Intrinsic Colors, Temperatures, and Bolometric Corrections of Pre-Main Sequence Stars
aa r X i v : . [ a s t r o - ph . S R ] J u l Draft version July 11, 2013
Preprint typeset using L A TEX style emulateapj v. 08/22/09
INTRINSIC COLORS, TEMPERATURES AND BOLOMETRIC CORRECTIONS OF PRE-MAIN SEQUENCESTARS
Mark J. Pecaut, Eric E. Mamajek
University of Rochester, Department of Physics and Astronomy, Rochester, NY 14627-0171, USA
Draft version July 11, 2013
ABSTRACTWe present an analysis of the intrinsic colors and temperatures of 5-30 Myr old pre-main sequence(pre-MS) stars using the F0 through M9 type members of nearby, negligibly reddened groups: η Chacluster, TW Hydra Association, β Pic Moving Group, and Tucana-Horologium Association. To checkthe consistency of spectral types from the literature, we estimate new spectral types for 52 nearbypre-MS stars with spectral types F3 through M4 using optical spectra taken with the SMARTS 1.5-m telescope. Combining these new types with published spectral types, and photometry from theliterature (Johnson-Cousins
BV I C , 2MASS JHK S and WISE W W W
3, and W eff ) and bolometric corrections(BCs) for our pre-MS star sample through comparing their photometry to synthetic photometrygenerated using the BT-Settl grid of model atmosphere spectra. We derive a new T eff and BC scalefor pre-MS stars, which should be a more appropriate match for T Tauri stars than often-adopteddwarf star scales. While our new T eff scale for pre-MS stars is within ≃
100 K of dwarfs at a givenspectral type for stars < G5, for G5 through K6, the pre-MS stars are ∼
250 K cooler than theirmain sequence counterparts. Lastly, we present (1) a modern T eff , optical/IR color, and bolometriccorrection sequence for O9V-M9V MS stars based on an extensive literature survey, (2) a revised Q-method relation for dereddening UBV photometry of OB-type stars, and (3) introduce two candidatespectral standard stars as representatives of spectral types K8V and K9V. Subject headings: open clusters and associations: individual( η Cha cluster, TW Hydra Association, β Pic Moving Group, Tucana-Horologium Association); — stars: pre-main sequence— stars: fundamental parameters (colors, temperatures) INTRODUCTION AND BACKGROUND
Knowledge of the stellar intrinsic color locus is an es-sential ingredient in studying young stellar populations.Recently-formed stars are typically either distant, andthus outside of the “Local Bubble” of low reddening inthe solar vicinity, or they are still embedded in their na-tal molecular cloud. Hence, we cannot assume negligi-ble reddening and extinction for most pre-main sequence(pre-MS) stars. Interstellar reddening is conventionallyestimated using tabulated intrinsic colors of dwarf fieldstars on the main sequence (e.g., Kenyon & Hartmann1995). However, this likely introduces systematic errorsin the analysis since the pre-MS stars are in a differentevolutionary stage than the main sequence calibratorsand may not exhibit “standard” dwarf colors. Accu-rate H-R diagram placement depends on accurate ex-tinction and effective temperature (T eff ) estimates. Ifthe extinction or T eff is systematically in error becauseof systematics in the intrinsic color and T eff tabulationsas a function of spectral type, this will obviously intro-duce systematic errors in the H-R diagram placement andages and masses inferred from comparison to evolution-ary tracks. For pre-MS stars, systematic errors in agesmay systematically shift the inferred timescales for pro-toplanetary disk dissipation and giant planet formation(e.g., Mamajek 2009; Bell et al. 2013). The H-R diagrampresents an opportunity for stellar theoretical evolution-ary models to make contact with observations, but if our H-R diagram placement is plagued with systematic er-rors, this makes testing evolutionary models impossible.Thus it is imperative that the intrinsic color and T eff scale be accurately known and as free of systematic er-rors as possible.Previous studies have noted that the intrinsic colorsof young stars differ from that of main sequence stars(e.g., Gullbring et al. 1998; Luhman 1999; Bell et al.2012). Stauffer et al. (2003) investigated Pleiades (age ∼
125 Myr; Stauffer et al. 1998) zero-age main sequenceK-stars exhibiting bluer B – V colors as a function ofspectral type than their counterparts in Praesepe (age ∼
750 Myr; G´asp´ar et al. 2009), and concluded thatthe effect was age-dependent. Their study identifiedstarspots and plages as the most likely cause of thebluer B – V colors and concluded that all young K dwarfswill exhibit this effect. Da Rio et al. (2010) constructeda young star intrinsic color sequence in their studyof the star-formation history of Orion Nebula Clusterby merging synthetic colors interpolated to a 2 Myrisochronal surface gravity with empirical colors fromKenyon & Hartmann (1995). However, this implicitlycharges the color discrepancy solely to lower surface grav-ity. Furthermore, synthetic near-infrared colors such as J – H and H – K S do not follow observed intrinsic colorsequences for M-dwarfs redder than V – K S > ∼ . η Cha cluster, the ǫ Cha cluster andthe TW Hydra Association (TWA). The Luhman et al.(2010b) tabulation is empirically derived and thus doesnot depend on synthetic colors.Here we offer an alternative and expanded pre-MS in-trinsic color tabulation by including optical
BV I C col-ors, including earlier spectral types, and using the youngstars’ spectral energy distributions to estimate effectivetemperatures and construct a temperature and bolomet-ric correction scale. In this work we examine spectraltypes F0 through M9.5, but our temperature scale onlyextends to types as late as M5. In Section 2 we describeour sample, and in Section 3 we describe the spectroscopyand photometry data used for our analysis. In Section 4we describe our spectroscopy, the derivation of our pre-MS intrinsic colors, and the derivation of our effectivetemperature and bolometric correction scale for pre-MSstars. Finally, in Section 5 we compare our temperaturescale and angular diameter estimates to previous resultsin the literature. SAMPLE SELECTION
Our sample consists of members of young ( < ∼
30 Myr),nearby moving groups including the β Pic moving group,TW Hydra Association (TWA), Tucana-Horologiummoving group (Tuc-Hor) and the η Cha cluster. Themembers of these groups are all predominantly pre-main sequence (with the exception of a handful ofintermediate-mass A-type stars, which we omit) andthus will allow us to study the observed color differ-ences between main sequence stars and pre-main se-quence stars. β Pic, TWA and Tuc-Hor members are lessthan 75 pc distant and thus lie within the Local Bubble,within which objects are subject to negligible reddening(E( B – V ) < H < ∼ cm − inside the localbubble from Cox & Reynolds 1987 and N(H I)/E( B – V )= 4.8 × cm − mag − from Savage & Mathis 1979). η Cha is slightly more distant ( ∼
95 pc) but also hasA V ≃ BV I C , Two Micron AllSky Survey (2MASS; Skrutskie et al. 2006) JHK S pho-tometric bands and the Wide-Field Infrared Survey Ex-plorer (WISE; Wright et al. 2010) W W W W µ m, 4.6 µ m, 12 µ m, and 22 µ m, re-spectively.Our sample was assembled from group membershiplists from Mamajek et al. (1999), Luhman & Steeghs(2004), Lyo et al. (2004), Song et al. (2004),Zuckerman & Song (2004), Scholz et al. (2005),Torres et al. (2006), L´epine & Simon (2009), Kiss et al.(2011), Schlieder et al. (2010), Rice et al. (2010b),Zuckerman et al. (2011), Shkolnik et al. (2011),Rodriguez et al. (2011), Schlieder et al. (2012b) andSchneider et al. (2012b). Following the Weinberger et al.(2012) and Mamajek (2005) studies, we reject TWA 22as a member of TWA based on its discrepant spacemotion. However, we retain it as a member of β Pic, following Teixeira et al. (2009). In addition, based onthe study of Mamajek (2005) and parallax data fromWeinberger et al. (2012), stars TWA 14, TWA 15A,TWA 15B, TWA 17, TWA 18, TWA 19A, TWA 19B,and TWA 24 are likely members of the Lower Centaurus-Crux subgroup of the Scorpius-Centaurus OB associationand thus may be subject to non-negligible reddening, sowe exclude them from our sample. We include TWA 9as a member of TWA, though Weinberger et al. (2012)reject it. We discuss our justification for including it inAppendix A. Our sample includes 54 members of β Picwith spectral types F0-M8, 34 members of TWA withspectral types K3-M9.5, 45 members of Tuc-Hor withspectral types F2-M2, and 15 members of η Cha withspectral types K5-M5.75. DATA
Spectroscopy
Though the objects in our sample have published spec-tral types, they are from a variety of sources and reso-lutions. In order to check the consistency of spectraltypes in the literature, we obtain new spectral typesusing a grid of standards from Keenan & Yorka (1988);Keenan & McNeil (1989), Kirkpatrick et al. (1991) andHenry et al. (2002). We acquired low-resolution blue( ∼ ∼ β Pic, TWA and η Cha. Thestars chosen for spectroscopy were selected based on (1)target brightness and (2) optimizing telescope time toavoid interfering with higher priority programs. Thefaintest targets would require prohibitively large expo-sure times with the RC spectrograph on the SMARTS1.5m telescope to obtain useful S/N for spectral classifi-cation. This spectroscopic sample includes stars down tom V ∼
14 mag, with spectral types F3-M4. Observationswere made in queue mode with the RC spectrograph be-tween February 2011 and July 2011. The blue spectrawere taken with the “26/Ia” setup which consists of agrating with groove density of 600 grooves mm − , blazewavelength 4450˚A and no filter. The red spectra weretaken with the “47/Ib” setup which consists of a gratingwith groove density of 831 grooves mm − , blaze wave-length 7100˚A, and a GG495 filter. Both used a slit withof 110.5 µ m. The resolution for the blue and red spectraare ∼ ∼ . The three images are median com-bined, bias-trimmed, overscan- and bias-subtracted andflat-fielded. The spectrum is wavelength-calibrated and,as a final step, we normalize the spectra to the contin-uum with a low order spline in preparation for spectralclassification. Photometry
After compiling the list of nearby < ∼
30 Myr old stars,we assembled the most precise available photometry fromthe literature, listed in Table 1. All stars in our list oung Stellar Colors and Temperatures 3have counterparts in the 2MASS Point Source Catalog.A few objects are known binaries but are unresolved inthe 2MASS catalog. In these cases, we retain the pri-mary in our lists but do not include the secondary sinceit would be of limited use without distinct near-infraredphotometry. Tuc-Hor member TYC 7065-0879-1 (K0V;Torres et al. 2006) is a 1.8 ′′ binary, resolved in Tycho-2(Høg et al. 2000) but unresolved in 2MASS. The 2MASSPSF photometry differs significantly from the 2MASSaperture photometry (e.g., H PSF − H AP = 0 .
356 mag),presumably due to a poorly fit PSF to the unresolvedbinary. Thus for TYC 7065-0879-1 we adopt unresolved
BV I C optical photometry and the unresolved 2MASSaperture photometry. All other objects in our samplehave 2MASS PSF photometry which agrees well withthe aperture photometry (when available) and there-fore we simply adopt the PSF photometry. We adopt WISE bands W W W
3, and W WISE
All-Sky Point Source Catalog, centered at 3.4,4.6, 12, and 22 µ m, respectively (Wright et al. 2010).Objects saturated in W < ∼ , so to avoid these biases we exclude W W < . Hip-parcos catalog entries, we adopt V and B – V photome-try from that catalog ESA (1997). We then fill miss-ing B – V photometry using Tycho-2 photometry ( B T , V T ) converted to Johnson B – V with the conversions ofMamajek et al. (2002, 2006), resorting to the conversionsin Høg et al. (2000) when B T − V T > .
0. We adoptedAAVSO Photometric All-Sky Survey (APASS) Data Re-lease 6 (Henden et al. 2012) BV and SACY (Torres et al.2006) BV I C photometry where available. Conservativeestimates for SACY BV I C photometric uncertainties ob-tained with the FOTRAP instrument are 0.01 mag forstars brighter than V ∼
12 mag (C.A.O. Torres, 2012private communication). We only adopted B – V colorswhen σ B − V < .
