Spectroscopic Abundances and Membership in the Wolf 630 Moving Group
aa r X i v : . [ a s t r o - ph . S R ] M a y Spectroscopic Abundances and Membership in the Wolf 630Moving Group
Eric J. Bubar
Department of Physics and Astronomy, University of Rochester, P.O. Box 270171,Rochester, NY 14627-0171 [email protected]
Jeremy R. King
Department of Physics and Astronomy, Clemson University, Clemson, SC 29630-0978 [email protected]
ABSTRACT
The concept of kinematic assemblages evolving from dispersed stellar clustershas remained contentious since Eggen’s initial formulation of moving groups inthe 1960’s. With high quality parallaxes from the Hipparcos space astrometrymission, distance measurements for thousands of nearby, seemingly isolated starsare currently available. With these distances, a high resolution spectroscopicabundance analysis can be brought to bear on the alleged members of thesemoving groups. If a structure is a relic of an open cluster, the members canbe expected to be monolithic in age and abundance inasmuch as homogeneityis observed in young open clusters. In this work we have examined 34 putativemembers of the proposed Wolf 630 moving group using high resolution stellarspectroscopy. The stars of the sample have been chemically tagged to determineabundance homogeneity and confirm the existence of a homogeneous subsampleof 19 stars. Fitting the homogeneous subsample with Yale-Yonsei isochronesyields a single evolutionary sequence of ∼ ± < [Fe/H] > =-0.01 ± ± λ Subject headings: stars: abundances - stars: kinematics and dynamics - stars:late-type 2 –
1. INTRODUCTION
A major goal of modern astronomy is to piece together the dynamic and chemicalevolution of the Galactic disk. To this end, one of the principle approaches for probingthe disk has been to study open clusters. Clusters are valuable astrophysical tools as theyshare common distances, common ages and common initial chemical abundances. With thedisk richly populated by both field stars and open clusters, and considering that clusters arerelatively well studied, the logical step in piecing together a more complete picture of thechemical and dynamical history of the disk is to study field stars.In recent years, the advent of large surveys such as
HIPPARCOS (Perryman & ESA1997) has yielded precise parallaxes for thousands of nearby field stars, and in doing so,provided the necessary tools for investigating the field. In particular, studies of the velocitydistributions of disk field stars in the solar neighborhood have identified stellar overdensitiesin kinematic phase space (Skuljan et al. (1999)). The potential application of these velocitystructures, commonly referred to as moving groups, was first identified by Eggen (1958)who considered these assemblages to be relic structures of dissolved open clusters. In thisparadigm, a moving group is essentially a spatially unassociated open cluster; thereforeit should possess some of the same characteristics that make open clusters such valuableastrophysical tools (common ages and common initial chemical abundances) and similartechniques that are useful for studying open clusters could be applied.Relatively little work has been done to explore the reality of smaller moving groups(kinematic assemblages of ∼
100 stars) as dissolved open clusters and their use in chem-ically tagging the galactic disk, with two notable exceptions: the Ursa Major Group andthe HR1614 Moving Group. Soderblom & Mayor (1993) examined the Ursa Major movinggroup and utilized age information inferred from chromospheric emission to constrain groupmembership in UMa. While this study did not utilize chemical tagging to constrain groupmembership, it did illustrate the viability of moving groups as dissolved populations of openclusters. King et al. (2003) and King & Schuler (2005) revisited the membership of the UMagroup, using new and extant abundances. They used the results to constrain membership inthe UMa group, showed the members to be chemically homogeneous, and noticed overexci-tation/overionization effects in the cooler field star members of the group, similar to thoseobserved in young ( <
500 Myr) cool open cluster dwarfs (Schuler et al. (2003), Schuler et al.(2004)). The first in depth application of chemical tagging to constrain moving group mem-bership was by De Silva et al. (2007), who derived abundances for various elements for theHR 1614 moving group. They found that for their 18 star sample, 14 stars were metal-rich([Fe/H] ≥ σ =0.03) leading to the conclusion that the HR 1614 moving group,with its distinct kinematics and distinctly super-solar chemical abundances, was a remnant 3 –of a dissolved open cluster.In the field of moving group populations, the classical Wolf 630 moving group is anintriguing target. The first identification of the Wolf 630 moving group was made by Eggen(1965) who noted that several K and M dwarfs and giants in the solar neighborhood appearedto have similar space motions to that of the multiple star system Wolf 630 ((U,V,W)=(23,-33, 18) kms − ). These kinematics, distinctive of membership in an old disk population,placed the stars in a relatively sparsely populated region of kinematic phase space (Eggen1969). Eggen also noted that the color magnitude diagram for the K and M dwarfs andgiants with kinematics similar to those of Wolf 630 appeared to trace an evolutionary se-quence similar to the old ( ∼ uvby- β photometry.Variations in the δ [m ] index were found to be comparable to the Hyades, Praesepe, and theComa Berenices clusters, implying chemical homogeneity.Tuominen & Vilhu (1979) studied the chemical composition of five field giant starsthat were alleged members of Wolf 630 using high dispersion coud´e spectra described inTuominen & Vilhu (1979). Employing a curve of growth approach and measured equivalentwidths, they found that three stars appeared to be chemically homogeneous with an overallmetallicity for Wolf 630 of [Fe/H] ∼ +0.23. However, it must be noted that their abundanceswere not measured with respect to the Sun, but are instead quoted with respect to a stan-dard star of presumed solar metallicity (HD 197989), which has since been determined to bea K0III. While they derived a metallicity of 0.00 for their reference star, literature determi-nations suggest a value of -0.24. This would lower the average metallicity for the group to[Fe/H] ∼ -0.02.McDonald & Hearnshaw (1983) revisited the membership of the Wolf 630 moving groupby recreating the approach presumably utilized by Eggen (1965) to find his original Wolfsample. In summary, they calculate the parallax that yields a V velocity for each groupcandidate equal to the assumed group velocity of V=-32.8 ± − . The final absolutemagnitudes they report assume these parallaxes. Typical uncertainties in their absolutemagnitudes appear to be between 0.2-0.4 magnitudes, larger than magnitude uncertaintiesobtainable with precise parallax information currently available from Hipparcos. The color-magnitude diagram assuming these M V values was compared to the scatter of apparentmembers with the observed scatter in the old open cluster M67. They concluded that either(1) the intrinsic scatter in the Wolf 630 moving group color-magnitude diagram was greater 4 –than that of M67, or (2) the errors in radial velocities and/or proper motions they utilizedmust have been underestimated by a factor of 2.4 or (3) many of the stars in their samplewere, in fact, non-members.Taylor (1994) examined metallicities from “published values of [Fe/H] from diverse pa-pers” of 40 members of the Wolf 630 group. His sample contains 26 % of Eggen’s originalobjects (Eggen 1969). He concluded that metallicity dispersions within his sample were toogreat for meaningful conclusions about the existence or non-existence of a genuine, chemi-cally distinct Wolf 630 moving group. This suggests the need to obtain high quality [Fe/H]determinations with minimal uncertainties in testing for chemical uniqueness in a putativeWolf 630 sample.The analysis of solar neighborhood Hipparcos data by Skuljan et al. (1999) indicatesa kinematic rediscovery of the Wolf 630 group. Their figure 10, showing the UV velocitydistribution for 3561 late type dwarfs in the solar neighborhood presents a clear overdensityof stars near the position of Wolf 630. Furthermore, this structure appears to be distinctlyseparated from any other known moving groups or stellar streams. This provides compellingevidence that Wolf 630 is a real kinematic structure. The question to be asked is if thiskinematic structure is composed of stars with a common origin?Despite the distinctive kinematics exhibited by the Wolf 630 moving group when exam-ined with updated Hipparcos parallaxes, it has not been specifically targeted in an abundancestudy which makes use of the modern astrometric and spectroscopic data. This is remedied inthis paper, where accurate parallaxes and photometry from the updated HIPPARCOS datareduction (van Leeuwen 2007) coupled with high precision radial velocities from CORAVEL(Nordstr¨om et al. (2004) and references therein) allow for developing a Wolf 630 sample withinternally consistent distances and absolute magnitudes, thereby removing the uncertaintiesfaced by McDonald & Hearnshaw (1983). Furthermore, our uniform high resolution spec-troscopic study of Wolf 630 moving group candidate members provides a single, consistentset of metallicites with low internal uncertainty to test chemical homogeneity in the group,removing the largest source of uncertainty from Taylor (1994).
2. DATA, OBSERVATIONS AND ANALYSIS2.1. Literature Data
The 34 stars in this sample, listed in Table 1, were previously identified as members ofthe Wolf 630 group (Eggen (1969), McDonald & Hearnshaw (1983)) according to their
UVW kinematics. In this study, we use updated parallaxes and proper motions from the latest 5 –reduction of Hipparcos data (van Leeuwen 2007). Precision radial velocities were taken fromthe compilation of Nordstr¨om et al. (2004). Visible band photometry ( B , V , B T ycho , V T ycho )was taken from the
HIPPARCOS catalogue (Perryman & ESA 1997). Near infrared J , H and K photometry was taken from the 2MASS Catalog (Cutri et al. 2003). We determined galactic
UVW kinematics from the proper motions, parallaxes and radialvelocities using a modified version of the code of Johnson & Soderblom (1987). Here, theU velocity is positive towards the Galactic center, the V velocity is positive in the directionof Galactic rotation and the W velocity is positive in the direction of the North GalacticPole (NGP). The relevant parameters for determination of these kinematics are presented inTable 1.
