\documentstyle[newpaasms4,pptwocol,newpub,epsf]{preprint}

\begin{document}

\title{SPECTROSCOPIC ABUNDANCES IN COOL PLEIADES DWARFS\\
~\\
AND NGC 2264 STARS\\
} 

\author{Jeremy R.\ King\/\altaffilmark{1} and David R.\ Soderblom}
\affil{~\\
Space Telescope Science Institute\\
3700 San Martin Drive\\
Baltimore, MD 21218;\\
email: jking@stsci.edu, drs@stsci.edu} 

\author{Debra Fischer}
\affil{~\\
Department of Physics and Astronomy\\
San Francisco State University\\
San Francisco, CA 94132;\\
email: fischer@stars.sfsu.edu}

\author{Burton F.\ Jones}
\affil{~\\
Board of Studies in Astronomy and Astrophysics,\\
University of California Observatories/Lick Observatory,\\
University of California, Santa Cruz, CA 95064;
email: jones@ucolick.org} 
\tobe{20 April 2000}{The Astrophysical Journal}
\recacc{28 July 1999}{30 November 1999}
\url{1406}
\altaffiltext{1}{Current address:  Department of Physics, University of Nevada Las Vegas, 4505 S. Maryland Parkway, Las Vegas, NV 89154-4002}
\newpage

\abstract{We derive parameters and abundances of several elements in 2~cool Pleiades dwarfs, 4~cool NGC~2264 pre-main-sequence stars, and a probable NGC~2264 non-member from high resolution, moderate S/N Keck/HIRES spectra.  Our Pleiades Fe abundance agrees with previous spectroscopic and photometric 
values of hotter stars, and does not resolve the 0.3~mag distance modulus discrepancy between main-sequence fitting and {\it Hipparcos\ \/}parallaxes.  Abundances of Cr/Ca/Ti/Al are subsolar, mimicing the pattern of ISM abundances. While modest temperature errors may contribute, the results (particularly for Al) could suggest an association with ionization potential; such effects might be related to the Pleiades Li scatter.  The cluster Fe scatter and its relation to Li scatter is discussed. Three NGC~2264 members suggest \mbox{[Fe/H]~$=-0.15$} and near-solar ratios of other elements.  Mildly supersolar abundances for another object support its probable non-membership.  A fourth member exhibits a Mg/Si/Fe/Ni and Cr/Ti/Ca/Al dichotomy opposite to the Pleiads'; a relation to ionization potential is again suggested.  A 0.15--0.20~dex scatter or steep decline, neither well accommodated by stellar models, in the NGC~2264 Li abundances with $T_{\rm eff}$ is indicated.  We note the surprising presence of the ${\lambda}7774$ \ion{O}{1} triplet in our Pleaids, one of the cool NGC~2264 
stars, and the K6 field dwarf GL~241.  The inferred LTE O~abundances are enhanced by 0.23--0.85~dex over solar, suggesting that even NLTE calculations of the \ion{O}{1} triplet are incomplete and perhaps implicating the influence of an overlying chromosphere.  Our results demonstrate the utility of cluster abundances besides Fe and Li in addressing fundamental issuses concerning stellar evolution and systematic errors in the analysis of cool young stars. }
\base

\section{Introduction}

As assemblages of stars of varying mass but common history, open clusters are valuable probes of stellar physics and Galactic evolution.  Important relevant parameters are the ``metal'' abundances of the cluster stars.  These abundances are of specific interest for several reasons.  First, they are fundamental to exploration of Galactic structure and evolution via constraints provided by cluster age-metallicity relations and spatial metallicity gradients ({\it e.g.,\/} \markcite{taa97}Twarog, Ashman, \& Anthony-Twarog 1997). 