08 mag. We adopted V – I C photometryfrom Torres et al. (2006), Lawson et al. (2001) and the Hipparcos catalog, when it was directly observed (value“A” in field H42), since a significant portion of the tab-ulated V – I C photometry in the Hipparcos catalog is in-ferred from photometry in other bands or from the spec-tral type of the star. Though it was available for many ofour objects, we did not adopt DEep Near-Infrared Surveyof the Southern Sky (DENIS) i band photometry since itsaturates at ∼ i photometry. ANALYSIS
Spectral Classification
The optical spectra were visually classified by directlycomparing them with spectral standards using a customspectral software tool, sptool , described in Pecaut et al.(2012). F- and G-type standards are taken from Table 2of Pecaut et al. 2012; K- and M-type standards are listedin Table 2. For the blue spectra, the F-type stars wereclassified using the strength and profile of the Balmerlines, with particular attention to the wings of the lines incase the line depths were filled in by chromospheric emis- http://wise2.ipac.caltech.edu/docs/release/allsky/expsup/sec6_3c.html See or rumtph.org/pecaut/sptool/ . Fig. 1.—
A spectrum of η Cha member RECX 7 (K5IV(e))with spectral standards K4Ve (TW PsA), K5III (HD 82668), K5V(HD 36003), and K6Va (GJ 529). The primary regions used forspectral classification of K-type stars are highlighted in grey. sion. In addition, we use the G-band at ∼ ∼ ∼ . We also made use of Ca I lines at λλ ∼ ∼ ∼ ∼ ∼ / Many of these lines were identified using the VALD service(Kupka et al. 1999). http://vald.astro.univie.ac.at/
Pecaut et al.perature subclass, assigning the luminosity class “IV” ifthe strength was intermediate between the dwarf and gi-ant, “IV-V” if the strength was very similar to that of adwarf but only slightly weaker, and “V” if the Na dou-blet was indistinguishable from a dwarf. The results ofour spectral classification are listed in Table 3.
Synthetic Colors
In order to compare observed colors to model at-mosphere predictions for the color locus and the pre-dicted effects of surface gravity, we compare our ob-served colors with synthetic colors calculated from the“BT-Settl” models from the Phoenix/NextGen group(Hauschildt et al. 1999; Allard et al. 2012) and the “AT-LAS9” models from Castelli & Kurucz (2004). The BT-Settl models offer synthetic spectra with 400
K < T eff < K , − . < log( g ) < . − . < [ M/H ] < +0 .
5, with α -element enhancement between +0.0 and+0.6 dex. The ATLAS9 models offer synthetic spec-tra with 3500 K < T eff < K , 0 . < log( g ) < . − . < [ M/H ] < +0 . α -element enhancement be-tween +0.0 and +0.4 dex. However, since our objectsare young and are in the solar neighborhood, we assumesolar metallicity with no α -element enhancement. Thisis consistent with the findings of Viana Almeida et al.(2009), who have spectroscopically analyzed a smallsample of these young stars, obtaining < [Fe/H] > =-0.06 ± ± β Pic member HD 322990.We computed synthetic colors, listed in Table 5, for so-lar metallicity models with with 3 . < log( g ) < . K < T eff < K for the BT-Settl models and3500 K < T eff < K for the ATLAS9 models, withno α -element enhancement. Pre-MS stars have lower sur-face gravities than main sequence stars at the same T eff but both should have 3 . < log( g ) < .
0. We wish toevaluate model predictions of color trends as a functionof surface gravity, so we plot synthetic colors for bothlog( g ) = 3 . V – K S . Wechose V – K S because it is available for nearly all ourobjects, and it offers a very large baseline compared toother colors so it is useful as a proxy for T eff . To com-pute the synthetic photometry for the models, we usethe updated BV I C normalized photonic bandpasses andzero points from Bessell & Murphy (2012), including theadditional zeropoints listed in their Table 5. To com-pute the 2MASS JHK S synthetic photometry, we usethe relative system response (RSR) curves available onthe IPAC website with the zero magnitude flux givenin Rieke et al. (2008). Similarly, for the WISE bands weuse RSR curves available on the IPAC website with thezero magnitude flux given in Jarrett et al. (2011). TheATLAS9 models are sparsely sampled past ∼ µ m, withonly 9 points representing the flux density from 10 µ mto 160 µ m, so we linearly interpolate λ F λ from 10 µ mto 160 µ m and divide by λ before performing synthetic http://wise2.ipac.caltech.edu/docs/release/prelim/expsup/sec4_3g.html photometry. This is not necessary for the BT-Settl mod-els because they are sampled at 0.2˚Aspectral resolutionfor λ > . µ m. The BT-Settl models shown adopt theAsplund et al. (2009) solar abundances while the AT-LAS9 models shown use the Grevesse & Sauval (1998)solar abundances. The computed synthetic colors arelisted in Table 5. Empirical Colors of Dwarfs Versus Pre-MS Stars
To compare dwarfs colors with pre-MS colors, we plotcolor-color diagrams for the young stars listed in Table 1.Figures 2 and 3 show V – K S versus B – V , V – I C , J – H , H – K S , K S – W K S – W K S – W K S – W B – V from Alonso et al. (1999) and for V – I C , J – H , and H – K S from Bessell & Brett (1988) con-verted to the 2MASS photometric system with the con-versions of Carpenter (2001). For reference we includethe BV I C solar colors estimated by Ram´ırez et al. (2012)and 2MASS JHK S and WISE W W W W V – K S versus B – V and V – K S versus J – H show the largest color difference between our youngstars and the dwarf locus. Redward of V – K S ∼ . B – V than the dwarf locus, andfor V – K S ≥ . B – V colors arebluer at lower surface gravity at a given V – K S , consis-tent with our observations, though the agreement is notperfect. Models predict little sensitivity to surface grav-ity for V – K S versus V – I C , consistent with the locationof the dwarf and giant locus as well as the placementof the young stars. For V – K S versus J – H locus, a bi-furcation between the dwarf and giant empirical locusoccurs at V – K S ∼ ∼ K5. This color split has been explained by themodels as an effect of surface gravity, due to the COand H O bands and H − opacity (Jorgensen 1996). Theyoung stars in our sample have surface gravities inter-mediate between that of the giants and dwarfs, and asa result they populate the region between the the dwarfand giant loci. For V – K S ≤ .
5, the young stars lie abovethe dwarf locus for colors H – K S and K S – W
1, indicatingthat these two colors are redder for young stars at a given V – K S . We exclude photometry for objects which havepreviously identified infrared excesses in that respectiveinfrared band, likely due to a dusty circumstellar disk.Excluded photometry is indicated in Table 1. Spectral Type-Color Sequence
To define the intrinsic color sequence empirically, withthe constraint of satisfying the color-color plots, we firstfit a spline to spectral type versus V – K S and spectraltype versus V – I C . We then verify that these relationsprovide a good fit to the V – K S versus V – I C color-colorrelation as well. We then fit splines to V – K S versus J – H and V – K S versus H – K S and use our spectral type- V – K S relation to anchor J – H and H – K S to spectral type.Finally, we fit splines to spectral type versus color for thecolors B – V , K S – W K S – W K S – W K S – W V – I C data is sparse for types earlier than G5, but appearsconsistent with the dwarf sequence, so we simply adoptoung Stellar Colors and Temperatures 5 Fig. 2.—
Comparison of B – V , V – I C , J – H , H – K S , K S – W
1, and K S – W V – K S of young stars from β Pic, η Cha, TWA andTuc-Hor moving groups (circles) with the dwarf color locus described in Appendix C and the giant color locus from Bessell & Brett (1988),except the B – V giant locus, which is from Alonso et al. (1999). Spectral types corresponding to the V – K S colors of dwarfs are plottedalong the top. Objects with a known near-IR or IR excess have been excluded (see Table 1). Pecaut et al.
Fig. 3.—
Same as Figure 2, except V – K S versus K S – W K S – W the dwarf V – I C sequence for spectral types F0 throughG5 discussed in Appendix C. In Figure 4 we see that pre-MS stars later than K3 become bluer in B – V than theirmain sequence counterparts, while those hotter than K2are nearly indistinguishable from main sequence stars.Figure 4 also shows that young stars G5 and later haveredder V – K S and J – H colors than field dwarfs, whilethose earlier than G5 have V – K S and J – H colors indis-tinguishable from field dwarfs. Pre-MS stars have H – K S colors redder than field dwarfs between spectral types F0and M2, shown in Figure 4. The spectral type sequencefor K S – W K S – W K S – W K S – W µ m excess. We find that it has a K S – W σ above the young color sequence. We alsoidentify HD 160305 and CD-54 7336 as having a K S – W σ and 5.4 σ above the young color sequence, so we also exclude them from the K S – W V – K S color 0.24 mag redder than a main-sequence K0. If oneestimated A V based on the stars E ( V − K S ) calculatedusing dwarf colors, then this star would appear to have A V =1.12 E ( V − K S ) ≃ V – K S color difference betweenpre-MS and a main-sequence K0 stars (assuming a stan-dard R V =3.1 reddening law). A 0.3 mag systematic shiftin H-R diagram placement would cause a 15 Myr old K-type star to erroneously appear 10 Myr old. Temperature Scale
Technique
The effective temperature (T eff ) scale for giants (e.g.,van Belle et al. 1999) as a function of spectral type is ∼ ∼ eff scale will be intermediatebetween dwarfs and giants (e.g., Luhman et al. 2003).All T eff scales depend on models (e.g., atmosphericmodels, limb-darkening models) to some degree. Ar-guably, the least model-dependent methods are those direct methods based on the angular diameter of thestar, measured interferometrically or by lunar occulta-tion methods. While some of the stars in our sam-ple are candidates for angular diameter measurements(see McCarthy & White 2012), only two have actualmeasurements in the literature (HR 9 and 51 Eri;Simon & Schaefer 2011; see Section 5 for details). Thereare also indirect methods, such as the infrared fluxmethod (IRFM), performed by Alonso et al. (1999), andmore recently for M-dwarfs by Casagrande et al. (2008),or directly fitting spectral energy distributions to syn-thetic model photometry, as described by Masana et al.(2006).Spectroscopically, young stars have been shown to ex-hibit more than one photospheric T eff (Gullbring et al.1998; Stauffer et al. 2003), so fitting synthetic spectra toobserved spectra will yield a different T eff depending onthe spectral region selected for fitting. An example of thisis TW Hydra, which has been consistently typed as a lateK star based on optical spectra (K7e, de la Reza et al.1989; K6e, Hoff et al. 1998; K6Ve, Torres et al. 2006;K8IVe, this work) but near-IR spectroscopy indicate aspectral type of M2.5V (Vacca & Sandell 2011). Weneed a method to infer temperatures that will simultane-ously take into account the observed optical-IR photom-etry. Therefore we attempt to infer the effective temper-atures by simultaneously fitting the observed photome-try to synthetic models (the “Spectral Energy Distribu-tion Fitting” (SEDF) method, see Masana et al. 2006).The downside of this method is that we are using modelswhich do not completely correctly predict the colors ofyoung stars. However, since the T eff is defined by theintegrated spectral energy distribution (SED) and thestellar radius, the observed photometry is the most di-oung Stellar Colors and Temperatures 7 F0 F5 G0 G5 K0 K5 M0 M5 L0 . . . . . . . B – V [ m a g ] HIP 47133
Dwarf ColorsYoung ColorsF0 F5 G0 G5 K0 K5 M0 M5 L0 . . . . . V – K S [ m a g ] AG Tri2MASS J11254754-4410267TWA 23
F0 F5 G0 G5 K0 K5 M0 M5 L0SpT . . . . . . H – K S [ m a g ] F0 F5 G0 G5 K0 K5 M0 M5 L0 . . . . V – I C [ m a g ] TWA 28RECX 15
F0 F5 G0 G5 K0 K5 M0 M5 L0 . . . . . . . . . . J – H [ m a g ] HD 164249BTWA 30B
F0 F5 G0 G5 K0 K5 M0 M5 L0SpT − . . . . . . K S – W [ m a g ] TWA 31 TWA 30ATWA 33TWA 34
Fig. 4.—
Comparison of B – V , V – I C , V – K S , J – H , H – K S , and K S – W β Pic, η Cha, TWA and Tuc-Hor movinggroups (circles) with the dwarf color sequence described in this work (dashed line). The outliers (filled squares) were excluded from the fit.
Pecaut et al.