Spectroscopy of the sample was obtained in March 2007 and November 2008 with theKPNO 4 meter Mayall telescope, the echelle spectrograph with grating 58.5-63 and a T2KB2048X2048 CCD detector. The slit width of ∼ ∼ echelle package of IRAF . These include bias correction, flat-fielding, scattered light correction, order extraction, and wavelength calibration. Samplespectra are presented in Figure 1. Spectroscopic physical parameters are typically determined by enforcing balance con-straints on abundances derived from lines of Fe, which has a plethora of suitable neutral(Fe I ) and ionized (Fe I I) features in the optical. We compiled low excitation potential( χ < I and Fe II lines from Thevenin (1990), the VIENNA Atomic Line IRAF is distributed by the National Optical Astronomy Observatories, which are operated by the As-sociation of Universities for Research in Astronomy, Inc., under cooperative agreement with the NationalScience Foundation. I lines and 11 Fe II lines was formed. These linelists are presented in Table 2. Theequivalent widths listed are for measurements in a high resolution solar spectrum.Linelists for other elements of interest have also been compiled from multiple sources(Thevenin (1990), King et al. (1998), De Silva et al. (2006)). These elements include Li, Na,Al, Ba, a selection of α elements (O, Mg, Si, Ca, Ti I and Ti II) and a selection of Fe peakelements (Cr, Mn and Ni). The lines are also given in Table 2. Equivalent widths are againfor measurements in the high resolution solar spectrum. The equivalent widths that weremeasurable for each individual star are given in Table 4, with corresponding abundancesderived from each equivalent width. Equivalent widths for the lines of interest were measured in each star and in a highresolution solar spectrum using the spectral analysis tool SPECTRE (Fitzpatrick & Sneden1987). Final abundances were obtained from equivalent widths through use of the MOOGLTE spectral analysis tool (Sneden 1973) with an input Kurucz model atmosphere character-ized by the four fundamental physical parameters: temperature, surface gravity, microturbu-lent velocity ( ζ ) and metallicity. Unless noted otherwise all abundances are differential withrespect to the Sun and are presented in the standard bracket notation ([X/H]=log( N ( X ) N ( H ) ) ∗ − log( ( N ( X ) N ( H ) ) ⊙ where logN(H) ≡ The color- T eff -[Fe/H] calibrations of Ram´ırez & Mel´endez (2005) were used to determinephotometric temperatures from Johnson B − V , Tycho B T − V T , Johnson/2MASS V − J , V − H and V − K . The color indices for 8 stars were outside of the calibrated ranges; 7 –consequently photometric temperatures were not derived. Uncertainties in the photometrictemperatures were conservatively taken as the standard deviation of the temperatures derivedfrom each of the respective colors. With the availability of high quality Hipparcos Parallaxes,physical surface gravities were calculated from: log gg ⊙ = log MM ⊙ + 4 log T eff T eff, ⊙ + 0 . V o + 0 . B.C. + 2 logπ + 0 . π is the paral-lax. Initial microturbulent velocities were found from the calibrations of Allende Prieto et al.(2004). These photometric parameters provided the initial guesses for physical parameterswhen deriving the final spectroscopic values. Additionally, the photometric calibrations pro-vided reasonable estimates to compare to spectroscopically derived results. Using the up-dated calibrations of Casagrande et al. (2010) does not change the results described herein. The refined Fe linelists discussed above acted as target lists for each of the stars in thesample. The typical star contained ∼
80 of the 145 good Fe I lines that were measurablein the solar spectrum. Several of the stars showed correlations between Fe I excitation po-tential and reduced equivalent width. If ignored, such correlations can be imposed onto thetemperatures and microturbulent velocities, resulting in non-unique solutions of physical pa-rameters. Consequently, two linelists for the Fe I lines were formed for each star; a correlatedand an uncorrelated sample. Final basic physical parameters for the sample were derivedusing a modification to the standard techniques of Fe excitation/ionization/line strengthbalance. In all the approaches described below, a differential analysis was used where thesame lines were measured in a solar spectrum and in the stellar spectra. Final abundanceswere then determined by subtracting the solar abundance from the stellar abundance in aline by line fashion.The first technique utilized the uncorrelated line sample and proceeded as follows: tem-peratures of input model atmospheres were adjusted to remove any correlation in solar-normalized abundances with respect to excitation potential; ζ is adjusted to remove anycorrelation with line strength and log g is adjusted until the mean abundance from Fe I I linesmatches the abundance from Fe II lines. This approach required simultaneously adjustingtemperatures, surface gravities, metallicities and microturbulent velocities to converge to acommon solution. Use of the uncorrelated line sample, as described above, is necessary toensure a unique solution. This approach will be referred to as the “classical” approach. 8 –The second approach used the correlated line sample and the Hipparcos-based physicalsurface gravities. The Fe II abundances are primarily set by this gravity. The temperaturewas adjusted to force the mean abundance from Fe I lines to match that from Fe I I lines. Themicroturbulent velocity was adjusted until the abundance from Fe I lines had no dependenceon reduced equivalent width. The advantage of this approach is that it does not requiresimultaneous solutions requiring excitation balance and equivalent width balance, allowinguse of a full correlated line sample. This approach will be referred to as the “physical surfacegravity” approach.When comparing results from the classical and physical surface gravity approaches it wasapparent that the microturbulent velocities were nearly identical ( δζ ≈ ± .
04 km s − . Thusour final spectroscopic parameters were determined as follows. The microturbulent velocitiesfrom the “classical” approach and the “physical surface gravity” approach were averaged toyield a final value. The correlated line sample was used to determine the temperature andsurface gravity using excitation/ionization balance. For the remainder of the work, theresults from this approach were used for the physical parameters of these 30 stars. Theremaining 4 stars would not converge to an acceptable solution and the following alternativeapproach was developed.The coolest stars in the sample (HIP105341-dwarf, HIP114155-giant and HIP5027-dwarf) had an insufficient number of well-measured Fe II lines for accurately determining thesurface gravity spectroscopically. Consequently, Hipparcos-based physical surface gravitieswere used to set the gravity, and the temperature and microturbulence were iterated to elim-inate correlations in [Fe I /H] versus excitation potential and versus the reduced equivalentwidth. This is the “physical surface gravity” approach.Finally, one of the stars in the sample (HIP 5027) had a microturbulence correlationwhich could not be removed without utilizing unreasonable surface gravities. For this star,the surface gravity was set based on Yale-Yonsei isochrones (Demarque et al. 2004). Themicroturbulent velocity was set to zero and the temperature was determined from excitationbalance.The final basic physical parameters (T Spec , log g , microturbulent velocity ( ξ ) and [Fe/H])are presented in Table 3 and final abundances are summarized in Table 5. For the interestedreader, we also provide plots of all abundances ([X/H]) versus [Fe/H] in an appendix. 9 – Abundances have been derived for lithium using spectral synthesis. We use the synth driver of MOOG to synthesize a spectrum of the lithium line at λ =6707.79 ˚A with anupdated version of the linelist from King et al. (1997). Appropriate smoothing factors weredetermined by measuring clean, weak lines in the lithium region. The lithium abundancewas varied until a best fit is obtained from visual inspection. A sample synthesis is presentedin Figure 5.Uncertainties in lithium abundances have been determined by examining the changein Li abundance in syntheses with arbitrary changes in physical parameters of ∆T=150 K,∆log g =0.12 cm s − and ∆ ξ =0.60 km s − , and adding the resultant abundance differencesin quadrature. Oxygen abundances for many stars have been derived from the near-IR λ λ λ blends driver of MOOG, following Schuler et al. (2006). TheNi abundance utilized to account for blending is the mean value derived from the EWs ofNi I lines in our sample.A possible CN feature at 6300.265 ˚A and two at 6300.482 ˚A with log(gf) values of -2.70,-2.24 and -2.17 are claimed by Davis & Phillips (1963). In order to explore these blends,multiple syntheses of the λ > ± − (Gray1981) and a rotational broadening characterized by vsin(i)=2.4 ± − (Gray 1981)with a limb darkening coefficient of 0.9 (from Gray (2005)). With this smoothing, the spec-trum for Arcturus in the forbidden oxygen region was fit by increasing the CN features gfvalues by ∼ The uncertainties in experimental and theoretical log(gf) values (likely at least 0.1 dex)can be a significant source of error; however, by performing a line-by-line differential analysiswith respect to the Sun, uncertainties due to transition probabilities are eliminated to firstorder.Here, then, it is the uncertainty in physical parameters that underlie the uncertaintiesin the abundances. Errors in the temperature were determined by adjusting the temperaturesolution until the correlation between [Fe/H] and excitation potential (excitation balance)reached a 1- σ linear correlation coefficient for the given number of lines. The uncertaintyin microturbulent velocity was determined in the same manner, by adjusting the microtur-bulence until the linear correlation coefficient for [Fe/H] versus equivalent width (equivalentwidth balance) resulted in a 1- σ deviation. For HIP 5027, which would not converge to a 11 –unique solution for microturbulence, an uncertainty in microturbulence of 0.20 kms − wasadopted.For the cases where the physical surface gravity was utilized, the uncertainty was es-timated by propagating the uncertainties in the temperature, mass, apparent magnitude,parallax and bolometric corrections. The uncertainties in the spectroscopically determinedsurface gravities required a deeper treatment. Since gravity is calculated by eliminating thedifference in iron abundance derived from [Fe I /H] and [Fe II /H], the uncertainty in surfacegravity is related to the quadratic sum of the the uncertainties in [Fe I /H] and [Fe I I/H].These abundances, in turn, have sensitivities that depend on the basic physical parameters.Proper uncertainty calculations, therefore, require an iterative procedure. The errors in [FeI/H] and [Fe II/H] are a combination of the measurement uncertainties and the uncertaintiesin the physical parameters. The line measurement uncertainties in Fe I and Fe II were esti-mated as the standard deviation of the abundances from all Fe I and Fe II lines, respectively.Abundance sensitivities for arbitrary changes in temperature ( ±
