifiedhere as it has dredged-up enough C to have approximately doubled its abundance relativeto its abundance on the RGB; the dashed-dotted line moves HD 199799 to its approximateposition when it was evolving along the RGB. 29 –
Fig. 7.— Testing if total N(C) + N(N) is conserved, as would be expected for pure CN-cycledredge-up. The dashed line represents metallicity scaled values of A(C) and A(N) with asolar C/N ratio. The blue curves map constant values of C+N scaled to the metallicity(taken to be [Fe/H]) of each star. Clearly, initial (main-sequence) values of [C+N/Fe]=0,followed by simple CN-cycle dredge-up, are adequate representations of what is observed inthese field RGB stars. The TP-AGB star HD 199799 has increased it C abundance byapproximately 0.35 dex and this shift in A(C) is shown by the dashed-dotted line. 30 –Fig. 8.— The abundances of oxygen, magnesium, and aluminium, relative to iron, shownas log(N(X)/N(H)) versus A(Fe). The 5 stars studied here are represented by the filled redcircles. The small blue symbols are results from Reddy et al. (2003; 2006) and the lowermetallicity small majenta symbols are from Fulbright (2002) and Johnson (2003). The solarposition is indicated with the dashed line showing the solar abundance ratio. The abundancesderived here using the APOGEE linelist are well mapped into the general trends observedfor Galatic field populations. The error bars are taken from the standard deviations in Table6. Errors for the ratios are quadrature sums of the element in-question and Fe I values of σ . 31 –Fig. 9.— The abundances of silicon, potassium, and calcium, relative to iron, versus A(Fe).The plotting format and symbols are the same as in Figure 8. 32 –Fig. 10.— The abundances of titanium, vanadium, and chromium, relative to iron, versusA(Fe). The plotting format and symbols are the same as in Figure 8. 33 –Fig. 11.— The abundances of manganese, cobalt, and nickel, relative to iron, versus A(Fe).The plotting format and symbols are the same as in Figure 8. 34 –Fig. 12.— The abundance of copper relative to iron versus A(Fe). The plotting format andsymbols are the same as in Figure 8. 35 –Table 1. Red Giant Standard Stars: Observed Properties Star HR SpT π (mas) a d(pc) J-K Sb M bolc β And 337 M0III 16.4 ± ± ± µ Leo 3905 K2III 26.3 ± ± ± α Boo 5340 K2III 88.8 ± ± ± δ Oph 6056 M0III 19.1 ± ± ± ± ±
92 1.25 -3.64 ± a Parallax from van Leeuwen (2007) b Johnson (1965) transformed to the 2MASS system using Carpenter (2001) c Calculated from M K with bolometric corrections from Bessell et al. (1998)
36 –Table 2. Red Giant Standard Stars: Derived Parameters
Star Log(L/L ⊙ ) M/M ⊙ T eff (K) Log g (cm-s − ) ξ (km-s − ) [Fe/H] a β And 3.11 ± ± ±
75 0.9 ± ± µ Leo 2.68 ± ± ±
50 2.1 ± ± α Boo 2.24 ± ± ±
50 1.7 ± ± δ Oph 2.75 ± ± ±
75 1.2 ± ± ± ± ±
75 0.5 ± ± a [Fe/H] represents the overall metallicity used in the final model atmosphere. Final Fe abun-dances are given in Table 6
37 –Table 3. Fe I Lines used in the Abundance Determinations λ (˚A) χ (eV) log gf α Boo β And δ Oph µ Leo HD 19979915194.492 2.223 -4.779 7.04 7.22 7.37 7.83 7.0515207.526 5.385 +0.080 7.04 7.29 7.40 bl a a a a bl: Feature too blended to use in cooler or very metal-rich red giants.
38 –Table 4. Molecular Lines and Features used for C, C, N, and O Molecular Lines λ –interval (˚A) α Boo β And δ Oph µ Leo HD 199799 C from C O lines (3-0) V-R 15578 – 15586 7.86 8.05 8.27 w a C/ C ratios from C O and C N lines C O(3-0) V-R 15922 – 15926 7 17 14 w a C O(4-1) V-R 16120 – 16125 6 15 11 w a C O(6-3) V-R 16740 – 16747 6 13 11 w a C N 15314 – 15315 w a w a w a
22 w a C N 15354 – 15356 w a w a w a
17 w a O from OH lines (2-0) P N from C N lines (1-2) Q2 41.5 15260. 7.61 8.10 8.23 8.83 8.31(1-2) P2 34.5 15322. 7.68 8.07 8.21 8.88 8.35(1-2) R2 56.5 15397. 7.65 8.06 8.26 8.63 8.28(0-1) R1 68.5 15332. 7.63 7.95 bl b b (0-1) Q2 59.5 15447. 7.64 7.98 8.28 8.73 bl b (0-1) Q1 60.5 15466. 7.65 8.06 8.23 8.51 bl b (1-2) P2 38.5 15472. 7.71 8.13 8.11 8.68 8.25(0-1) P1 51.5 15482. 7.56 8.03 8.15 8.68 8.25 a w: Feature too weak to use. b bl: Feature too blended by nearby lines.
39 –Table 5. Atomic Lines used and Derived Abundance
Element λ (˚ A ) χ (eV) log gf α Boo β And δ Oph µ Leo HD 199799
Mg I a a a s a b b b b c b Al I a b Si I b bl b bl b w c c c b bl b b K I b bl b bl b Ca I
Ti I b b b b V I b Cr I
Mn I b b Co I
Ni I b bl b b b Cu I a s: Feature too strong to use. b bl: Feature too blended by nearby lines.
40 – c w: Feature too weak.
41 –Table 6. Chemical Abundances
Element α Boo β And δ Oph µ Leo HD 199799Fe 6.98 ± ± ± ± ± C 7.96 ± ± ± ± ± C/ C 6.3 ± ± ± ± ± N 7.64 ± ± ± ± ± O 8.64 ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ±
42 –Table 7. Abundance Sensitivity: T eff =4000K, Log g=1.3, ξ =2.0 km-s − , [m/H]=0.0 Species ∂A/∂T (+50K) ∂A/∂G (+0.2 dex) ∂A/∂ξ (+0.2 km-s − ) ∂A/∂m∂A/∂m