F0 F5 G0 G5 K0 K5 M0 M5 L0 . . . . . . K S – W [ m a g ] TWA 31TWA 33TWA 34
Young ColorsDwarf ColorsF0 F5 G0 G5 K0 K5 M0 M5 L0 − . . . . . . K S – W [ m a g ] HD 222259 TWA 13ATWA 13B2MASS J11254754-4410267 HD 164249B
F0 F5 G0 G5 K0 K5 M0 M5 L0SpT − . . . . . . . . K S – W [ m a g ] HD 160305 CD-54 7336 AG Tri TYC 1766-1431-1
Fig. 5.—
Same as in Figure 4, except showing colors K S – W K S – W
3, and K S – W
4. Outliers (filled squares) have been excludedfrom the fit, and objects with known infrared excesses are notshown. rect link to the effective temperature of objects in oursample. We closely follow the formalism and methods ofMasana et al. (2006) and fit the observed photometry tomodels by minimizing χ , defined as χ = X i (cid:18) m i − m i, syn − A σ m i (cid:19) With m i = B , V , I C , J , H , K S , W W W
3, and W i, syn = B syn , V syn , I C syn , J syn , H syn , K S syn , W syn , W syn , W syn , and W syn are the synthetic apparent magnitudes at the stel-lar surface, and A is the magnitude difference betweenthe flux observed on Earth (obs) and the theoretical fluxat the surface of the star (surface) : A = − . F surface /F obs )related to the angular semi-diameter: θ = Rd = 10 − . A We fit the observed photometry to synthetic photom-etry from two different libraries of synthetic spectra:the BT-Settl models of Allard et al. (2012) with theAsplund et al. (2009) solar composition and the AT-LAS9 models of Castelli & Kurucz (2004) with theGrevesse & Sauval (1998) solar composition. The dif-ferences in the solar composition are particularly im-portant for low-mass stars and brown dwarfs, dueto the importance of TiO and VO in their spec-tra. The solar oxygen abundance was revised down-ward by 38% by Asplund et al. (2009) compared to theGrevesse & Sauval (1998) oxygen abundances. Anothermajor difference between the ATLAS9 models and theBT-Settl models is the treatment of line opacities. TheATLAS9 models include opacity distribution functions(ODFs) to account for line blanketing, whereas the BT-Settl models are generated by the PHOENIX code inwhich the individual contribution of atoms and moleculesis directly sampled over all computed points in the spec-trum (Hauschildt et al. 1997). Given that the BT-Settlmodels offer continuity in our ability to model SEDs of F-type down to M-type stars, and the recent successes theBT-Settl models have had fitting NIR colors of low-massstars down to ∼ Testing Technique on Objects with Measured AngularDiameters This flux is the unresolved flux integrated over the disk of thestar and does not represent the resolved flux one would observe ifplaced on the stellar surface. The flux we are referring to is thecounterpart to the apparent magnitude at the stellar surface (e.g., B syn ). http://phoenix.ens-lyon.fr/Grids/BT-Settl/AGSS2009/ oung Stellar Colors and Temperatures 9As a reliability check for the usefulness of ourmethod, we use the estimated solar BV I C colors fromRam´ırez et al. (2012) together with the solar 2MASS JHK and WISE W W W W eff , as-suming log( g ) = 4 .
44 and adopting the apparent V bandmagnitude of -26.74 ± eff ⊙ =5776 ±
22 K (remarkably within 4 K of themodern solar T eff of 5771.8 ± ′′ ± ′′ . The ATLAS9 mod-els give T eff ⊙ =5737 ±
21 K, 35 K too low but still within2 σ , and an angular diameter of 1953 ′′ ± ′′ . Both an-gular diameter measurements are systematically higherthan the 1918.3 ′′ ± ′′ angular diameter implied by thesolar radius estimate of Haberreiter et al. (2008), whichstrongly suggests that our adopted V ⊙ is too high. If weinstead adopt V ⊙ ≡ -26.71 ± ′′ ± ′′ and1926 ′′ ± ′′ with the BT-Settl and ATLAS9 models, re-spectively, consistent with the modern solar angular di-ameter estimates. Thus for consistency with the solarvalues, also consistent with the Engelke et al. (2010) syn-thetic solar V ⊙ , we adopt V ⊙ =-26.71 ± .We also check our technique on nearby K- and M-typefield dwarfs with directly measured angular diametersfrom the recent work of Boyajian et al. (2012b). Weuse photometry from Table 7 of Boyajian et al. (2012b),converting Johnson I to the Cousins system using theconversions in Bessell (1979) and converting Johnson JHK to the 2MASS system using the conversions ofCarpenter (2001). We adopt WISE W W W g ) = 4 . σ log( g ) =0.2 dex and σ [ m/H ] =0.1 dex. OurSEDF-derived T eff for these stars are listed in Table 7,and plotted with the Boyajian et al. (2012b) T eff valuesin Figure 10. The mean difference between our SEDF-derived T eff values and those based on angular diametermeasurements from Boyajian et al. (2012b) is 13 K witha 1 σ dispersion of 108 K. We conclude that our techniqueworks well for the Sun and nearby dwarfs with angular di-ameter measurements, and gives us some confidence thatthis method will accurately predict the effective temper-atures of our pre-MS stars. Analysis
For many objects in our sample, one or more bands ofphotometry are not available. In those cases we simplyomit the term containing the missing band data. We donot fit bands with poor quality photometry (in 2MASS,anything other than quality flag ‘A’; for WISE bands,anything other than contamination and confusion flag‘0’). We have again excluded photometry for objectswith infrared excesses, flagged in Table 1. RECX 11 andRECX 15 have K S -band excesses, so we exclude them V ⊙ =-26.71 ± V, ⊙ =4.862 ± ± × erg s − ) leads toM bol, ⊙ =4.7554 ± V, ⊙ =-0.107 ± V magnitude estimates is available at https://sites.google.com/site/mamajeksstarnotes/basic-astronomical-data-for-the-sun from SED fitting entirely. We also exclude TWA 30Adue to its time variable extinction (Looper et al. 2010b)and TWA 30B due to the time variable near-infraredexcess (Looper et al. 2010a). TWA 31, TWA 33 andTWA 34 have W W W W W
3, and W JHK S photometry, so we exclude them en-tirely. TWA 29 had only 2MASS JHK S photometry,and HD 139084B and HD 164249B had 2MASS photom-etry and only two bands of WISE photometry with largeuncertainties ( > σ T eff >
300 K) so we ex-cluded them from SED fitting as well. Objects excludedfrom SED fitting are listed in Table 8. The behaviorof χ as a function of T eff is consistent with Gaussianerrors and χ has a quadratic dependence on T eff nearthe best-fit value. A representative SED from our sam-ple with the observed and best-fit model are shown inFigure 6.In general the synthetic photometry is a function oflog( g ), T eff , and metallicity ([m/H]). As discussed previ-ously, we use solar metallicity synthetic models. Pre-main sequence evolutionary tracks from Baraffe et al.(1998) between 8-30 Myr predict that log( g ) varies be-tween 4.1 dex and 4.5 dex so we simply adopt 4.3 ± eff and log( g )simultaneously, this often gives spuriously large or smalllog( g ) values, and even when the values of log( g ) ob-tained from the fit are within an expected range, they arenot well-constrained (e.g., formal errors on log( g ) ∼ . g ), and fur-thermore, we found that the best-fit T eff did not varysignificantly between log( g ) = 4 . g ) = 4 .
5. Themean difference in T eff between log( g ) =4.1 and 4.5 is4 K with a dispersion of 31 K. Therefore, in our fittingprocedure we set T eff as the only free parameter. Duringthe fitting procedure, we first determine A as the inverse-variance weighted mean difference between the observedand synthetic photometry at the stellar surface. How-ever, rather than numerically minimizing χ (as done inMasana et al. 2006) we simply find the minimum valueover our grid, interpolated to T eff increments of 20 Kfrom 1400 K to 9800 K for the BT-Settl models and from3500 K to 9750 K for the ATLAS9 models. We then fita parabola in the region surrounding the minimum. Results
The effective temperatures from the SEDF techniqueare listed in Table 9. We estimate our uncertainties byperforming a Monte Carlo simulation. For each object,we select trial photometry values from a distribution withmean and standard deviation equal to the observed pho-tometry value and uncertainty, and use the trial pho-tometry values to obtain the best-fit T eff and angulardiameter estimate. We perform 300 trials for each objectand use the standard deviation of the resulting T eff andangular diameter distribution as our statistical uncer-tainties. However this does not account for systematicscaused by uncertainties in our assumed surface gravityand metallicity. To account for these systematics, we re-peat our fitting procedure for each object, varying log( g )from 4.1 dex to 4.5 dex and [m/H] from +0.2 dex to -0 Pecaut et al. Fig. 6.—
SED for β Pic member V1005 Ori (K8IVe). Observedphotometry (circles) with the best fit BT-Settl model photometry(crosses) from a T eff =3866 ±
18 K model (interpolated). Uncertain-ties are smaller than the symbol markers. eff and angular di-ameter obtained for the systematic uncertainty, typically ∼
11 K in T eff and ∼ µ as in angular diameter. Theuncertainties quoted in Table 9 are the statistical and(internal) systematic uncertainties added in quadrature.This does not account for any systematic uncertaintiesfrom the underlying Phoenix/NextGen models or the as-sumed solar abundances.Similar to other studies, we find that V – K S providesthe closest correlation to temperature with relatively lit-tle scatter. To take advantage of the utility of V – K S as aproxy for T eff , we estimate the spectral type-temperaturecalibration by fitting a polynomial to T eff as a function of V – K S . The coefficients for this polynomial are listed inTable 10. We then apply this polynomial to our spectraltype-intrinsic color sequence. Unfortunately only one ob-ject in our sample later than spectral type M5.5 has V band photometry, so we do not provide effective tem-perature estimates for spectral types M6-M9, though wedo provide intrinsic colors for those spectral types. Ourspectral type, intrinsic color and T eff sequence for youngstars is listed in Table 6. For comparison, in Figure 7 wehave plotted the new temperature scale for 5-30 Myr pre-MS stars described in this work, the giant temperaturescale of van Belle et al. (1999), a new “consensus” dwarfT eff scale described in Appendix C, and the young starscale of Luhman et al. (2003) (appropriate for ∼ eff scale is within ∼
100 K of the dwarf scale as a functionof spectral type, except for spectral types G5 throughK6, which are ∼
250 K cooler than their main-sequencecounterparts.
Bolometric Corrections
As a byproduct of estimating the effective temperatureof stars in our sample using the method of SED fitting,we also obtain an estimate of each object’s angular di-ameter. This can then be used to estimate the apparentbolometric magnitude (m bol ) and the bolometric correc-tion in any band ( BC x ). The basic equation that relates F0 F5 G0 G5 K0 K5 M0 M5SpT T e ff [ K ] Pre-MS (this work)Dwarfs (this work)Giants (van Belle et al. (1999)Luhman et al. (2003)
Fig. 7.—
Spectral type versus T eff for the pre-MS (solid blackline) and dwarf (solid grey line) T eff scales derived in this work.For comparison we plot the giant T eff scale of van Belle et al.(1999) (dotted line) and the T eff scale of Luhman et al. (2003)(dashed line), appropriate for ∼ eff scale is within 100 K of the dwarf scale as a function of spectraltype, except for spectral types G5 through K6, which are ∼
250 Kcooler than their main-sequence counterparts. stellar bolometric magnitude to luminosity is M bol = − . (cid:18) LL ⊙ (cid:19) + M bol , ⊙ = −
10 log (cid:18) T eff T ⊙ (cid:19) − (cid:18) RR ⊙ (cid:19) + M bol , ⊙ . We can also write this in terms of apparent magnitude m x in band x with the distance d and bolometric correc-tion BC x : M bol = m x − (cid:18) d pc (cid:19) + BC x . Equating these two, using the angular semi-diameter θ = Rd = 10 − . A , and solving for BC x we find BC x = A + 5 log (cid:18) R ⊙ pc (cid:19) + M bol , ⊙ −
10 log (cid:18) T eff T eff , ⊙ (cid:19) − m x . We use consistent solar values of T eff , ⊙ = 5772 K , R ⊙ = 695660 km, m V, ⊙ from Section 4.5.2, and M bol , ⊙ =4 .