150 K), surface gravity ( ± ± − ) were determined by adjusting each parameterindividually and recording the resultant difference in abundance. To determine abundanceuncertainties the abundance differences must be properly normalized by the respective pa-rameter’s uncertainty. For example, in HIP3455 the total temperature uncertainty was foundto be 35 K. The final abundance uncertainty introduced by the arbitrary temperature changewould, therefore, be equal to the difference in abundance multiplied by K K , where 35 K isthe temperature uncertainty and 150 K is the arbitrary temperature change introduced todetermine the temperature sensitivity. For the first calculation the uncertainties in temper-ature and microturbulent velocity were determined as above and the uncertainty in surfacegravity was unknown; consequently its contribution to abundance uncertainty was initiallyignored. Adding the measurement errors in [Fe I/H] and [Fe II/H] in quadrature with thephysical parameter abundance uncertainties from temperature and microturbulence yields afirst estimate for the uncertainty in the surface gravity. This gravity uncertainty can then beadded in quadrature to the line measurement uncertainty, the temperature uncertainty andthe microturbulent uncertainty to yield a final uncertainty for the surface gravity. The sur-face gravity in the model atmosphere was adjusted until the difference in abundance between[Fe I/H] and [Fe II/H] was equal to their quadrature added uncertainties. The differencebetween this gravity and the spectroscopically derived gravity provides the final uncertaintyin surface gravity.Uncertainties in abundances were found by introducing arbitrary changes in T, micro-turbulence and surface gravity (∆ T=150 K, ∆ ξ =0.60 kms − , and ∆log g =0.12 cm s − ),normalized by the respective parameter uncertainties. The uncertainties introduced by eachof these parameter changes was added in quadrature to obtain total parameter-based uncer- 12 –tainties. Measurement uncertainties were taken as the uncertainty in the weighted mean forall lines of a given element. For elements with only a single line available, the standard devi-ation of all Fe I abundances was utilized as an estimate of the line measurement uncertainty.The final uncertainties in the abundances were determined by adding the parameter-basedabundance uncertainties with the measurement uncertainties in quadrature.A sample table of the normalized parameter changes and their final resultant [Fe I /H]errors on a given star is presented in Table 6. The temperatures for the stars in the sample were determined from photometric cali-brations as well as through spectroscopic excitation balance. In Figure 2 the spectroscopictemperature is plotted versus the photometric temperature. The line represents perfectagreement between the two temperatures. It can clearly be seen that the temperatures fromthe two techniques are equivalent within their respective uncertainties. There is a slightindication that spectroscopic temperatures may be systematically higher, with 66 % of thestars lying above the line, however the effects on the abundance analysis are negligible anddo not change any conclusions.
The surface gravity was determined from Hipparcos data (i.e. physical surface gravities)and spectroscopically via ionization balance. In Figure 2, the spectroscopic surface gravity isplotted versus the physical surface gravity. The line shows the trend for the gravities beingequal. Within their respective uncertainties, the surface gravities are equal.
3. RESULTS AND DISCUSSION
The primary goal of the paper is to determine if the kinematically defined Wolf 630Moving Group represents a stellar population of a single age and chemical composition.The sample stars have been plotted in the HR diagram (Figure 7) to determine if they arecoincident with a single evolutionary sequence. The sequence traced by the majority ofstars coincides with a Yale-Yonsei isochrone (Demarque et al. 2004) of 2.7 ± blends driver (O). In order to visually present the abundance results,the metallicity distribution of the entire sample is presented in the form of a “smoothedhistogram” in Figure 3. This distribution has been generated by characterizing each starwith a gaussian centered on its mean [Fe/H] with standard deviation equal to the [Fe/H]uncertainty. The distributions are summed to yield a final smoothed histogram and havebeen renormalized to a unit area. In this manner, the distributions include uncertainties inabundances, making them useful for a visual examination of the complete sample to discernif any stars yield abundances that deviate from the sample as a whole. The distribution isclearly not unimodal or symmetric. It is dominated by a near-solar metallicity peak andtwo smaller peaks at [Fe/H] ∼ -0.50 and [Fe/H] ∼ +0.30. It is clear that our Wolf 630 movinggroup sample is not characterized by a single chemical composition. While our entire sample cannot be characterized by a single chemical abundance, wecan investigate whether there is a dominant subsample having common abundances and age.This is done by eliminating stars that are clearly outliers, using arguments based on extremeabundances, evolutionary state (inferred from HR diagram positions, lithium abundance,chromospheric acitivities and surface gravities) or a combination thereof. These memberswill be classified as “unlikely” members of a dominant homogeneous group. In this waywe can, for example, establish a subsample that is characterized by a dominant [Fe/H],if it exists. Stars with such an [Fe/H] will be classified as either “possible” or “likely”members of a chemically homogeneous, isochronal population having common kinematics.The final distinctions between “possible” and “likely” will be made based on evolutionarystatus and additional abundance information inferred from lithium, alpha elements and ironpeak elements. Particular interest is paid to the iron abundance, [Fe/H], as it is consideredthe most well determined abundance, primarily due to the quality and size of the Fe linesample.The quantitative constraint adopted for determining chemical homogeneity was to re- 14 –quire that a star’s abundance, within its uncertainty, rest within a metallicity band centeredon the weighted mean abundance of stars in the sample. The half-width of this band wasconservatively taken to be 3 times the uncertainty in the weighted mean. This approach wasfollowed in an iterative fashion where whenever a star was determined to be an “unlikely”member of a dominant chemical group it was removed from the sample and a new weightedmean and band size was found. In this manner, a common abundance for the sample wasconverged to for each element (except Lithium and Oxygen). Examples of the band plotsfor [Fe/H] versus T eff is given in Figure 4, where [Fe/H] is plotted versus temperature. Thesolid line gives the weighted mean [Fe/H] while the dotted lines give the 3- σ uncertaintiesin this mean, i.e. the abundance band.This visual analysis from examining the abundance distributions served as a guide foridentifying the clearly unlikely members. Abundance information alone was used to constraingiant star membership in a dominant chemical group, as robust discriminants of age areunavailable. Many of the dwarfs lay above the main sequence, leading to the question of ifthey might be pre-or-post main sequence objects. Consequently a diagnostic was needed toconstrain evolutionary status for these dwarf and subgiant stars. The full analysis, therefore,examined each star individually, utilizing abundances and information on evolutionary status(inferred from chromospheric activities, isochrone ages and surface gravities) to classify eachstar in its appropriate category (unlikely, possible or likely).Figure 6 shows the absolute lithium abundance versus effective temperature for the litle-evolved stars in our sample and for a sample of dwarf stars in the Pleiades, Hyades, NGC752and M67. The lithium abundances of the sample stars are plotted with each cluster: filledhexagons are dwarfs, filled triangles are upper limits for dwarfs, open hexagons are subgiants(as inferred from HR-diagram positions and apparently low levels of chromospheric activity)and open triangles are upper limits for subgiants. Accepted ages are given for each of therespective clusters, with the Pleiades trend being used as a baseline to indicate that a star islikely to be young (i.e. if a star has a lithium abundance which rests in the Pleiades lithiumabundance trend it is likely a young star). With the considerations above, the 34 stars in the sample have been classified as unlikely,possible and likely members of a common chemical, temporal and kinematic assemblage.There were a total of 13 stars removed from group membership due to classification asunlikely members. If the remaining 21 stars classified as possible and likely are considered torepresent a chemically distinct group, then out of the original kinematically defined sample, 15 – ∼
60% remain members of a kinematically and chemically related group with a common 2-3Gyr age insofar as we can tell.The final evolutionary sequence traced by the possible and likely members is presentedin Figure 7, with possible members plotted in red and likely members plotted in green.The group is reasonably well traced by an evolutionary sequence of ∼ ± − , respectively. In the final subsample of group mem-bers, U RMS =25.21 kms − and V RMS =35.8 kms − , therefore the kinematic identity has notbeen significantly altered by the requirement of chemical and temporal coherence to establishgroup membership, which points to the necessity to utilize criteria other than kinematics torobustely link members of moving groups.The weighted mean abundances of the final possible and likely members of a dominantchemical group are presented in Table 10. The quoted errors are the uncertainties in theweighted mean. In order to explore the homogeneity of our samples a reduced chi-squaredstatistic is presented for each element assuming a constant mean abundance. Performingthis test for [Fe/H] for warm stars (T ≥ χ ν of 1.303. For a set of 7 Pleaides stars from Schuler et al.(2003), the reduced chi-squared in [Fe/H] is 1.818. Note that the cool stars were removedfrom the calculation as they are believed to be impacted by overexcitation/ionization ef-fects. From these chi-squared values, we estimate the Hyades and Pleiades are chemicallyhomogeneous with a roughly 2-sigma significance. With these open clusters assumed to bechemically homogeneous, an approximate reduced chi-squared of ≤
2, therefore, provides arough quantitative indication of homogeneity. The χ ν is presented for the full sample of 34stars ( χ ν all ), the final sample of 21 possible and likely group members ( χ ν group ) and the11 likely members ( χ ν likely ). First, the very large χ nu for the full sample confirms that theinitial kinematically defined sample of alleged Wolf 630 members is clearly not chemicallymonolithic. The decrease in reduced chi-squared between the full sample and the chemically 16 –distinct subsample demonstrates that chemically discrepant stars have been removed. Evenin the likely subsample the χ ν values remain uncomfortably large for Na and Al. Discussionof these discrepancies is reserved for a later section.Considering the reduced chi-squared for other homogeneous open cluster samples iscomparable to the reduced chi-squared for the possible and likely members of the sampleacross multiple elements, the chosen sample is considered to represent a chemically consis-tent group with a weighted average metallicity of [Fe/H]=-0.01 ± U sing precise chemical tagging of the 34 star sample of the Wolf movinggroup, a single evolutionary sequence of 2.7 ± ± We present additional results here that illustrate the application of moving group fieldstar members in exploring stellar and chemical evolution in the Galactic disk.