755 mag as adopted by Mamajek (2012) . The uncer-tainties in BC x are( σ BC x ) = (cid:18) σ T eff T eff ln 10 (cid:19) + ( σ A ) + ( σ m x ) . See also “Basic Astronomical Data for the Sun”, https://sites.google.com/site/mamajeksstarnotes/basic-astronomical-data- more complete discussion on solar data, including motivation forthe values adopted here. oung Stellar Colors and Temperatures 11
Fig. 8.—
Bolometric corrections for V and J band magnitudesas a function of effective temperature. Note that for T eff < ∼ BC V becomes a sensitive function of T eff and therefore it is prefer-able to use M bol = M J + BC J for cooler stars. Coefficients forpolynomial fit are listed in Table 10. Though the zero point of the bolometric correctionscale is arbitrary, the combination of bolometric cor-rection and solar absolute bolometric magnitude is not(see Torres 2010 and Appendix D of Bessell et al. 1998).In Table 9 we give the calculated individual bolomet-ric corrections in both Johnson V band and 2MASS J band. We also provide log(L/L ⊙ ) for stars with measuredtrigonometric parallaxes. For F- and G-type stars (T eff > ∼ M bol = M V + BC V since the V band correctionis not a sensitive function of T eff for 5000 K < T eff < eff < ∼ BC V becomes a steep function of T eff and therefore itis better to use M bol = M J + BC J . Plots of BC V and BC J versus T eff are shown in Figure 8. Polynomial fitsto BC V and BC J as a function of T eff and V − K S aregiven in Table 10. DISCUSSION
Consistent with previous studies (e.g., Da Rio et al.2010; Luhman et al. 2010b), we have found that pre-MS stars do not have the same intrinsic colors as fielddwarfs for certain spectral types and colors. There aretwo likely main reasons for the differences in colors. Thefirst and most important cause is the different surfacegravities of pre-MS stars compared to main sequencedwarfs. The striking bifurcation in the V – K S versus J – H color-color diagram between dwarfs and giants hasbeen explained as an effect of CO and H O bands andH − opacity (Jorgensen 1996). The B – V colors for pre-MS stars with V – K S > . B – V colors are predicted to be bluer ata given V – K S than higher-surface gravity stars. Ournew spectral type-color relations take these importantsurface-gravity effects for young stars into account. How-ever, this does not explain the origin of redder colors,particularly H – K S , for F- and G-type stars, which havesurface gravities very close to main sequence dwarfs.The second possible explanation for color differencesbetween young stars and older main sequence starssuggested by Gullbring et al. (1998) and Stauffer et al.(2003) is the greater abundance of stellar spots on youngstars. Young stars show evidence of stronger mag-netic activity than older main sequence stars, which isexhibited by hotter plage and cooler spot regions onthe surface. In particular, these plage regions havebeen suggested as contributing to the systematicallybluer B – V colors observed in the Pleiades open cluster(Stauffer et al. 2003). Gullbring et al. (1998) estimateda ∼
50% spot coverage to account for the mean V – J color anomaly in weak-lined T Tauri stars. However, theStauffer et al. (2003) study is the most comprehensiveattempt to date to investigate the contribution of coolspots to stellar colors. Stauffer et al. (2003) found thatPleiades K star red spectra (5700-8400˚A) had system-atically later spectral types than the blue (3300-5300˚A)spectra, whereas the older Praesepe K stars did not suf-fer from this effect . Stauffer et al. (2003) addition-ally modeled the BV RIJHK
SEDs of several Pleiades,combining observed SEDs of an earlier field dwarf anda later field dwarf to obtain a fit. The best-fit mod-els obtained in the Stauffer et al. (2003) study indicatedthat the K-type Pleiades were covered in > ∼
50% “coolspots”, consistent with the Gullbring et al. (1998) re-sults. They use
BV RIJHK photometry to fit observedPleiad SEDs. On the basis of their spectroscopy and SEDfitting, Stauffer et al. (2003) concluded that the PleiadesK stars had more than one photospheric temperature,and that spottedness was well-correlated with the B – V color anomaly. While these results point convincinglyto stellar spots as a significant contributing factor, espe-cially to bluer B – V colors, we do not attempt to quantifythe relative contribution of spots or surface gravity effectsto the intrinsic colors of pre-MS stars. Disentangling theeffects of surface gravity and spots would require time-series multi-band photometry for most of the objects inour sample. Quantifying the specific contribution of thespots and plages to the stellar colors is beyond the scopeof this study.McCarthy & White (2012) published predicted angu-lar diameters for many of the β Pic moving group mem-bers in our sample using estimated H-R diagram po-sitions and revised
Hipparcos parallaxes (van Leeuwen2007). In addition, Lafrasse et al. (2010) have estimated This effect is also seen in G and K stars from the youngerScorpius-Centaurus OB association, where blue spectra ( ∼ ∼ Fig. 9.—
The individual angular diameter estimates from thiswork compared with estimates from McCarthy & White (2012)and Lafrasse et al. (2010). the angular diameters of thousands of dwarfs and gi-ants with V and V – K surface brightness relations (e.g.,Barnes & Evans 1976). We compare our results to theMcCarthy & White (2012) and Lafrasse et al. (2010) re-sults in Figure 9. Our angular diameter estimates followthe Lafrasse et al. (2010) estimates very closely, thoughours are systematically smaller by 4%. There is a trendwith T eff , with hotter objects tend to be more dis-crepant than cooler objects, however, the origin of thisdiscrepancy is unclear. Our angular diameter estimatesalso compare well with the results of McCarthy & White(2012), with our estimates being 6% larger on average,but with much larger scatter, however, this differenceis not statistically significant. The larger scatter be-tween our angular diameter estimates and those fromMcCarthy & White (2012) are likely due to the differentmethods used to infer the stellar T eff s. For example, wepredict TYC 1208-468-1 to have a diameter of 241 ± µ as,but McCarthy & White (2012) predict 120 µ as. Thisstar has BV JHK colors consistent with a spectral typeof ∼ K6, but it has a reported spectral type of K3Ve(Jeffries 1995). The ∼
600 K difference in the assumedT eff translates to a large difference in the predicted an-gular diameter.There is considerable overlap between our sample andthe sample of Mentuch et al. (2008), who examined Lidepletion in several nearby young associations. TheMentuch et al. (2008) study estimated T eff for each starin their sample by least-squares fitting synthetic spec-tra to spectral regions λλ eff values obtained by fitting multi-band photometry tothe BT-Settl NextGen model colors with the T eff val-ues obtained by Mentuch et al. (2008). Overall thereis a systematic difference – the values obtained byMentuch et al. (2008) are systematically ∼
150 K hot-ter than the values we obtain, with a larger difference( ∼ ∼
120 K)below 4500 K. This discrepancy could be due to the dif- ferent synthetic models used. The latest BT-Settl mod-els use the revised solar abundances from Asplund et al.(2009) and include more complete molecular opacity lists,though these updated opacities would mostly affect thelower-mass stars and are unlikely to account for the dif-ferences above ∼ eff values with those of Casagrande et al. (2008) andCasagrande et al. (2011), where possible (Figure 10).Both studies used synthetic spectra with an implemen-tation of the Infrared Flux Method (IRFM) or a closelyrelated method (Multiple Optical Infrared TEchniqueor “MOITE”) to estimate the stellar effective temper-ature for a large number of objects. The IRFM com-pares the ratio of the observed bolometric flux to the ob-served monochromatic flux density at the Earth (“R obs ”)to the ratio of theoretical bolometric flux to monochro-matic flux density at the surface of the star (“R theo ”)(Blackwell & Shallis 1977). R theo is a function of theT eff , and is compared to the R obs ratio to obtain theT eff of the star. For hotter stars the sensitivity tothe model in the IR is very minimal and thus onlythese flux ratios in IR bands are used to determine theT eff . For cooler stars, Casagrande et al. (2008) haveadapted this method to use optical and infrared bands(called “MOITE”). Casagrande et al. (2008) assumed log ( g )=5.0 dex throughout with the ‘Cond’ variant ofNextGen models (we have used the ‘BT-Settl’ varianthere with revised solar abundances from Asplund et al.2009), whereas the Casagrande et al. (2011) study usedthe Castelli & Kurucz (2004) models which used theGrevesse & Sauval (1998) solar abundances. Stellar T eff estimates from this work are typically ∼
40 K lower thanthe values from the Casagrande et al. (2011) study (sixstars in common), and within 2 σ of the values fromthe Casagrande et al. (2008) study (stars TX PsA andHIP 107345 in common). A comparison of stellar T eff estimates from this work and the literature is shown inFigure 10.For the few objects with spectral types M8 or laterwe obtain cooler temperatures than expected from thetemperature scale of Luhman et al. (2003) or the dwarftemperature scale. Rice et al. (2010a) fit PHOENIXdusty synthetic spectra to high-resolution observed spec-tra to find the best-fit T eff and log( g ) of sample ofyoung late M-type objects. Two objects in our samplewith SEDF-determined T eff , 2MASS J06085283-2753583(2M0608-27; M8.5 γ ; Rice et al. 2010b) and TWA 26(M8IVe; Barrado Y Navascu´es 2006), are included inthe Rice et al. (2010a) study. For 2M0608-27, assum-ing log( g )=4.3 dex, we find T eff =2118 ±
20 K, whereasRice et al. (2010a) adopt log( g )=3.98 and T eff =2529 K,much hotter than our results and consistent with thetemperature scale of Luhman et al. (2003). We findT eff =2176 ±
17 K for TWA 26 but Rice et al. (2010a)find log( g )=3.98 and T eff =2609 K, again much hotterthan our results and consistent with the T eff scale ofLuhman et al. (2003). These four objects lack BV I C photometry and thus do not have any SED fitting con-straints blueward of their SED peak; this could be a con-tributing factor in their discrepantly cool T eff fit. Be-cause of these discrepancies, we do not include T eff es-timates for M6 through M9 objects in our pre-MS tem-oung Stellar Colors and Temperatures 13 Fig. 10.—
The individual T eff values from this work comparedwith values obtained by least-squares fitting to synthetic NextGenspectra from Mentuch et al. (2008) (crosses) and those in thestudy of Casagrande et al. (2011) (triangles) and Casagrande et al.(2008) (stars). We also compare a sample of K- and M-typedwarfs which have angular diameter-based T eff estimates fromBoyajian et al. (2012b) with estimates using our SEDF implemen-tation (circles). The values in Mentuch et al. (2008) are systemat-ically higher than those estimated in this work, with the difference ∼
230 K above 4500 K, reducing to ∼
120 K below 4500 K. Thosefrom the Casagrande et al. (2011) study are typically ∼
40 K higherthan the values from this work. perature scale (Table 6).Only two stars in our sample, HR 9 and 51 Eri, havedirect angular diameter measurements available fromthe literature. Simon & Schaefer (2011) measured angu-lar diameters of 492 ± µ as and 518 ± µ as for HR 9and 51 Eri, respectively. Our angular diameter esti-mates of 346 ± µ as for HR 9 and 471 ± µ as for 51 Eriare much lower than the interferometric measurements.While we have no reason to suspect the direct mea-surements are unreliable a priori , the angular diame-ter of 492 µ as for HR 9 (F3Vn; Gray et al. 2006) war-rants some discussion. If we adopt the estimated bolo-metric flux at Earth for HR 9 from Casagrande et al.(2011) of 8.6609 × − mW m − , and again use the mea-sured angular diameter of 492 µ as, we obtain a T eff of5724 K, similar to a G3V! This is ∼ ±
91 K estimated by Casagrande et al. (2011)and our estimate of 6761 ±
28 K, both of which are consis-tent with the F3Vn spectral type. The Simon & Schaefer(2011) results indicate a larger angular diameter at H band than K band, which points to unusual calibrationerrors (M. Simon, private communication 2012). We sus-pect that our predicted angular diameters are closer tothe actual diameters and until updated measurementsare published, we recommend our predicted angular di-ameter. CONCLUSIONS
We can summarize our conclusions as follows:1. 5-30 Myr old pre-main sequence stars followslightly different spectral type-intrinsic color se-quences than that of main sequence stars. Pre-MS colors follow the dwarf sequence for some colors andspectral types, but for other optical/infrared colorsand spectral types, deviations can exceed 0.3 mag.In Table 6 we provide an empirical tabulation of theintrinsic colors of young stars for spectral types F0through M9, including B - V , V – I C , V – K S , J – H , H – K S , K S – W K S – W K S – W
3, and K S – W J – H colorsand bluer B – V colors for lower surface gravity ob-jects, consistent with observations. However, wecannot exclude hotter plage and cooler spot regionson the stellar surface as contributing factors.3. A pre-MS T eff scale derived from fitting SEDs tosynthetic spectral models is within ∼
100 K of mainsequence stars as a function of V – K S . As a func-tion of spectral type, the effective temperatures ofF0 through G4 and K7 through M5 pre-MS starsare within ∼
100 K of their main sequence coun-terparts, whereas G5 through K6 pre-MS stars are ∼
250 K cooler at a given spectral type. We providenew spectral type-T eff relations and color-T eff re-lations appropriate for 5-30 Myr old pre-MS stars.We also provide bolometric corrections appropri-ate for PMS stars as polynomial functions of T eff and V – K S in Table 10 and as part of our spectraltype-intrinsic color sequence in Table 6.We thank Fred Walter for the use of his IDL reduc-tion pipeline and the Stony Brook Spectral StandardsLibrary. We also thank the referee, Kevin Luhman, forhis very thorough and prompt review which greatly im-proved the paper. Spectra taken for this study wereobserved on the 1.5-m telescope on Cerro Tololo viathe Small and Moderate Aperture Research TelescopeSystem (SMARTS) Consortium. We thank Duy Nyu-gen for helpful discussions regarding χ fitting. Thiswork was supported by funds from NSF grant AST-1008908. 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MEMBERSHIP OF TWA 9 TO THE TW HYA ASSOCIATION
The membership of TWA 9 to the TW Hya Association merits some discussion. Weinberger et al. (2012) showedthat the space motion of TWA 9A is more than 3 σ from the mean of the association, and concluded that it was eithernot a member or the Hipparcos distance is underestimated. However, when considering the TWA centroid space motion(Weinberger et al. 2012), the Tycho-2 proper motion ( µ α ∗ =-55.4 ± − , µ δ =-17.7 ± − ; Høg et al.2000) of TWA 9A seems consistent with membership in TWA. Assuming it is a member and adopting the TWAmean group space motions from Weinberger et al. (2012) of (U, V, W) = (-10.9 ± ± ± − , weestimate a kinematic distance of 70.0 ± ± − , the3-D space motion of TWA 9A is then (U, V, W) = (-10.2 ± ± ± − . This is consistent with themean TWA space motions in the Weinberger et al. (2012) study. Furthermore, the kinematic distance would decreasethe absolute magnitude M H by ≈ Hipparcos distance (using d= π − , where π is the trigonometric Hipparcos parallax), and thus the isochronal age of TWA 9A would be ∼
16 Myr, much closer to the isochronal agesobtained by Weinberger et al. (2012) for other TWA members. TWA 9A exhibits very high Li (EW(Li 6708˚A)=470 m˚A;Torres et al. 2006), lies in the direction of other TWA members, has proper motion consistent with membership inTWA, and, adopting the kinematic distance of 70.0 ± Hipparcos parallax is most likely ∼ σ in error. SPECTRAL TRANSITION FROM K7 TO M0
Some spectral surveys implicitly or explicitly do not recognize or use spectral types K8 and K9. While spectral typesK8 and K9 are not considered full subtypes of the spectral classification system (Keenan 1984), it should be pointedout that neither are G1, G3, G4, G6, G7 or G9, yet these classifications are consistently recognized and used (e.g.,Gray et al. 2003). Keenan (1984) noted that subdivisions such as G3 simply means the star is closer to G2 than G5,and that they should be used when it is possible to classify the stars accurately enough to justify their use. Keenan(1984) considered K5 and M0 one subtype apart even though the difference in their B – V color is 0.3 mag, largerthan the difference between M0 and M4 (see Table 4). From the standpoint of spectral classification, there is nothingdifferent about the K7 to M0 transition that merits such a gap in spectral types. Therefore we find no compellingreason to omit spectral types K8 and K9 from use and we include them here in our analysis.With low-resolution red optical spectra we can distinguish between subtypes K7, K8, K9 and M0. Unfortunately,K8V and K9V spectral standards do not appear in the literature (e.g., Gray & Corbally 2009). For these subtypes,we adopted stars as standards which were assigned this classification by an expert classifier. For K8V we adoptedHIP 111288 ( V – K S =3.52 ± V – K S =3.70 ± V – K S colors intermediate between K7 and M0. We visually compared the spectra of both adoptedstandards and verified that they were morphologically intermediate between the K7V and M0V standards in Table 2.Figure 11 shows the red spectral sequence from K7 to M0. The spectra show a distinct progression in the CaHband at λλ ∼ Fig. 11.—
The spectral transition from K7 to M0, with regions most useful for discriminating among the different spectral types highlightedin grey. GJ 673 (K7V), HIP 111288 (K8Vk), HIP 3261 (K9V) and GJ 701 (M0.0V) are shown.
DWARF COLORS AND TEMPERATURES
In order to accurately compare the empirical intrinsic colors and T eff scale of pre-MS stars to dwarfs and quantifytheir differences, the empirical intrinsic colors of dwarf stars must be accurately tabulated. Here we describe theconstruction of a modern dwarf color, T eff , and bolometric correction sequence, which has been an ongoing processcarried out over several years. While other compilations are available (e.g. Schmidt-Kaler 1982; Kenyon & Hartmann1995; Worthey & Lee 2011), it was our goal to incorporate a detailed review of the color/temperature placement ofmodern spectral standard stars and assess their pedigree as standards. A preliminary version of this sequence (A0V-G9V) was previously published in Pecaut et al. (2012). The primary motivations for constructing this sequence werethat (1) color sequences over the O-M range of spectral types had not been constructed explicitly including 2MASS andWISE bands, (2) the methodology of the construction of previous sequences was not always made clear, (3) systematicdifferences exist between some of the widely cited past sequences, (4) there have been sizable shifts in T eff s reportedfor some stars (especially among the hottest and coolest dwarf stars) over the past few decades, and (5) there havebeen subtle changes to the dwarf spectral sequence over the decades, especially among the M dwarfs. In light of thesubtle shifts of the MK system over the past decades, improvements in the modeling of stellar atmospheres, and giventhe large volume of optical-IR photometry and derived stellar parameters in the literature now, a reevaluation of thetemperature and colors scales is overdue.We present our modern intrinsic color-T eff -spectral type tabulation for dwarfs in Table 4. This color tabulationwas independently derived, and is not dependent on previous compendia of dwarf photometric properties. There wereseveral stages that went into assembling Table 4. When discussing samples of “nearby stars”, we assumed that starswith trigonometric parallax distances within 75 pc had negligible reddening (e.g. Reis et al. 2011), and so could beused to estimate intrinsic colors. While we often quote the intrinsic stellar colors to 0.001 mag precision (to ensureconstruction of smooth sequences on color-color plots), the uncertainty in the mean colors is typically at the ∼ s colors for M dwarfs. Spectral Standard Stars
Spectral standard stars for stars of F-type and earlier were mostly drawn from Johnson & Morgan(1953); Morgan & Hiltner (1965); Garrison (1967); Lesh (1968); Abt et al. (1968); Hiltner et al. (1969);Cowley et al. (1969); Garrison (1972); Cowley (1972); Morgan & Keenan (1973); Cowley & Fraquelli (1974);Houk & Cowley (1975); Garrison et al. (1977); Morgan et al. (1978); Garrison & Schild (1979); Gray & Garrison(1987); Walborn & Fitzpatrick (1990); Garrison & Gray (1994); Garrison (1994); Gray & Corbally (2009). For M-typestars, the primary sources of standard stars were Kirkpatrick et al. (1991, 1997); Henry et al. (2002). Some M-typestandards from Keenan’s papers (e.g. Keenan & McNeil 1976; Keenan & Pitts 1980; Keenan 1983; Keenan & Yorka1988; Keenan & McNeil 1989) that have conflicting types compared to the newer classifications by Kirkpatrick, Henry,and collaborators, have been deprecated (e.g. GJ 15A, 172, 250B, 526) and were not considered in assessing mediancolors and T eff . Given the immense volume of recent M-star classifications that have been done on the Kirkpatrick &Henry grid (e.g. Reid et al. 1995; Hawley et al. 1996; Henry et al. 2002), these should be preferred to the Keenan types8 Pecaut et al.where there is disagreement. Classifications of AFGK field dwarfs by Gray et al. (2003, 2006) were generally preferredover those of the Michigan Atlas (Houk & Cowley 1975), as it appears that the Gray et al. classifications more closelyfollow the Morgan-Keenan standards. Differences between Gray et al. and Houk et al. classifications are especially pro-nounced amongst the early G-type stars. Part of this may stem from disagreement between Morgan and Keenan on theF/G boundary (e.g. see the example of η Cas A previously mentioned). More problematically, Houk & Cowley (1975)considered β Com to be their main G2V standard, but it was considered to be a G0V standard by Johnson & Morgan(1953); Morgan et al. (1971); Morgan & Keenan (1973); Keenan & McNeil (1976). This appears to explain why themedian B-V color for nearby G2V stars in the Hipparcos catalog (dominated by Michigan Atlas classifications) is B-V ≃ ≃ ± Assessing the Pedigree of Spectral Standard Stars
An extensive literature search was conducted to assemble notes on the published classifications and col-ors for all known O- through M-type dwarf spectral type standard stars (The notes have been compiled at and will be periodically updated as needed). All the dwarfspectral standards were assessed for continuity in their spectral classifications over the decades, and standards weregraded as “anchor standards” (Garrison 1994), “primary standards”, “secondary standards”, “tertiary standards”,“variant standards”, or “deprecated standards”. Our terminology is a variation on the hierarchy scheme of Garrison(1994), and the goal of assessing the pedigree of the various spectral standards was to help in the estimation of thebest stellar parameters reflective of a given spectral subtype. While the grading of the individual standards is notprovided here, the reader is referred to the website mentioned. “Anchor standards” are those rare standard stars listedby Garrison (1994) whose spectral types have remained unchanged since Morgan et al. (1943), and which essentiallydefine the MK system. “Primary standards” typically showed very strong continuity in adopted spectral types amongexpert classifiers, often going back to Johnson & Morgan (1953). “Secondary standards” usually appeared severaltimes in the literature as spectral standards, but sometimes expert classifiers assigned slightly different spectral typesto the star (usually at the ± η Cas A, considered anF9V standard by Keenan & Yorka (1988); Keenan & McNeil (1989) and Gray et al. (2001b), but considered a G0Vstandard by Morgan et al. (1943); Johnson & Morgan (1953); Morgan & Keenan (1973); Keenan & McNeil (1976);Morgan et al. (1978); Keenan (1983), and Keenan & Yorks (1985). Another example is σ Boo (HR 5447), which wasconsidered a F2V standard by Morgan et al. (1943) and Johnson & Morgan (1953), but two later studies found thestar to appear spectrally metal poor (F3V vw; Barry 1970) and (F4V kF2 mF1; Gray et al. 2001b). Use of suchstandards should probably be avoided in the future, if possible.While estimating the parameters for a given dwarf spectral subtype, more weight was assigned to the individualparameters (e.g. colors, T eff s) of the anchor and primary standards compared to the secondary and tertiary standards,and the properties of the variant and deprecated standards were largely ignored. While estimating the typical propertiesof non-standard stars of a given spectral subtype, we employed median values throughout, in order to avoid the effectsof interloper data (Gott et al. 2001). The properties of both standard and non-standard stars were incorporated intoestimation of typical colors and T eff s, and their properties usually agreed well with very few exceptions (e.g. B7V,where the lone good standard star HD 21071 appears to be significantly bluer and hotter than the majority of fieldstars classified B7V). Color Sequences
The intrinsic ( B – V ) o and ( U – B ) o colors can be derived for OB dwarfs via the Q -method (e.g., Johnson & Morgan1953; Johnson 1958; Hiltner & Johnson 1956), where the reddening-free index Q is calculated using the observedcolors as Q =( U – B )-0.72( B – V ). Functions of ( B – V ) o and ( U – B ) o as linear functions of Q , especially those thatare forced through the origin (( B – V ) o , ( U – B ) o ), produce poor fits to the colors of real unreddened OB stars. Wecalibrated new Q versus intrinsic color relations using U BV photometry from Mermilliod & Mermilliod (1994) ofnearby negligibly reddened B-type dwarfs within 75 pc (
Hipparcos catalog; ESA 1997), and lightly reddened hotter Oand early-B luminosity class V and IV stars in nearby associations. The more distant OB stars were dereddened usingoung Stellar Colors and Temperatures 19published H I column densities (e.g., Fruscione et al. 1994) and the strong correlation between N(H I) and E ( B – V );Diplas & Savage (1994). The improved Q-method fits are:( B − V ) o = − . × − +0 . Q +1 . Q +1 . Q +0 . Q for − . < ( B − V ) o < .