The abundances of Na and Al appear to be enhanced for some of the stars in thesample. Similar enhancements have been observed in many open clusters. Most recently, ananalysis of abundances in the Hyades cluster found abundance enhancements in Na and Alof 0.2-0.5 dex in giant stars when compared with dwarfs (Schuler et al. (2009)) in line withobservations of giant stars in old open clusters (Friel et al. (2005), Jacobson et al. (2008)).These enhancements can be compared to those observed in the group members of this work.Plots of [Na/Fe] (top panel) and [Al/Fe] (bottom panel) versus surface gravity arepresented in Figure 9. For the members of the group, the Na and Al enhancements arerelatively modest, as seen in a relatively slight upward shift in abundances between dwarfsand subgiants. The giant abundances, in general, can be brought into agreement with dwarfabundances with downward revisions of 0.1-0.2 dex, consistent with NLTE corrections foundin field clump giants with surface gravities down to log g =2.10 (Mishenina et al. (2006)).The single star which has greatly enhanced [Na/Fe] and [Al/Fe], HIP 114155, is an evolved,metal poor red giant with enrichments of 0.53 dex and 0.51 dex, comparable to those foundby Schuler et al. (2009). According to the NLTE correction table of Takeda et al. (2003),the recommended NLTE correction is at most -0.10, although the calculations performed donot extend below a temperature of 4500 K. Gratton et al. (1999) performed an extensive 17 –set of NLTE corrections for Na, and based on their results, there is a recommended NLTEcorrection of ∼ ≈ -2.00). Further NLTE calculations for cool, moderately low metallicity giant likeHIP114155 are needed to determine whether the enhanced abundances in this star are aresult of NLTE effects.The other points of interest in Figure 9 are the two dwarfs with the greatest surfacegravities ([Na/Fe]=-0.38 in HIP105341 and [Na/Fe]=-0.33 in HIP5027). Closer inspectionshows that these are the two coolest dwarfs in the sample, perhaps pointing to overexcita-tion/ionization as a culprit for decreased abundances, similar to overexcitation/ionization ef-fects observed in cool open cluster dwarfs (Schuler et al. (2003), Yong et al. (2004), King & Schuler(2005) and Schuler et al. (2006)).Similar effects are not apparent for [Al/Fe]. A single Na line was measurable with arelatively low excitation potential of 2.10 eV, while two Al lines of 3.14 eV and 4.02 eV wereused. Additionally, the ionization potential of Al is ∼ I and Fe I I. I and Fe I I Abundances
In order to more closely examine the possible effects of overexcitation and overionizationfor the sample, abundances have been derived from Fe I and Fe II lines using physical surfacegravities (spectroscopic gravities are unsuitable for this purpose since ionization balanceforces agreement between abundances of Fe I and Fe II ). Refer to Figure 10 where thedifference in abundances between ionized and neutral Fe are plotted versus temperature. Forstars warmer than 4500 K the general trend reveals no overionization within the uncertainties.The same two coolest dwarfs which evince unusually low [Na/Fe], show large degrees of Feoverionization,Source of overionization in cool dwarfs are not well-understood, however, one possibleexplanation is that the stars are active young dwarfs and, thus, heavily spotted. Recent worksuggests that heavily spotted stars have radii which are “puffed” compared to standard stellarmodels (Torres & Ribas (2002), Morales et al. (2008). An increased radius would decreasethe surface gravity of the star compared to unspotted analogs, which would result in increased 18 –Fe II line strengths via overionization. In order to explore the viability of this explanation, theradius that corresponds to the surface gravity needed to eliminate the abundance differencebetween [Fe II /H] and [Fe I /H] was determined for HIP5027. A surface gravity of 3.57 wasfound to produce agreement between abundances from Fe I and Fe II , holding temperatureand microturbulence constant. From Yale-Yonsei isochrones, a mass of 0.66 M ⊙ is assumed.The radius for this gravity is R=2.19 R ⊙ . The radius corresponding to this mass and thephysical surface gravity of log g =4.70 is R=0.60 R ⊙ . From Morales et al. (2008) an upperlimit that can be expected for radius changes in this “spotted” regime is ∼ Abundances for the λ λ λ λ λ I feature at 7774.00 ˚A. While thenature of any blending for the reddest feature (7775 ˚A) is unclear, visual inspection of thespectral line reveals a slight asymmetry, possibly indicating a blend. The distinct increasein [O/H] abundances derived from the red features of the triplet as a function of decreasingtemperature suggest that only the blue line (7771.1 ˚A) of the triplet should be used foroxygen abundance determinations in cooler stars.In order to test the possibility of an Fe blend as discussed above, two cool stars of thesample with no measurable oxygen abundances (HIP5027 and HIP105341) were examinedto see if they showed any indications of an Fe blending feature near 7774 ˚A. In HIP5027 apossible detection of a feature at 7774 ˚A was found to have a measured equivalent width ofroughly 6.0 m˚A. This strength is not inconsistent with the expected contribution required 19 –from two nearby Fe I features at 7773.979 ˚A and 7774.06 ˚A for the derived Fe abundance.Neglecting the two red triplet lines in the cool dwarfs, the [O/H] trend of the ourdwarf sample is plotted along with the Pleiades trend from Schuler et al. (2004) (where[Fe/H]=0.00 was assumed to calculate [O/Fe]), and the Hyades trend of Schuler et al. (2006)(where [Fe/H]=+0.13 was assumed to calculate [O/Fe]) in Figure 12. Using λ eff ≤ ∼
120 Myr old Pleiades appeared to be steeper than that in the ∼
625 Myr oldHyades, perhaps pointing to an age-related effect whereby [O/H] enhancements in coolerstars decrease as a function of increasing age.Our field dwarfs do not show a drastic increase in abundance as a function of decreasingtemperature. The single star that appears to reside within the increasing Hyades trend atcooler temperatures is metal weak (HIP 42499, [Fe/H]=-0.56), resulting in [O/Fe]=+0.47.The enhanced [O/Fe] ratio at this low metallicity is unsurprising and coincides with thecharacteristic field dwarf enhancements observed as a function of decreasing temperaturefor oxygen in other metal poor field stars (Abia & Rebolo 1989). If the abundance trendobserved by Schuler et al. (2004) and Schuler et al. (2006) is age dependent, the lack of adistinct trend of increasing [O/Fe] with decreasing abundance may point to the stars in thesample being older than the Hyades, not inconsistent with the 2.7 Gyr age of the dominantsubsample identified above. If not an age-related effect, then an as yet unknown dichotomybetween oxygen abundances in field stars and cluster stars would have to be explored withabundances of field stars of quantifiable age.For the giant stars in the sample, oxygen abundances have been derived from the infraredtriplet and from the forbidden line at λ blends driver of MOOG,following the approach of Schuler et al. (2006).In examining the giant triplet abundances, a similar effect as in the dwarfs is observedas temperatures decrease with enhancements in oxygen abundances derived from both the7774 ˚A and 7775 ˚A lines. NLTE corrections were applied to the λ ∼
4. SUMMARY
The existence of spatially unassociated groups of stars moving through the solar neigh-borhood with common U and V kinematics has been explored for over half a century (Eggen1958). Despite this long history, the exact origins of these so called moving groups is stilla matter of some debate. The classical view contends that they are dissolved open clusterswhich have retained common kinematics and drifted into spatially elongated stellar streams.If this is indeed true, moving group members should possess similar characteristics to those ofopen cluster stars: particularly, common chemical abundances and residence along a distinctevolutionary sequence in an HR diagram. 21 –In order to address the viability of moving groups being dissolved open clusters, wehave performed a high resolution spectroscopic abundance analysis of a 34 star sample ofthe kinematically distinct Wolf 630 moving group, selected for its residence in a sparselypopulated region of the UV plane in the solar neighborhood. Our abundance measurementsreveal that the sample can not be characterized by a uniform abundance pattern. Theindividual stars have been closely scrutinized, making use of abundances, evolutionary stateand qualitative age information to constrain membership as an unlikely, possible or likelymember of a subsample with a dominant abundance trend and consistent age. There appearsto be a group with a weighted mean of [Fe/H]=-0.01 ± ± I versus Fe II abundances in the coolest dwarfs of thesample, likely attributable to increasing NLTE effects as a function of cooling temperature.We find the necessity to apply NLTE corrections of 0.10-0.20 dex to Na abundances ingiant stars. Finally, we derived oxygen abundances for the stars in the sample from boththe forbidden line at 6300 ˚A and the near-IR triplet. First, we find evidence for blendingin the IR triplet in both dwarfs and giant stars, possibly by Fe I features near the λ I in low log g cool giants are importantand cannot be accounted for by extrapolating current NLTE calculations. Finally, we findreliable oxygen abundances from the forbidden line in giant stars and again find evidence ofincreased NLTE effects as a function of cooling temperature manifested in increased tripletderived abundances.The authors would like to gratefully acknowledgement support for this work providedby NSF grants AST-0908342 and AST-0239518. Furthermore, we would like to thank thereferee for many useful comments which place the work into a broader context. REFERENCES
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This preprint was prepared with the AAS L A TEX macros v5.2.
26 –
Fig. 1.— Sample normalized spectra of our 34 star sample. The top panel shows a giantstar and the bottom displays a dwarf. The typical continuum level S/N in these spectra are ∼ Physical log g0 1 2 3 4 5012345 Spec g higherPhot g higher4000 5000 6000 70004000500060007000 4000 5000 6000 70004000500060007000 Photometric Temperature (K)Spec T HigherPhot T Higher
Fig. 2.— The spectroscopic temperatures are plotted versus photometric temperatures inthe left plot and the spectroscopic gravities versus physical gravities are plotted in the rightplot. The line is plotted to show perfect agreement between the two values. The differencesbetween the spectroscopic and photometric parameters agree within the uncertainties in therespective mean differences. 28 – [Fe/H]-1 -0.5 0 0.5 100.010.020.03 -1 -0.5 0 0.5 100.010.020.03
Fig. 3.— The metallicities of our 34 stars are plotted as guassians with central peaks ata given star’s metallicity and σ equal to the uncertainty in the [Fe/H]. The guassians arenormalized to unit area and summed to yield the smoothed abundance histogram. The peakat [Fe/H] ∼ -0.50 is from 3 low metallicity stars and the bump at [Fe/H] ∼ Fig. 4.— The [Fe/H] is plotted versus temperature for the full sample of stars (top) and thepossible (red) and likely (green) homogeneous members (bottom). The solid line gives theweighted mean of the sample while the dotted lines are 3- σ deviations from this mean. Ifa star rests within the dotted lines (i.e. the abundance band) within its respective uncer-tainty, then it is considered homogeneous with the dominant sample. Those stars which restfar outside the abundance band in the full sample plot are iteratively removed as unlikelymembers until convergence to a dominant abundance is achieved, as seen in the bottom plot. 30 – Fig. 5.— Sample lithium synthesis for HIP 23852. The crosses are the observed spectrumwhile the lines are lithium abundances of logN(Li)=2.30, 2.00 (best fit) and 1.97. 31 –
Pleiades - 120 Myr7000 6000 5000 40000123 Temperature (K) M67 - 5 GyrTemperature (K)7000 6000 5000 40000123 Hyades 625 MyrTemperature (K)7000 6000 5000 40000123NGC 752 - 2.5 GyrTemperature (K)7000 6000 5000 40000123
Fig. 6.— Lithium abundances for the Pleiades (top left-King et al. (2000)), the Hyades (topright-Balachandran (1995), NGC752 (bottom left-Sestito et al. (2004)) and M67 (bottomright-Jones et al. (1999) (plotted as crosses) and our Wolf 630 candidates. Filled hexagonsare for dwarfs, filled triangles are upper limits for dwarfs, open hexagons are for subgiantsopen triangles are upper limits for subgiants. Specific abundances for individual stars arediscussed in more detail in the text. 32 –
Fig. 7.— The HR diagram of the final candidate members of a common chemical groupwith the distinct UV kinematics of the classical Wolf 630 group. Green points are likelymembers while red points are possible members. Unlikely members are plotted as blackpoints. Yale-Yonsei isochrones of 2.2, 2.7 and 3.2 Gyr iare shown. 33 –
Fig. 8.— Plot of the U and V kinematics for the sample with likely members plotted in redand possible members in green. Black points are non-members 34 –
Fig. 9.— The abundances [Na/Fe] (top) and [Al/Fe] (bottom) for all stars with measurablelines of Na and/or Al are plotted versus surface gravity. The solid line gives the weightedmean [X/Fe] for the dwarfs, neglecting the two with unusually low [Na/Fe]. The dottedline gives the weighted mean [X/Fe] for the subgiants and giants, neglecting the giant withunusally high [Na/Fe] and [Al/Fe]. Subgiant and giant abundances are ∼ Temperature (K)4000 4500 5000 5500 6000 65000123
Fig. 10.— The difference [Fe I I/H]-[Fe I /H] is plotted versus temperature. Notice the clearoverionization in the two coolest dwarfs of the sample. 36 – Fig. 11.— Differences in oxygen abundances for dwarf stars derived from the infrared triplet.The top plot shows the difference in the abundance from the 7774 line and the 7771 line.The difference in abundance between these two lines for the coolest two stars in the sampleis of order 0.20 dex. The difference between the 7775 line and the 7771 line is slightly moremodest, but the general trend is for the cooler stars to yield slight abundance enhancements. 37 –
Fig. 12.— Oxygen abundances [O/Fe] versus temperature for the Wolf 630 sample that weredetermined to be chemically homogeneous (black), the Pleiades (blue) and the Hyades (red). 38 –
Fig. 13.— The [O/H] abundance from the forbidden line (black hexagons) for subgiant andgiant stars is plotted versus temperature in all windows. The top plot shows the NLTE[O/H] abundances from the 7771 line of the triplet (blue triangles), the middle plot showsthe NLTE [O/H] from the 7774 line of the triplet (green triangles) and the bottom plotgives NLTE [O/H] from the 7775 line (red triangles). The abundances derived from the7771 line agree well, after NLTE corrections, with abundances from the forbidden line, but[O/H] abundances from the redder lines of the triplet increase as a function of decreasingtemperature. 39 –
Temperature (K)4500 5000 5500 6000-1.5-1-0.500.5 Surface Gravity1 2 3 4-1.5-1-0.500.5
Fig. 14.— Differences in oxygen abundance between the forbidden line at 6300 ˚A and 7771˚A for the giant and subgiant stars. The NLTE corrected abundances from the triplet linegenerally agree with the 6300 forbidden line, with the exception of the metal weak cool star,HIP 114155. The clear agreement between the forbidden and the NLTE triplet abundanceuntil reaching a low surface gravity possibly indicates that greater than expected NLTEeffects impact the triplet abundances in more evolved stars.