02, and ( U − B ) o = 6 . × − +1 . Q +1 . Q +2 . Q +1 . Q for -1.13 < ( U – B ) o < B – V ) o colors of O9/B0 dwarfs are -0.32 to -0.31 (among thecalibrator stars e.g. 10 Lac, σ Sco, τ Sco, and υ Ori), in agreement with Johnson’s classic work (e.g. Johnson & Morgan1953; Johnson 1966, but at odds with the recent work of Martins & Plez (2006) who claim that ( B – V ) o colors ofGalactic O stars go no bluer than -0.28.Deriving the main sequence color sequence was fairly straightforward. Photometry for nearby stars came from thefollowing sources: U BV
Mermilliod (1991),
BV I C (ESA 1997), JHK S (Skrutskie et al. 2006), W W W
3, and W V – K S vs. B – V , V – I C , J – H , H – K S , B – V vs V – I C and U – B ). We fit polynomial relations to V – K S versus K S – W K S – W K S – W
3, and K S – W Hipparcos catalog and the catalog of bright M dwarfs fromL´epine & Gaidos (2011). We adopted V magnitudes from the APASS Data Release 6 catalog (Henden et al. 2012) forobjects not present in the Hipparcos catalog, and only fit objects with high quality photometry in the relevant band(for 2MASS bands, quality flag ‘A’; for WISE bands, contamination and confusion flag ‘0’). We restricted the data toWISE magnitudes W > . W > . W > . W > . V – K S < . V – K S > . Effective Temperatures
Subtype T eff s were estimated by considering published T eff s for individual stars of a given subtype, though greaterweighting was given to T eff values for spectral standards which were vetted for consistent classifications in the lit-erature. Our search for published T eff s was extensive, though not exhaustive, and given time constraints we areadmittedly limited by what values were published in electronic tables that could be easily queried with e.g. Vizier .Many T eff s came from large catalogs by e.g. Philip & Egret (1980); Sokolov (1995); Cayrel de Strobel et al. (1997);Blackwell & Lynas-Gray (1998); Gray et al. (2001a, 2003); Taylor (2005); Valenti & Fischer (2005); Paunzen et al.(2006); Gerbaldi et al. (2007); Fitzpatrick & Massa (2007); Prugniel et al. (2007); Zorec et al. (2009); Soubiran et al.(2010); Casagrande et al. (2011), and the authors calculated photometric T eff s for OB dwarf standards using pho-tometry from Hauck & Mermilliod (1998), dereddening relations from Castelli (1991), and color-temperature relationsfrom Balona & Shobbrook (1984); Napiwotzki et al. (1993); Balona (1994). T eff s were also estimated for OB dwarfstandards using U – B vs. T eff data in Bessell et al. (1998).Here is an example of our evaluation of the median T eff for A0V stars. We find very consistent effective temperaturesamong A0V standards within a few hundred K of each other. The A0V standard Vega has had a very precise apparentT eff measured by Monnier et al. (2012) of 9660 K (in good agreement with many previous estimates), and we findthe literature median T eff for the other widely used MK standards γ UMa and HR 3314 to be 9361 K and 9760 K.While there are other A0V standards, two of these ( γ UMa, Vega) are considered “anchor” standards by Garrison(1994) (i.e. their classifications have remained the same over seven decades of use), and HR 3314 has retained its A0Vstandard status throughout (Morgan et al. 1953; Johnson & Morgan 1953; Garrison & Schild 1979; Gray & Garrison1987; Houk & Swift 1999; Gray et al. 2003). An exhaustive search for T eff s for A0V stars in the literature (265estimates) yields a median T eff of 9707 K. Based on these values, we adopt a median T eff of 9700 K for A0V stars. http://vizier.u-strasbg.fr/viz-bin/VizieR eff could be as high as 10000 K (Bessell 1979; Crowther 1997), nor as lowas 9394 K (Boyajian et al. 2012a), 9520 K (Schmidt-Kaler 1982), or 9530 K (Theodossiou & Danezis 1991). We notein particular that the recently published T eff scale by Boyajian et al. (2012a) appears to be most deviant among theA0V T eff values, and while that study relies on new interferometric observations, their survey contained only a singlenon-standard A0V star (HD 177724). Similarly sized deviations at the hundreds of K level were seen between our T eff scale and the Boyajian et al. (2012a) T eff scale. So while there are other modern color/T eff scales in the literature,we believe that ours is based on a very broad (but vetted) amount of photometric/T eff literature and classifications. Bolometric Corrections
The bolometric corrections (BCs) listed in Table 4 are derived for each spectral type by adopting the median BCamong several scales as a function of the adopted T eff , including Balona (1994); Bertone et al. (2004); Flower (1996);Bessell et al. (1998); Masana et al. (2006); Schmidt-Kaler (1982); Code et al. (1976); Casagrande et al. (2006, 2008,2010); Lanz & Hubeny (2007); Vacca et al. (1996); Lanz & Hubeny (2003) where they applicable. For the M dwarfs,the BC V scale was estimated via V-K s colors and the BC K results from Leggett et al. (2001), Dahn et al. (2002), andGolimowski et al. (2004), as well as the authors’ SEDF fits compiled in Table 7. Extensive notes and discussion can be found for each spectral type at o un g S t e ll a r C o l o r s a nd T e m p e r a t u r e s TABLE 1Spectral Types and Optical/Near-IR Photometry for Young,Nearby, Moving Group Members.
Star Grp 2MASS SpT Ref.
V B − V V − I C Ref. J a H a K aS W b (mag) (mag) (mag) (mag) (mag) (mag) (mag)HIP 490 TH 00055255-4145109 G0V 1 7.510 ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± γ
21 13.595 ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± P ec a u t e t a l. TABLE 1 — Continued
Star Grp 2MASS SpT Ref.
V B − V V − I C Ref. J a H a K aS W b (mag) (mag) (mag) (mag) (mag) (mag) (mag)HIP 30034 TH 06191291-5803156 K2V 4 9.130 ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± † RECX 16 EC 08440914-7833457 M5.75 22 12.505 ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± † ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± o un g S t e ll a r C o l o r s a nd T e m p e r a t u r e s TABLE 1 — Continued
Star Grp 2MASS SpT Ref.
V B − V V − I C Ref. J a H a K aS W b (mag) (mag) (mag) (mag) (mag) (mag) (mag)TWA 10 TWA 12350424-4136385 M2Ve 6 12.960 ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± P ec a u t e t a l. TABLE 1 — Continued
Star Grp 2MASS SpT Ref.
V B − V V − I C Ref. J a H a K aS W b (mag) (mag) (mag) (mag) (mag) (mag) (mag) Note . — Groups: (BP) – β Pic Moving Group; (EC) – η Cha Cluster; (TWA) – TW Hydra Association; (TH) – Tucana–Horologium Association;References for Spectral Type and optical
BV I C photometry: (1) Houk (1978); (2) Perryman & ESA (1997); (3) This Work (4) Houk & Cowley (1975); (5)Converted from Tycho-2 using Mamajek et al. (2002, 2006); (6) Torres et al. (2006); (7) Hawley et al. (1996); (8) Riaz et al. (2006); (9) Gray et al. (2006);(10) Jeffries (1995); (11) Henden et al. (2012); (12) Zuckerman & Song (2004); (13) Weis (1993); (14) Vyssotsky (1956); (15) Robertson & Hamilton (1987);(16) Houk & Swift (1999); (17) Stephenson & Sanwal (1969); (18) Gray (1989); (19) Houk & Smith-Moore (1988); (20) Schlieder et al. (2012a); (21)Rice et al. (2010b); (22) Luhman & Steeghs (2004); (23) Lyo et al. (2004); (24) Lawson et al. (2001); (25) Lawson et al. (2002); (26) Stephenson (1986);(27) Reid et al. (1995); (28) Koen et al. (2010); (29) Rodriguez et al. (2011); (30) Schneider et al. (2012b); (31) Scholz et al. (2005); (32) Teixeira et al.(2008); (33) Looper et al. (2010b); (34) White & Hillenbrand (2004); (35) Barrado Y Navascu´es (2006); (36) Shkolnik et al. (2011); (37) Kastner et al.(2008); (38) Looper et al. (2007); (39) Song et al. (2002);( a ) JHK S photometry from the 2MASS All-Sky Point Source Catalog (Cutri et al. 2003; Skrutskie et al. 2006);( b ) W W W W WISE
All Sky Data Release (Cutri & et al. 2012);( † ) Photometry excluded due to identified infrared excess (Hutchinson et al. 1990; Megeath et al. 2005; Riaz et al. 2006; Sicilia-Aguilar et al. 2009;Rebull et al. 2008; Gautier et al. 2008; Zuckerman et al. 2011; Schneider et al. 2012a,b). The binary pair HIP 10679 and HIP 10680 were studied inRebull et al. (2008), with only HIP 10679 identified as having a 24 µ m excess. However, they are separated by ∼ ′′ and thus it is likely that theHIP 10680 W W W TABLE 2Spectral standard stars used for classification.
Standard Spectral Telescope/Source ReferencesTypeHD 8512 K0IIIb SMARTS 1.5m/Stony Brook 1HD 3651 K0V SMARTS 1.5m/Stony Brook 2HD 10476 K1V SMARTS 1.5m/Stony Brook 1HD 153210 K2III SMARTS 1.5m/Stony Brook 1HD 22049 K2V SMARTS 1.5m/Stony Brook 1HD 16160 K3V SMARTS 1.5m/Stony Brook 1 α Sct K3III SMARTS 1.5m/Rochester 1TW PsA K4Ve SMARTS 1.5m/Rochester 1 β Cnc K4III Ba0.5 SMARTS 1.5m/Rochester 1HD 82668 K5III SMARTS 1.5m/Rochester 1HD 36003 K5V SMARTS 1.5m/Stony Brook 1GJ 529 K6Va SMARTS 1.5m/Rochester 1GJ 673 K7V SMARTS 1.5m/Rochester 2HIP 111288 K8V k SMARTS 1.5m/Rochester 3HD 142574 K8IIIb SMARTS 1.5m/Rochester 1HIP 3261 K9V SMARTS 1.5m/Stony Brook 3GJ 701 M0.0V SMARTS 1.5m/Rochester 4 υ Gem M0III SMARTS 1.5m/Rochester 1GJ 229A M1.0V SMARTS 1.5m/Rochester 4 ν Vir M1III SMARTS 1.5m/Rochester 1GJ 411 M2+V SMARTS 1.5m/Rochester 1GJ 752A M3-V SMARTS 1.5m/Rochester 1GJ 402 M4.0V SMARTS 1.5m/Rochester 4Gl 9066 M5-V SMARTS 1.5m/Stony Brook 1HD 151061 M5-M5.5IIIb SMARTS 1.5m/Rochester 1GJ 406 M6.0V SMARTS 1.5m/Rochester 4HD 118767 M6III SMARTS 1.5m/Rochester 1
Note . — References: (1) Keenan & McNeil (1989); (2) Gray et al. (2003); (3) Gray et al. (2006); (4) Henry et al. (2002)Spectral standards for F- and G-type stars were taken from Table 2 of Pecaut et al. (2012).