Table 1. Kinematic InformationHIP π PM RA PM DEC Radial Velocity U V W(mas) (mas/yr) (mas/yr) (kms − ) (kms − ) (kms − ) (kms − )102531 31.69 ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± Table 1—ContinuedHIP π PM RA PM DEC Radial Velocity U V W(mas) (mas/yr) (mas/yr) (kms − ) (kms − ) (kms − ) (kms − )34440 10.68 ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± Table 1—ContinuedHIP π PM RA PM DEC Radial Velocity U V W(mas) (mas/yr) (mas/yr) (kms − ) (kms − ) (kms − ) (kms − ) 43 –Table 2. Solar Equivalent WidthsWavelength Ion Excitation Potential log(gf) Equivalent Width˚A eV m˚A5505.881 Fe I I I I I I I I I I I I I I I I I I I I I I I I I I I I I I I I I I I I I I I I I I I I I I I I I I I I I I I I I I I I I I I I I I I I I I I I I I I I I I I I I I I I I I I I I I I I I I I I I I I I I I I I I I I I I I I I I I I I I I I I I I I I I I I I I I I I I I I I I I I I I I I I I I I I I I I I I I I I II II II II II II II II II II II I I I I I I I I I I I I I I I I I I I I I I I I I I I I I I I I I I I I I I II II II II II II II II I I I I I I I I I I I I I I I I I I I II II II spec Log(g) χ [Fe/H]K cms − kms − ±
59 3.80 ± ± ± ±
56 4.52 ± ± ± a ±
88 4.67 ± ± ± ±
69 2.40 ± ± ± ±
100 4.10 ± ± ± ±
100 3.75 ± ± ± ±
86 2.10 ± ± ± a ±
75 1.34 ± ± ± ±
40 4.36 ± ± ± ±
54 2.68 ± ± ± ±
30 2.71 ± ± ± ±
57 4.44 ± ± ± ±
32 2.09 ± ± ± ±
38 4.22 ± ± ± ±
31 4.57 ± ± ± ±
60 4.11 ± ± ± ±
55 4.40 ± ± ± ±
39 2.43 ± ± ± ±
35 2.53 ± ± ± ±
38 4.07 ± ± ± ±
62 2.54 ± ± ± ±
53 2.58 ± ± ± ±
37 3.77 ± ± ± ±
39 4.41 ± ± ± ±
32 4.41 ± ± ± a ±
200 1.39 ± ± ± ±
42 4.52 ± ± ± ±
55 4.35 ± ± ± b ±
79 4.70 ± ± ± ±
41 4.42 ± ± ± a Surface gravities for these stars are physical, calculated as discussed in the test. b Surface gravity for this star is physical, calculated as discussed in the text. The microtur-bulence was set to 0. Table 3—ContinuedHIP T spec
Log(g) χ [Fe/H]K cms − kms − ±
74 4.56 ± ± ± ±
34 2.61 ± ± ± ±
65 2.50 ± ± ± ±
42 2.45 ± ± ± Table 4. All Equivalent WidthsWavelength Ion Excitation Potential log(gf) HIP102531 EQW HIP102531 LogN(x) HIP105341 EQW HIP105341 LogN(x) HIP11033 EQW HIP11033 LogN(x) HIP112222 EQW HIP112222 LogN(x) HIP112447 EQW HIP112447 LogN(x) HIP113622 EQW HIP113622 LogN(x) HIP114155 EQW HIP114155 LogN(x) HIP114924 EQW HIP114924 LogN(x) HIP12784 EQW HIP12784 LogN(x) HIP13701 EQW HIP13701 LogN(x) HIP14501 EQW HIP14501 LogN(x) HIP17792 EQW HIP17792 LogN(x) HIP23852 EQW HIP23852 LogN(x) HIP29525 EQW HIP29525 LogN(x) HIP29843 EQW HIP29843 LogN(x) HIP33671 EQW HIP33671 LogN(x) HIP34440 EQW HIP34440 LogN(x) HIP3455 EQW HIP3455 LogN(x) HIP3559 EQW HIP3559 LogN(x) HIP36732 EQW HIP36732 LogN(x) HIP3992 EQW HIP3992 LogN(x) HIP40023 EQW HIP40023 LogN(x) HIP41484 EQW HIP41484 LogN(x) HIP42499 EQW HIP42499 LogN(x) HIP4346 EQW HIP4346 LogN(x) HIP43557 EQW HIP43557 LogN(x) HIP45617 EQW HIP45617 LogN(x) HIP5027 EQW HIP5027 LogN(x) HIP50505 EQW HIP50505 LogN(x) HIP5286 EQW HIP5286 LogN(x) HIP53229 EQW HIP53229 LogN(x) HIP53465 EQW HIP53465 LogN(x) HIP6732 EQW HIP6732 LogN(x)˚A eV m˚A m˚A m˚A m˚A m˚A m˚A m˚A m˚A m˚A m˚A m˚A m˚A m˚A m˚A m˚A m˚A m˚A m˚A m˚A m˚A m˚A m˚A m˚A m˚A m˚A m˚A m˚A m˚A m˚A m˚A m˚A m˚A m˚A6154.230 Na I . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . I . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . I I . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . I . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . I . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . I . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . I . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . I . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . I . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . I . . . . . . . . . . . . I . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . I . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . I . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . I . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . I . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . I I . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . I . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . I . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . I . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . I . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . Table 4—ContinuedWavelength Ion Excitation Potential log(gf) HIP102531 EQW HIP102531 LogN(x) HIP105341 EQW HIP105341 LogN(x) HIP11033 EQW HIP11033 LogN(x) HIP112222 EQW HIP112222 LogN(x) HIP112447 EQW HIP112447 LogN(x) HIP113622 EQW HIP113622 LogN(x) HIP114155 EQW HIP114155 LogN(x) HIP114924 EQW HIP114924 LogN(x) HIP12784 EQW HIP12784 LogN(x) HIP13701 EQW HIP13701 LogN(x) HIP14501 EQW HIP14501 LogN(x) HIP17792 EQW HIP17792 LogN(x) HIP23852 EQW HIP23852 LogN(x) HIP29525 EQW HIP29525 LogN(x) HIP29843 EQW HIP29843 LogN(x) HIP33671 EQW HIP33671 LogN(x) HIP34440 EQW HIP34440 LogN(x) HIP3455 EQW HIP3455 LogN(x) HIP3559 EQW HIP3559 LogN(x) HIP36732 EQW HIP36732 LogN(x) HIP3992 EQW HIP3992 LogN(x) HIP40023 EQW HIP40023 LogN(x) HIP41484 EQW HIP41484 LogN(x) HIP42499 EQW HIP42499 LogN(x) HIP4346 EQW HIP4346 LogN(x) HIP43557 EQW HIP43557 LogN(x) HIP45617 EQW HIP45617 LogN(x) HIP5027 EQW HIP5027 LogN(x) HIP50505 EQW HIP50505 LogN(x) HIP5286 EQW HIP5286 LogN(x) HIP53229 EQW HIP53229 LogN(x) HIP53465 EQW HIP53465 LogN(x) HIP6732 EQW HIP6732 LogN(x)˚A eV m˚A m˚A m˚A m˚A m˚A m˚A m˚A m˚A m˚A m˚A m˚A m˚A m˚A m˚A m˚A m˚A m˚A m˚A m˚A m˚A m˚A m˚A m˚A m˚A m˚A m˚A m˚A m˚A m˚A m˚A m˚A m˚A m˚A5999.658 Ti I . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . I . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . I . . . . . . I . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . I . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . I . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . I . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . I . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . I . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . I . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . I . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . II . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . II . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . II . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . II . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . II . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . II . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . I . . . . . . . . . . . . . . . . . . I . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . I . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . I . . . . . . . . . . . . . . . . . . . . . . . . I . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . Table 4—ContinuedWavelength Ion Excitation Potential log(gf) HIP102531 EQW HIP102531 LogN(x) HIP105341 EQW HIP105341 LogN(x) HIP11033 EQW HIP11033 LogN(x) HIP112222 EQW HIP112222 LogN(x) HIP112447 EQW HIP112447 LogN(x) HIP113622 EQW HIP113622 LogN(x) HIP114155 EQW HIP114155 LogN(x) HIP114924 EQW HIP114924 LogN(x) HIP12784 EQW HIP12784 LogN(x) HIP13701 EQW HIP13701 LogN(x) HIP14501 EQW HIP14501 LogN(x) HIP17792 EQW HIP17792 LogN(x) HIP23852 EQW HIP23852 LogN(x) HIP29525 EQW HIP29525 LogN(x) HIP29843 EQW HIP29843 LogN(x) HIP33671 EQW HIP33671 LogN(x) HIP34440 EQW HIP34440 LogN(x) HIP3455 EQW HIP3455 LogN(x) HIP3559 EQW HIP3559 LogN(x) HIP36732 EQW HIP36732 LogN(x) HIP3992 EQW HIP3992 LogN(x) HIP40023 EQW HIP40023 LogN(x) HIP41484 EQW HIP41484 LogN(x) HIP42499 EQW HIP42499 LogN(x) HIP4346 EQW HIP4346 LogN(x) HIP43557 EQW HIP43557 LogN(x) HIP45617 EQW HIP45617 LogN(x) HIP5027 EQW HIP5027 LogN(x) HIP50505 EQW HIP50505 LogN(x) HIP5286 EQW HIP5286 LogN(x) HIP53229 EQW HIP53229 LogN(x) HIP53465 EQW HIP53465 LogN(x) HIP6732 EQW HIP6732 LogN(x)˚A eV m˚A m˚A m˚A m˚A m˚A m˚A m˚A m˚A m˚A m˚A m˚A m˚A m˚A m˚A m˚A m˚A m˚A m˚A m˚A m˚A m˚A m˚A m˚A m˚A m˚A m˚A m˚A m˚A m˚A m˚A m˚A m˚A m˚A5846.990 Ni I . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . I . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . I . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . I . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . I . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . I . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . I . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . I . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . I . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . I . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . I . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . I . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . I . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . I . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . II . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . II . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . II . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