TABLE 3Observations and New Spectral Types
Object Spectral Type Spectral Spectral Type Ref.(This Work) Coverage (Literature)HR 9 F3V Blue F3Vn 1TYC 1186-706-1 K7V(e) Blue K5 2HIP 10679 G3V Blue G2V 3HIP 10680 F7V Blue F7V 3HIP 12545 K5IVe Blue/Red K6Ve 4HIP 12925 F7V Blue F8 5GJ 3305 M0Ve Blue/Red M1.1 6V1005 Ori K8IVe Blue/Red M0Ve 4HIP 23309 b K8Ve Blue/Red K8V kee 1HIP 23418 M4IVe Red M4 2HIP 25486 F7V Blue F8V(n)k 1HIP 29964 K3.5V Blue K3.5V ke 1RECX 1 K5IVe Red K6 7RECX 3 M3.5IV-Ve Red M3.25 7RECX 4 M0IVe Red M1.75 7RECX 7 K5IV(e) Red K6 7RECX 10 K9IV-Ve Red M1 7RECX 11 K5IVe Red K5.5 7TWA 21 K3IV(e) Red K3Ve 4TWA 7 M3IVe Red M2Ve 4TWA 1 K8IVe Red K6Ve 4TWA 2 M1.5IVe Red M2Ve 4TWA 3 M4IVe Red M4Ve 4TWA 12 M2IVe Red M2 8TWA 4 K6IV(e) Red K5V 4TWA 5A M2IVe Red M2Ve 4TWA 8A M3IVe Red M3 9TWA 9A K7IVe Red K7 9TWA 25 K9IV-Ve Red M1Ve 4TWA 20 M3IVe Red M2 10TWA 16 M2IVe Red M1.5e 11HD 139084 G8V Blue K0V k 1HD 155555AB a G5V, K1V Blue/Red G5IV+K0IV-V 12HD 161460 G9V Blue K0IV 4HIP 88399 F4.5V Blue/Red F6V 4V4046 Sgr K4IVe Blue K1e 13
TABLE 3 — Continued
Object Spectral Type Spectral Spectral Type Ref.(This Work) Coverage (Literature)GSC 7396-0759 M1IVe Red M1Ve 4HD 168210 G3IV Blue G5V 4CD-64 1208 K4V(e) Blue K5Ve 4TYC 9073-0762-1 M1Ve Red M1Ve 4TYC 7408-0054-1 K8IVe Red K8Ve 4TYC 6872-1011-1 K8IVe Red M0Ve 4CD-26 13904 K3.5IV(e) Red K4V(e) 4HIP 95270 F6V Blue F6V 4TYC 7443-1102-1 K9IVe Red M0.0Ve 14AT Mic A M4IVe Red M4Ve 4AT Mic B M4IVe Red M4Ve 4AU Mic M0Ve Red M1Ve Ba1 15AZ Cap K5IVe Red K6Ve 4TYC 2211-1309-1 c K8IVe Red M0.0Ve 14CPD-72 2713 K7IVe Red K7Ve 4BD-13 6424 M0V-IVe Red M0Ve 4
Note . — References: (1) Gray et al. (2006); (2) Stephenson (1986); (3) Harlan (1969); (4) Torres et al. (2006); (5) Cannon & Pickering (1918);(6) Shkolnik et al. (2009); (7) Luhman & Steeghs (2004); (8) Sterzik et al. (1999); (9) White & Hillenbrand (2004); (10) Reid (2003);(11) Zuckerman et al. (2001); (12) Strassmeier & Rice (2000); (13) Stephenson & Sanduleak (1977); (14) L´epine & Simon (2009); (15)Keenan & McNeil (1989); a Our blue spectrum of this star was classified as G5V while our red spectrum was classified as a K1V. Thereforewe adopted an overall spectral type of G5V,K1V. b Our spectrum of this star did not cover the Na doublet feature; c McCarthy & White (2012) were unable to detect Li in a high-resolution spectrum of this star, casting doubt on its membership in β Pic; however,we retain it as a member for the purposes of this study.
TABLE 4Intrinsic Colors of O9-M9 Dwarfs and Adopted T eff , BolometricCorrection Values SpT T eff BC V U – B B – V V – R C V – I C V – J V – H V – K S K – W K – W K – W K – W oung Stellar Colors and Temperatures 27 TABLE 4 — Continued
SpT T eff BC V U – B B – V V – R C V – I C V – J V – H V – K S K – W K – W K – W K – W P ec a u t e t a l. TABLE 5Synthetic Color Indices From BT-Settl and ATLAS9 models T eff log( g ) U – B B – V V – I C R C – I C J – H H – K S V – K S K – W K – W K – W K – W V – V T B T – V T V – H p Model(K) (dex) (mag) (mag) (mag) (mag) (mag) (mag) (mag) (mag) (mag) (mag) (mag) (mag) (mag) (mag)1400 3.5 3.475 3.191 5.213 2.454 1.916 1.344 13.151 1.379 1.967 2.652 2.921 -0.405 3.747 0.979 BT-Settl1400 4.0 4.909 0.795 7.250 2.854 1.149 0.635 13.788 0.983 1.083 2.226 2.292 0.091 1.668 2.550 BT-Settl1400 4.5 4.959 -0.165 7.918 3.184 0.919 0.306 13.929 0.923 1.148 2.326 2.426 0.561 1.123 2.903 BT-Settl1400 5.0 4.957 -0.945 8.421 3.433 0.893 0.098 14.184 0.873 1.363 2.474 2.631 1.085 0.836 3.189 BT-Settl1500 3.5 3.642 3.306 5.223 2.433 1.467 1.048 12.157 1.088 1.594 2.323 2.577 -0.412 3.929 1.006 BT-Settl
Note . — All synthetic colors are computed using solar metallicity models. The BT-Settl model colors presented here adopt the Asplund et al. (2009) solar composition, whereas the ATLAS9model colors presented here adopt the Grevesse & Sauval (1998) solar composition. Table 5 is published in its entirety in the electronic edition of ApJS. A portion is shown here for guidanceregarding its form and content. oung Stellar Colors and Temperatures 29
TABLE 6Intrinsic colors of 5-30 Myr old Stars and Adopted T eff , BolometricCorrection Values Spec. T eff B – V V – I C V – K S J – H H – K S K S – W K S – W K S – W K S – W V BC J Type (K) (mag) (mag) (mag) (mag) (mag) (mag) (mag) (mag) (mag) (mag) (mag)F0 7280 0.28 0.34 0.73 0.11 0.06 0.04 0.03 0.00 0.08 0.01 0.57F1 6990 0.34 0.39 0.89 0.14 0.07 0.04 0.04 0.00 0.09 0.00 0.68F2 6710 0.38 0.43 0.99 0.15 0.08 0.04 0.05 0.00 0.09 -0.01 0.75F3 6660 0.41 0.45 1.01 0.16 0.08 0.04 0.05 0.01 0.09 -0.01 0.76F4 6590 0.43 0.48 1.05 0.17 0.08 0.04 0.06 0.01 0.09 -0.01 0.79F5 6420 0.47 0.51 1.14 0.19 0.08 0.04 0.03 0.01 0.10 -0.02 0.85F6 6250 0.50 0.55 1.25 0.21 0.09 0.04 0.04 0.02 0.10 -0.04 0.91F7 6140 0.53 0.58 1.31 0.22 0.09 0.05 0.04 0.02 0.10 -0.05 0.95F8 6100 0.55 0.60 1.34 0.23 0.09 0.05 0.04 0.03 0.10 -0.06 0.96F9 6090 0.56 0.62 1.35 0.23 0.09 0.05 0.04 0.03 0.10 -0.06 0.97G0 6050 0.57 0.66 1.37 0.24 0.09 0.06 0.04 0.03 0.11 -0.06 0.98G1 5970 0.59 0.67 1.42 0.25 0.10 0.06 0.03 0.04 0.11 -0.07 1.00G2 5870 0.60 0.71 1.49 0.27 0.10 0.07 0.03 0.04 0.11 -0.09 1.03G3 5740 0.63 0.72 1.58 0.29 0.10 0.07 0.03 0.05 0.12 -0.11 1.08G4 5620 0.66 0.73 1.68 0.31 0.11 0.07 0.03 0.05 0.12 -0.14 1.12G5 5500 0.70 0.76 1.77 0.33 0.11 0.08 0.03 0.06 0.13 -0.17 1.16G6 5390 0.74 0.79 1.86 0.35 0.12 0.08 0.03 0.06 0.13 -0.20 1.19G7 5290 0.77 0.83 1.95 0.37 0.12 0.09 0.04 0.07 0.14 -0.23 1.23G8 5210 0.79 0.87 2.02 0.39 0.12 0.09 0.04 0.08 0.14 -0.26 1.25G9 5120 0.80 0.91 2.10 0.41 0.13 0.09 0.05 0.08 0.15 -0.29 1.27K0 5030 0.82 0.93 2.19 0.43 0.13 0.09 0.06 0.09 0.16 -0.33 1.30K1 4920 0.86 0.96 2.32 0.46 0.14 0.09 0.06 0.10 0.18 -0.38 1.34K2 4760 0.93 1.01 2.49 0.49 0.14 0.09 0.07 0.12 0.19 -0.46 1.40K3 4550 1.02 1.12 2.75 0.55 0.16 0.09 0.08 0.13 0.21 -0.60 1.44K4 4330 1.11 1.27 3.06 0.60 0.17 0.09 0.09 0.14 0.22 -0.77 1.52K5 4140 1.18 1.44 3.35 0.64 0.18 0.09 0.10 0.16 0.24 -0.95 1.58K6 4020 1.24 1.57 3.54 0.66 0.19 0.10 0.10 0.17 0.27 -1.08 1.61K7 3970 1.28 1.66 3.62 0.66 0.19 0.10 0.12 0.19 0.29 -1.14 1.63K8 3940 1.32 1.74 3.67 0.67 0.20 0.10 0.13 0.21 0.32 -1.17 1.63K9 3880 1.37 1.83 3.77 0.67 0.20 0.11 0.15 0.23 0.35 -1.24 1.66M0 3770 1.41 1.95 3.96 0.68 0.21 0.11 0.17 0.25 0.38 -1.38 1.69M1 3630 1.45 2.11 4.22 0.68 0.22 0.12 0.20 0.27 0.42 -1.58 1.74M2 3490 1.46 2.28 4.50 0.67 0.23 0.14 0.23 0.31 0.47 -1.80 1.80M3 3360 1.47 2.48 4.78 0.66 0.25 0.16 0.28 0.36 0.51 -2.03 1.84M4 3160 1.53 2.78 5.23 0.62 0.27 0.19 0.35 0.43 0.56 -2.43 1.91M5 2880 1.65 3.31 6.08 0.55 0.31 0.22 0.43 0.52 0.62 -3.21 2.01M6 ... ... ... 7.38 0.54 0.36 0.27 0.54 0.63 ... ... ...M7 ... ... ... 8.47 0.58 0.41 0.33 0.67 0.77 ... ... ...M8 ... ... ... 9.28 0.65 0.45 0.40 0.84 0.93 ... ... ...M9 ... ... ... 9.80 0.70 0.47 0.49 1.05 1.13 ... ... ...