57 –
Table 5. AbundancesHIP [Mg/H] [Si/H] [Ca/H] [Ti/H] [Ti2/H] [Cr/H] [Mn/H] [Ni/H] [Na/H] [Al/H] [Ba/H]102531 0.13 ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± . . . . . . -0.17 ± . . . . . . -0.20 ± . . . . . . ± . . . ± ± ± ± . . . -0.01 ± . . . ± ± ± . . . ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± . . . ± ± ± . . . ± ± . . . . . . . . . -0.51 ± . . . -0.61 ± ± ± ± ± . . . ± ± ± ± . . . ± ± ± ± ± ± ± ± ± ± . . . -0.02 ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± . . . ± ± . . . ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± . . . -0.12 ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± . . . ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± . . . -0.03 ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± Table 5—ContinuedHIP [Mg/H] [Si/H] [Ca/H] [Ti/H] [Ti2/H] [Cr/H] [Mn/H] [Ni/H] [Na/H] [Al/H] [Ba/H]40023 -0.01 ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± . . . ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± . . . ± ± ± ± . . . ± ± ± ± ± . . . ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± . . . -0.13 ± ± ± ± ± ± ± ± ± ± ± . . . -0.26 ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± . . . ± ± ± ± ± ± ± ± ± ± . . . ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± . . . -0.10 ± ± ± ± ± ± ± ∓ g = ± ∓ ζ = ± ± < [Fe/H] > ± ± ± ∓ g = ± ∓ ζ = ± ± < [Fe/H] > ± ± ± ∓ g = ± ∓ ζ = ± ± < [Fe/H] > ± ± Table 7. Oxygen AbundancesHIP EW EW EW LTE [O/H]
LTE [O/H]
LTE [O/H] [O/H]
NLTE [O/H]
NLTE [O/H]
NLTE [O/H] m˚A m˚A m˚A102531 157.7 139.7 114.0 0.48 ± ± ± ± ± ± ± ± ± ± . . . ± ± ± ± ± ± ± ± ± ± ± ± ± . . . -0.05 ± ± ± ± ± ± . . . -0.02 ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± . . . ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± . . . ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± . . . ± ± ± ± ± ± . . . ± ± ± ± ± ± . . . -0.08 ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± Table 7—ContinuedHIP EW EW EW LTE [O/H]
LTE [O/H]
LTE [O/H] [O/H]
NLTE [O/H]
NLTE [O/H]
NLTE [O/H] m˚A m˚A m˚A42499 21.2 17.7 13.5 -0.18 ± ± ± . . . -0.09 ± ± ± . . . . . . . . . . . . . . . . . . -0.10 ± . . . . . . . . . ± ± ± . . . ± ± ± ± ± ± . . . -0.17 ± ± ± ± ± ± . . . ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ≤ -0.205286 ≤ ≤ ± ± ≤ ± ± ≤ ≤ ≤ ≤ ± ± ≤ -0.25112447 2.30 ± ± ≤ spec [Fe/H]KUNLIKELY3559 5800 ±
38 -0.18 ± ±
200 0.24 ± ±
79 -0.08 ± ±
74 0.29 ± ±
69 0.12 ± ±
32 -0.52 ± ±
38 -0.14 ± ±
60 0.12 ± ±
55 -0.21 ± ±
32 -0.56 ± ±
55 -0.12 ± ±
41 -0.14 ± ±
56 0.02 ± ±
100 -0.34 ± ±
75 -0.58 ± ±
53 -0.15 ± ±
54 0.09 ± ±
39 -0.15 ± ±
62 0.10 ± ±
39 0.08 ± ±
34 -0.10 ± ±
88 -0.05 ± ±
40 0.06 ± ±
35 0.00 ± ±
42 -0.03 ± A. Notes on Individual Stars: Unlikely MembersA.1. HIP 3559: T=5800 logg=4.07 ξ =1.27 [Fe/H]=-0.18 This star resides above the ZAMS in the HR diagram. Ca II H and K measure-ments indicate an inactive chromosphere (logR’ HK =-5.16); the activity-age calibrations ofMamajek & Hillenbrand (2008) suggest an age of 9.4 ± ± g ≈ ± ⊙ . Assumingthis mass, as a ZAMS star of Pleiades age ( ≈
120 Myr), HIP 3559 would have had a main-sequence temperature of 6158 K. From the lithium abundance trend traced by the Pleiades(6), we infer that this star would have possessed an abundance of logN(Li) ∼ ∼ ⊙ at an age of ∼ ∼ A.2. HIP 4346: T=3820 logg=1.39 ξ =1.33 [Fe/H]=0.24 The metallicity of [Fe/H]=0.24 ± spec [Fe/H]K13701 4675 ±
30 -0.03 ± ±
57 -0.08 ± ±
31 -0.03 ± ±
37 -0.05 ± ±
42 -0.03 ± ±
65 -0.08 ± ±
59 0.07 ± ±
100 0.04 ± ±
86 0.00 ± χ νall χ νgroup χ νprobable Al/H -0.01 ± ± ± ± ± ± ± ± ± ± ± ± A.3. HIP 5027: T=4398 logg=4.70 ξ =0.00 [Fe/H]=-0.08 Hip 5027 has an Fe abundance which is consistent with the dominant [Fe/H] valuesexhibited by our sample. However, the abundances of other elements (Na, Al, Mn, Ni, Mgand Si) are all markedly sub-solar, and rest outside of the abundance bands for the fullsample. The lithium upper limit of logN(Li) ≤ -0.20 may place the star in the trend tracedby the Pleiades (Figure 6); however the significant spread in the Pleiades lithium abundancesas a function of decreasing temperature makes a firm conclusion regarding age difficult todraw. With the majority of elements disagreeing with the dominant abundance trends of theentire sample, this star is classified as an unlikely member of a chemically dominant group. A.4. HIP 5286: T=4683 logg=4.56 ξ =0.54 [Fe/H]=0.29 HIP 5286 is a member of the high metallicity “bump” at [Fe/H]= ∼ ≤ ± A.5. HIP 11033: T=4510 logg=2.40 ξ =1.60 [Fe/H]=0.12 The metallicity of HIP 11033 ([Fe/H] = 0.12 ± A.6. HIP 17792: T=4416 logg=2.09 ξ =1.50 [Fe/H]=-0.52 HIP 17792, with [Fe/H]=-0.52 ± α elements lead to the conclusion that thisstar is an unlikely candidate that is part of a dominant chemical subsample. A.7. HIP 23852: T=5778 logg=4.22 ξ =1.22 [Fe/H]=-0.14 This star resides above the ZAMS, raising the question of pre-main sequence or subgiantstatus. An isochrone age of 8.8 Gyr was estimated from Padova isochrones by Nordstr¨om et al.(2004). Using Yonsei-Yale isochrones, we find an age of 7.9 ± g =4.22, is consistent with a super ZAMS classification. Usinga stellar mass of 1.08 M ⊙ , inferred from the Yale-Yonsei isochrones, the ZAMS tempera-ture of this star would have been 6052 K. The ZAMS lithium abundance, inferred from thePleiades trend of Figure 6, logN(Li)=3.00, suggests a factor of 10 lithium depletion, consis-tent with theoretical calculations. The current lithium abundance of logN(Li)=2.00 is toolow for a PMS star, and appears consistent with the M67 Li-T eff trend, which implies thestar is an ∼ A.8. HIP 33671: T=6040 logg=4.40 ξ =1.38 [Fe/H]=-0.21 The metallicity of HIP 33671 is [Fe/H]=-0.21 ± A.9. HIP 42499: T=4994 logg=4.41 ξ =0.59 [Fe/H]=-0.56 HIP 42499 is a member of the metal-weak peak in the full sample [Fe/H] distribution(Figure 3). A Li upper limit of logN(Li) ≤ HK =-4.98) is much lower than the activitytrend for the Hyades, suggesting it is older than the Hyades. With an [Fe/H]=-0.56 ± A.10. HIP45617 T=4855 logg=4.35 ξ =1.01 [Fe/H]=-0.12 This star resides above the lower main sequence. The lithium upper limit (logN(Li) ≤ HK =-4.60, from the Ca II H and K survey of the solar neighborhood of D. Soderblom (privatecommunication), would place the star below the activity trend of the Hyades, qualitativelysuggesting that a Hyades age would be a reasonable lower limit. However, the spectroscopicsurface gravity is somewhat low for a dwarf star. A possible explanation is that overioniza-tion, observed in many cool cluster dwarfs (Schuler et al. (2003)), is yielding spuriously lowsurface gravities. With a greater number of atoms in ionized states, the gravity would haveto be artificially lowered to obtain ionization balance. However excellent agreement is seenbetween the spectroscopic gravity (log g =4.35) and the physical gravity (log g =4.38). Thestar’s [Fe/H]=-0.12 ± A.11. HIP 50505: T=5655 logg=4.42 ξ =1.16 [Fe/H]=-0.14 This star clearly resides on the main sequence, with a low Li upper limit of logN(Li) ≤ ≤ M67 age) dwarf. The star is clearly metal poor([Fe/H]=-0.14 ± A.12. HIP 112447: T=6095 logg=3.75 ξ =1.82 [Fe/H]=-0.34 This star has a distinctly low [Fe/H]=-0.34 ± A.13. HIP 114155: T=4348 logg=1.34 ξ =2.19 [Fe/H]=-0.58 The [Fe/H] of HIP 114155 is clearly low [Fe/H]= − ± ± ± B. Notes on Individual Stars: Possible MembersB.1. HIP 3992: T=4772 logg=2.58 ξ =1.59 [Fe/H]=-0.15 HIP 3992 has an [Fe/H]= − ± B.2. HIP 12784: T=4701 logg=2.68 ξ =1.49 [Fe/H]=0.09 The uncertainty associated with the metallicity of HIP 12784 ([Fe/H]=0.09 ± B.3. HIP 29843: T=6130 logg=4.11 ξ =1.52 [Fe/H]=0.12 HIP 29843 has a Li upper limit of logN(Li) ≤ ± ⊙ . Using this mass to determine the ZAMS temperature of the staryields T ZAMS =6678 K. This temperature, when compared to Figure 6, would have placedthis star in or on the blue-edge of the lithium dip while a dwarf. Currently, as a subgiant thathas emerged from the lithium dip, the lack of lithium suggests that the deepening convectionzone in the subgiant has not brought lithium back to the surface. This appears consistentwith the findings of Balachandran (1990) who also inferred little transport of lithium tothe surface in subgiants emerging from the lithium dip in M67. The metallicity of the star([Fe/H]=0.12 ± B.4. HIP 34440: T=4757 logg=2.43 ξ =1.46 [Fe/H]=-0.15 Fe, Ti and Cr for this star lay outside of the respective abundance bands. The otherelements all reside within their bands, consistent with homogeneity. We consider this staronly a possible member of a dominant homogeneous chemical group in our sample.