TABLE 7T eff Comparison: SEDF versus Diameter-Derived T eff Object SpT Ref. T eff a T eff b m cbol (K) (K) (mag)GJ 15A M1.5V 1 3535 ±
14 3567 ±
11 6.560 ± ±
19 4950 ±
14 5.459 ± ±
22 5348 ±
26 4.981 ± ±
20 5398 ±
75 5.452 ± ±
18 4662 ±
17 5.406 ± ±
19 5077 ±
35 3.457 ± ±
21 5147 ±
14 4.169 ± ±
36 3801 ± ± ±
24 3907 ±
35 6.471 ± ±
25 3867 ±
37 6.479 ± ±
23 4085 ±
14 5.544 ± ±
14 3464 ±
15 5.873 ± ±
13 3497 ±
39 7.308 ± ±
17 3416 ±
53 8.778 ± ±
15 3618 ±
31 7.028 ± ±
12 3054 ±
79 7.280 ± ±
16 4507 ±
58 5.245 ± ±
13 3442 ±
54 8.560 ± ±
19 5337 ±
41 5.527 ± ±
13 3413 ±
28 7.231 ± ±
11 3222 ±
10 7.173 ± ±
12 3407 ±
15 7.007 ± ±
12 3104 ±
28 7.587 ± TABLE 7 — Continued
Object SpT Ref. T eff a T eff b m cbol (K) (K) (mag)GJ 764 K0V 4 5364 ±
21 5246 ±
26 4.498 ± ±
14 3692 ±
22 7.205 ± ±
14 4361 ±
17 4.599 ± ±
23 3932 ±
25 5.082 ± ±
16 4555 ±
24 4.217 ± ±
16 3713 ±
11 4.217 ± ±
19 3676 ±
35 5.871 ± ±
18 4699 ±
16 5.191 ± Note . — a : T eff from this work using the SEDF method. See Section 4.5.2 for details. b : T eff from Boyajian et al. (2012b) computed usingdirect angular diameter measurements. c : apparent bolometric magnitude estimated from our SED fit.Spectral Type References: (1) Henry et al. (2002); (2) Gray et al. (2006); (3) Gray et al. (2003); Keenan & McNeil (1989), (5) Torres et al. (2006); TABLE 8Objects Rejected From SED-T eff fitting Object Rejection ReasonHD 139084B Uncertain photometry resulting in poorly constrained T eff HD 164249B Uncertain photometry resulting in poorly constrained T eff RECX 11 K S band excessRECX 15 K S band excessRECX 16 IRAC µ m and 4.5 µ m excessTWA 27 IRAC µ m and 4.5 µ m excessTWA 29 Only three good bands of photometry ( JHK S )TWA 30A time-variable extinctionTWA 30B time-variable NIR excessTWA 31 Only three good bands of photometry ( JHK S )TWA 34 Only three good bands of photometry ( JHK S ) o un g S t e ll a r C o l o r s a nd T e m p e r a t u r e s TABLE 9T eff , Bolometric Magnitudes, Bolometric Corrections and AngularDiameter Estimates From SED Fitting Object 2MASS BT-Settl KuruczT eff θ BC V BC J m bol log( L/L ⊙ ) T eff θ BC V BC J m bol log( L/L (K) ( µ as) (mag) (mag) (mag) (dex) (K) ( µ as) (mag) (mag) (mag) (dex)HIP 490 00055255-4145109 5990 ±
16 250 ± ± ± ± ± ±
15 250 ± ± ± ± ±
25 346 ± ± ± ± ± ±
21 348 ± ± ± ± ±
15 180 ± ± ± ± ± ±
16 180 ± ± ± ± ±
18 242 ± ± ± ± ± ±
17 242 ± ± ± ± ±
15 189 ± ± ± ± ±
42 197 ± ± ± ± ±
18 188 ± ± ± ± ± ±
24 189 ± ± ± ± ±
14 159 ± ± ± ± ± ±
17 161 ± ± ± ± ±
10 252 ± ± ± ± ± ± ± ± ± ± ±
23 190 ± ± ± ± ± ±
32 190 ± ± ± ± ±
33 208 ± ± ± ± ±
29 214 ± ± ± ± ±
12 120 ± ± ± ± ±
11 121 ± ± ± ± ±
17 187 ± ± ± ± ± ±
15 187 ± ± ± ± ±
10 206 ± ± ± ± ± ±
10 208 ± ± ± ± ±
14 241 ± ± ± ± ±
16 242 ± ± ± ± ±
17 218 ± ± ± ± ± ±
16 218 ± ± ± ± ±
20 316 ± ± ± ± ± ±
21 317 ± ± ± ± ±
16 181 ± ± ± ± ± ±
15 182 ± ± ± ± ±
20 238 ± ± ± ± ± ±
19 239 ± ± ± ± ±
11 163 ± ± ± ± ±
10 166 ± ± ± ± ±
19 237 ± ± ± ± ± ±
19 237 ± ± ± ± ±
35 283 ± ± ± ± ± ±
36 284 ± ± ± ± ±
20 192 ± ± ± ± ± ±
24 192 ± ± ± ± ±
21 173 ± ± ± ± ±
12 195 ± ± ± ± ± ±
10 197 ± ± ± ± ±
25 181 ± ± ± ± ±
22 183 ± ± ± ± ±
20 209 ± ± ± ± ±
14 212 ± ± ± ± ± ±
16 215 ± ± ± ± ±
14 239 ± ± ± ± ±
14 244 ± ± ± ± ±
20 215 ± ± ± ± ±
11 144 ± ± ± ± ±
10 147 ± ± ± ± ±
16 207 ± ± ± ± ±
17 207 ± ± ± ± ±
23 250 ± ± ± ± ± ±
21 250 ± ± ± ± ±
15 252 ± ± ± ± ± ±
16 251 ± ± ± ± ±
14 203 ± ± ± ± ± ±
12 203 ± ± ± ± ±
16 179 ± ± ± ± ± ±
17 179 ± ± ± ± ±
24 471 ± ± ± ± ± ±
25 472 ± ± ± ± ±
29 331 ± ± ± ± ±
38 317 ± ± ± ± ±
18 178 ± ± ± ± ± ±
16 178 ± ± ± ± ±
19 247 ± ± ± ± ± ±
22 247 ± ± ± ± ±
28 117 ± ± ± ± ±
19 174 ± ± ± ± ± ±
19 174 ± ± ± ± ±
18 324 ± ± ± ± ± ±
23 327 ± ± ± ± ±
17 322 ± ± ± ± ± ±
22 326 ± ± ± ± ±
15 392 ± ± ± ± ± ±
17 243 ± ± ± ± ± ±
17 243 ± ± ± ± ±
18 427 ± ± ± ± ± ±
18 427 ± ± ± ± ±
14 160 ± ± ± ± ±
12 161 ± ± ± ± ±
16 89 ± ± ± ± ±
17 89 ± ± ± ± ±
17 234 ± ± ± ± ± ±
18 235 ± ± ± ± ±
20 46 ± ± ± ± ± ± ± ± ± ±
11 235 ± ± ± ± ±
16 208 ± ± ± ± ± ±
17 208 ± ± ± ± ±
13 189 ± ± ± ± ± ±
12 189 ± ± ± ± P ec a u t e t a l. TABLE 9 — Continued
Object 2MASS BT-Settl KuruczT eff θ BC V BC J m bol log( L/L ⊙ ) T eff θ BC V BC J m bol log( L/L (K) ( µ as) (mag) (mag) (mag) (dex) (K) ( µ as) (mag) (mag) (mag) (dex)HIP 32235 06434625-7158356 5661 ±
17 149 ± ± ± ± ± ±
16 149 ± ± ± ± ±
17 219 ± ± ± ± ± ±
19 220 ± ± ± ± ±
12 141 ± ± ± ± ± ±
13 142 ± ± ± ± ± ± ± ± ± ± ±
20 189 ± ± ± ± ± ±
17 190 ± ± ± ± ± ± ± ± ± ± ±
10 58 ± ± ± ± ± ±
23 94 ± ± ± ± ± ±
25 117 ± ± ± ± ± ±
32 116 ± ± ± ± ±
16 87 ± ± ± ± ± ±
23 94 ± ± ± ± ± ±
20 159 ± ± ± ± ± ±
20 160 ± ± ± ± ±
13 114 ± ± ± ± ± ±
21 105 ± ± ± ± ± ±
26 106 ± ± ± ± ±
22 144 ± ± ± ± ± ±
19 205 ± ± ± ± ±
24 206 ± ± ± ± ±
34 177 ± ± ± ± ± a ±
32 178 ± ± ± ± ± DK Leo 10141918+2104297 3861 ±
17 325 ± ± ± ± ±
23 328 ± ± ± ± ± ± ± ± ± ± c ... ... ... ... ...TWA 6 10182870-3150029 3982 ±
15 140 ± ± ± ± ±
19 142 ± ± ± ± ±
22 116 ± ± ± ± ±
27 117 ± ± ± ± ±
27 141 ± ± ± ± ±
29 142 ± ± ± ± ±
19 284 ± ± ± ± ±
22 204 ± ± ± ± ± ±
24 205 ± ± ± ± ±
160 59 ± ± ± ± ± b ... ... ... ... ...TWA 2 11091380-3001398 3556 ±
22 288 ± ± ± ± ± a ±
31 288 ± ± ± ± ± TWA 3 11102788-3731520 3112 ±
13 355 ± ± ± ± ±
18 159 ± ± ± ± ± a ... ... ... ... ...TWA 13A 11211723-3446454 3812 ±
21 179 ± ± ± ± ± a ±
26 181 ± ± ± ± ± TWA 13B 11211745-3446497 3549 ±
24 191 ± ± ± ± ± a ±
34 188 ± ± ± ± ± TWA 4 11220530-2446393 4223 ±
15 389 ± ± ± ± ± ±
15 395 ± ± ± ± ±
16 97 ± ± ± ± ±
22 299 ± ± ± ± ± a ... ... ... ... ...TWA 8B 11324116-2652090 3191 ±
63 122 ± ± ± ±
18 221 ± ± ± ± ± ± ± ± ± ±
17 66 ± ± ± ± a ... ... ... ... ...TWA 9B 11482373-3728485 3288 ±
20 109 ± ± ± ± ± ±
16 150 ± ± ± ± ± ±
15 153 ± ± ± ± ±
21 195 ± ± ± ± ± a ... ... ... ... ...TWA 25 12153072-3948426 3704 ±
20 203 ± ± ± ± ± a ±
28 205 ± ± ± ± ± ±
89 105 ± ± ± ±
23 141 ± ± ± ± ± a ... ... ... ... ...TWA 16 12345629-4538075 3475 ±
21 160 ± ± ± ± ± a ... ... ... ... ...TWA 10 12350424-4136385 3367 ±
18 157 ± ± ± ± ±
18 130 ± ± ± ± ± a ... ... ... ... ...HD 139084 15385757-5742273 4986 ±
15 307 ± ± ± ± ± ±
15 308 ± ± ± ± ±
24 116 ± ± ± ±
65 116 ± ± ± ±
13 537 ± ± ± ± ± ±
12 536 ± ± ± ± ±
14 142 ± ± ± ± ± ±
14 142 ± ± ± ± ±
13 156 ± ± ± ± ±
12 156 ± ± ± ± ±
21 166 ± ± ± ± ± ±
20 166 ± ± ± ± ±
22 203 ± ± ± ± ±
21 204 ± ± ± ± ±
11 256 ± ± ± ± ± ±
17 256 ± ± ± ± ±
45 201 ± ± ± ± ±
54 204 ± ± ± ± o un g S t e ll a r C o l o r s a nd T e m p e r a t u r e s TABLE 9 — Continued
Object 2MASS BT-Settl KuruczT eff θ BC V BC J m bol log( L/L ⊙ ) T eff θ BC V BC J m bol log( L/L (K) ( µ as) (mag) (mag) (mag) (dex) (K) ( µ as) (mag) (mag) (mag) (dex)GSC 7396-0759 18142207-3246100 3629 ±
23 121 ± ± ± ± ±
31 121 ± ± ± ± ±
26 170 ± ± ± ± ± ±
23 170 ± ± ± ± ±
21 317 ± ± ± ± ±
39 321 ± ± ± ± ±
23 166 ± ± ± ± ±
31 166 ± ± ± ± ±
15 179 ± ± ± ± ±
19 183 ± ± ± ± ±
15 247 ± ± ± ± ± ±
12 248 ± ± ± ± ±
23 144 ± ± ± ± ±
26 144 ± ± ± ± ± ± ± ± ± ± ± ± ± ± ±
20 262 ± ± ± ± ± ±
17 263 ± ± ± ± ±
48 164 ± ± ± ±
19 154 ± ± ± ± ±
24 155 ± ± ± ± ±
22 136 ± ± ± ± ±
29 137 ± ± ± ± ±
12 790 ± ± ± ± ± ±
22 757 ± ± ± ± ± ±
31 753 ±
10 -1.56 ± ± ± ±
19 290 ± ± ± ± ± ±
20 290 ± ± ± ± ±
13 215 ± ± ± ± ± ±
15 218 ± ± ± ± ±
17 182 ± ± ± ± ± ±
15 182 ± ± ± ± ±
13 228 ± ± ± ± ± ±
11 229 ± ± ± ± ±
18 154 ± ± ± ± ± ±
24 155 ± ± ± ± ±
18 250 ± ± ± ± ± ±
17 250 ± ± ± ± ±
19 405 ± ± ± ± ± ±
19 406 ± ± ± ± ±
14 207 ± ± ± ± ± ±
12 208 ± ± ± ± ±
22 160 ± ± ± ± ±
27 162 ± ± ± ± ±
14 235 ± ± ± ± ±
19 238 ± ± ± ± ±
14 308 ± ± ± ± ± ±
10 223 ± ± ± ± ± ±
20 288 ± ± ± ± ±
26 290 ± ± ± ± ±
18 176 ± ± ± ± ± ±
18 180 ± ± ± ± ±
21 196 ± ± ± ± ± ±
20 195 ± ± ± ± Note . — Teff values were fit at log( g ) = 4 . a Indicates log(
L/L ⊙ ) estimates use parallaxes from Weinberger et al. (2012). b Indicates log(
L/L ⊙ ) estimate uses parallax from Teixeira et al. (2008). c Indicates log(
L/L ⊙ ) estimate uses parallax from Teixeira et al. (2009).In addition, we have used a weighted mean distance of 94.3 ± η Cha cluster members. For HD 139084B, HD 164249B, AZ Cap, and HD 222259B we have adopted parallaxes from their brighter companions. P ec a u t e t a l. TABLE 10T eff , Bolometric Correction, and Bolometric Magnitude PolynomialCoefficients for 5-30 Myr Old Stars Y X