B.5. HIP 36732: T=4667 logg=2.54 ξ =1.44 [Fe/H]=0.10 Fe, Mn, Ni, Na and Mg all appear slightly enriched when compared to the dominantabundance bands. While the other elements have abundances within their respective bandsthe consistent overabundances for Fe, Mn, Ni, Na and Mg suggest this star be classified asonly a possible member of a chemically dominant subsample. 72 –
B.6. HIP 41484: T=5855 logg=4.41 ξ =1.17 [Fe/H]=0.08 The Fe abundance of HIP 41484 ([Fe/H]=0.08 ± ± B.7. HIP 43557: T=5816 logg=4.52 ξ =1.15 [Fe/H]=-0.03 The [Fe/H] of HIP 43557 matches the mean abundance of our entire sample. Mg, Naand Si however, do not appear to lay within their respective abundance bands. The averageabundances for Ti, Ti II , Cr and Ba all rest near the sample mean abundances irrespectiveof their uncertainties, suggesting a high degree of homogeneity. In examining the lithiumabundance, the star rests below the 5 Gyr trend in the Li-T eff relation, perhaps suggestingan older age. Although the [Fe/H] agrees well with the mean metallicity and the averageabundances of multiple elements are close to the respective mean abundances for the group,the evidence from Mg, Na and Si and the lower lithium abundance make HIP 43557 a possiblegroup member. B.8. HIP103983: T=5750 logg=4.52 ξ =1.16 [Fe/H]=0.02 The status of this star is somewhat of an enigma. While an isochrone fit is consis-tent with placement on the subgiant branch of an 8.5 ± g =4.52 suggests a dwarf luminosity class. Note the significant uncertainty inthe surface gravity measurement (0.20 dex). Valenti & Fischer (2005) find a surface gravityof 4.37, consistent with the lower limit of the spectroscopic gravity derived here. Furthercomparing surface gravity estimates, the physical surface gravity derived for this star is log g =4.22 ± ± ⊙ . This yields a ZAMS tempera-ture, T ZAMS =5754 K, which coincides with a Pleiades lithium abundance of logN(Li)=3.00on the so-called “lithium plateau”. Assuming this as a reasonable ZAMS lithium abundance,this star would have ≈ ⊙ obtained from the Clemson-American Universityof Beirut Stellar Evolution Code. This perhaps points to the star not being a clear subgiant,however, the evolutionary status of this star remains uncertain. Examining the abundances,the star has an [Fe/H]=0.02 ± α and Fe peakelements, likewise, yield abundances that reside within the respective abundance bands thatare used to characterize homogeneity. Considering the uncertainties in the surface gravitiesand the potential that the lithium abundance negates a subgiant classification and that theabundances are homogeneous with the rest of the sample, we consider this star a possiblemember of a chemically homogeneous sub-group. B.8.1. HIP 105341: T=4005 logg=4.67 ξ =0.83 [Fe/H]=-0.05 This star is the coolest dwarf in the sample. The chromospheric activity (logR’ HK =-4.552) from Gray et al. (2006) suggests this is a relatively active star, which may be consistentwith PMS status, although it is not inconsistent with a main sequence age. The activityderived age, using the updated age-activity relation of Mamajek & Hillenbrand (2008) is0.85 Gyr ± ≤ -0.25) may plausibly place the star inthe lithium trend traced by the Pleiades in Figure 6, but without lithium abundances formore cool Pleids the picture is unclear. The Fe abundance of the star ([Fe/H]=-0.05 ± B.9. HIP 114924: T=6179 logg=4.36 ξ =1.59 [Fe/H]=0.06 With [Fe/H]=0.06 ± ± C. Note on Individual Stars: Likely MembersC.1. HIP 3455: T=4860 logg=2.53 ξ =1.49 [Fe/H]=0.00 This star has an [Fe/H]=0.00 ± C.2. HIP 6732: T=4665 logg=2.45 ξ =1.58 [Fe/H]=-0.03 This star has a metallicity ([Fe/H]=-0.03 ± C.3. HIP 13701: T=4675 logg=2.71 ξ =1.37 [Fe/H]=-0.03 This star clearly resides within the dominant [Fe/H] band. Indeed, its abundance isnearly identical to the weighted mean of the sample. It appears consistent with the metallic-ity bands for all elements. This homogeneity with the rest of the sample leads to classifyingHIP 13701 as a likely group member. 75 –
C.4. HIP 14501: T=5785 logg=4.44 ξ =1.24 [Fe/H]=-0.08 Having [Fe/H]=-0.08 ± ≤ C.5. HIP 29525: T=5710 logg=4.57 ξ =1.28 [Fe/H]=-0.03 This star resides on the main sequence of the isochrones in Figure 7. The lithiumabundance (logN(Li)=2.03 ± ± C.6. HIP 40023: T=5290 logg=3.77 ξ =1.21 [Fe/H]=-0.05 HIP 40023 has an [Fe/H]=-0.05, which is within the metallicity band of the sample.Indeed, its abundances across multiple elements fit inside the respective metallicity bands.The small spread in abundances for the star itself and relative to the overall sample abun-dance bands, lead to classification of this star as a likely member of a dominant chemicallyhomogeneous 2-3 Gyr subsample.
C.7. HIP 53229: T=4690 logg=2.61 ξ =1.47 [Fe/H]=-0.10 The Fe abundance of HIP 53229 ([Fe/H]=-0.10 ± C.8. HIP 53465: T=4570 logg=2.50 ξ =1.30 [Fe/H]=-0.08 The metallicity of HIP 53465 ([Fe/H]=-0.08 ± II are foundto lay outside of the sample abundance bands, the remaining elements show a high degreeof homogeneity. For most elements, the abundances lay within the abundance band. Thus,this star is considered a likely member of a chemically dominant group in our sample. C.9. HIP 102531: T=6238 logg=3.80 ξ =1.85 [Fe/H]=0.07 The metallicity of HIP 102531 ([Fe/H]=0.07 ± σ cutoffof the mean Fe abundance of the whole sample. However, this star resides within the meanabundance bands of Al, Ba, Ca, Mg, Mn, Ni, Si, Ti and Ti 2. In Figure 6, this is the warmestsample star that has lithium, and can be seen to lay significantly beneath any trend traced byany of the plotted open cluster dwarf abundances. From the HR diagram, this star lies alongthe early subgiant branch of a 2.7 Gyr isochrone, which indicates a mass of 1.5 ± ⊙ .In comparing this star with Figure 11 of Balachandran (1995), who plot lithium abundancesfor open clusters versus stellar mass, the lithium abundance for the derived mass appears tobe between the trends for M67 and NGC 752, consistent with the estimated isochrone ageof 2.7 Gyr. This would suggest that the star has suffered subgiant and/or main sequenceLithium dip depletion. Recognizing that the majority of elements suggest this star is partof a chemically homogeneous subsample, and the ∼ C.10. HIP 112222: T=6369 logg=4.10 ξ =1.69 [Fe/H]=0.04 HIP 112222, with [Fe/H]=0.04 ± ∼ ⊙ consistent with this possibly being a lithium dip star, providing anexplanation for the apparently low upper limit lithium abundance of logN(Li) ≤ C.11. HIP 113622: T=4295 logg=2.10 ξ =1.52 [Fe/H]=0.00 With [Fe/H]=0.00 ± -0.6 -0.4 -0.2 0 0.2 0.4-1-0.500.51 [Fe/H] -0.6 -0.4 -0.2 0 0.2 0.4-1-0.500.51 [Fe/H]-0.6 -0.4 -0.2 0 0.2 0.4-1-0.500.51 [Fe/H] -0.6 -0.4 -0.2 0 0.2 0.4-1-0.500.51 [Fe/H]-0.6 -0.4 -0.2 0 0.2 0.4-1-0.500.51 [Fe/H] -0.6 -0.4 -0.2 0 0.2 0.4-1-0.500.51 [Fe/H] Fig. 15.— Abundance trends for all stars versus Fe/H. 79 – -0.6 -0.4 -0.2 0 0.2 0.4-1-0.500.51 [Fe/H] -0.6 -0.4 -0.2 0 0.2 0.4-1-0.500.51 [Fe/H]-0.6 -0.4 -0.2 0 0.2 0.4-1-0.500.51 [Fe/H] -0.6 -0.4 -0.2 0 0.2 0.4-1-0.500.51 [Fe/H]-0.6 -0.4 -0.2 0 0.2 0.4-1-0.500.51 [Fe/H]-0.6 -0.4 -0.2 0 0.2 0.4-1-0.500.51 [Fe/H]