Second, open cluster star metallicities are important for constraining stellar physics, which has important implications for diverse issues such as the Universal age and the primordial Li abundance.  Recent work demonstrates the importance of open cluster abundances for stellar evolution in three different contexts.  First, ages of the oldest open clusters set a lower age to that of the Galactic disk; cluster metallicity is an important parameter in age determinations from theoretical isochrones (\markcite{dgg92}Demarque, Green \& Guenther 1992).  Second, as has been recently emphasized by \markcite{rapps}Randich \etal  (1997) and \markcite{cdp95}Chaboyer, Demarque \& Pinsonneault (1995; CDP95), metallicities are important parameters in constraining stellar models using open cluster main-sequence Li-$T_{\rm eff}$ morphologies. \markcite{sfira}Swenson \etal (1994) have called attention to the potential importance of abundances of elements besides Fe in interpreting cluster Li observations in terms of stellar models.  Third, \markcite{psskh}Pinsonneault \etal (1998) and \markcite{skhjfsp}Soderblom \etal (1998) have recently considered the 0.3~mag discrepancy between the Pleiades distance modulus derived from 
main-sequence fitting and from Hipparcos parallaxes. If this fundamental discrepancy is not due to systematic errors in the parallaxes, then a fundamental failure in our understanding of stellar evolution is suggested unless abundances in the Pleiades are vastly different than presently believed.  Unfortunately, little is known about the abundances of elements other than Li and Fe in even the most well-studied open clusters.  

Third, independent determinations of cluster metal abundances derived from data of a variety of stars and using different approaches provide a way to assess the confidence with which we can derive stellar abundances.  For example, abundances derived from cluster stars spanning a range in evolutionary state (e.g., giants, F-dwarfs, K-dwarfs) provide an evaluation of the relative adequacy of model atmospheres, $T_{\rm eff}$ calibrations, and other assumptions. Such tests would be particularly interesting for pre-main sequence (PMS) stars, but even less is known about detailed abundances of numerous elements in these stars than in well-studied open clusters.  Finally, the very way in which we measure metallicity---in proportion to the Sun---is not just because 
we can determine solar abundances on an absolute basis.  We are also naturally curious about how our Sun compares to the other stars around us. 

In previous work, we have utilized high resolution, moderate S/N spectra obtained with the Keck~I 10-m telescope and its HIRES spectrograph to determine Li abundances in solar-type stars of several open clusters.  These echelle data also contain considerable information on the abundances of other elements.  Here, we take a 
first step towards exploiting this information and addressing some of the above issues.  A natural starting point is the Pleiades, a cluster of fundamental importance and having previous abundance analyses.  The Pleiades abundances provide a good check on our analysis procedure and results, which we then apply to the more challenging PMS NGC~2264 stars.  Our results are some of the first abundances of elements other than Fe and Li in PMS field or cluster stars. 

We emphasize that the analysis has focussed on quality rather than quantity.  We have analyzed only those data which would yield the most reliable abundance determinations from which to draw conclusions.  For the Pleiades, we have confined ourselves to the two slowest rotating K-dwarfs from the dataset summarized below.  This is done to derive the most reliable abundances possible for an important comparison of our initial results with previous work.  Restricting ourselves to the slowest rotators does not just minimize uncertainties in the line strengths due to photon-noise.  More important is that even modest rotation makes many of the lines in our 
cool ($T_{\rm eff}{\sim}4500$) dwarfs unusable due to the effects of blending; we also wish to avoid any systematic errors caused by unrecognized effects of blending.  Even more restrictive criteria were mandated for the PMS NGC~2264 data.  These spectra have lower S/N than those for the Pleiades; we believed only the higher S/N spectra were suitable for the present purpose.  From the surviving slowly rotating objects having sufficient S/N, we eliminated those which were suspected of being SB2 (to avoid the effects of dilution and
blending) and those evincing strong H$\alpha$ emission (to avoid the effects of veiling).  This left four certain NGC~2264 members that we could analyze; additionally, one suspected non-member was purposefully
included to see if this could be confirmed from the abundances.  The objects are listed in the first column of Table~1.  
%\marginpar{Tab. 1}

\section{Observational Data}

Spectra of the cool Pleiades dwarfs HII~97 and HII~676 are those described in \markcite{j96}Jones \etal (1996; hereafter JSFS).  Spectra of the cool NGC~2264 stars J407, 428, 680, 682, and 1088 are 
those described in \markcite{S99}Soderblom \etal (1999; their Table~1 provides cross-identifications).  The data were acquired with the HIRES echelle spectrograph on the Keck~I 10~m telescope, and have a ${\sim}3$-pixel resolving power of ${\sim}45,000$.  The typical per-pixel S/N ratio is 70 for the Pleiads, and 45--75 for the NGC~2264 stars.  These values are modest, but wholly adequate for our analysis (perhaps marginally so for J680).  We use daytime sky spectra obtained with the same instrumentation during the 
same programs as solar proxy data.  These spectra are much higher quality; the Poisson-based per-pixel S/N is typically near ${\sim}700$.  Portions of the spectra can be seen in Figures~1 of JSFS and Soderblom \etal (1999). 

\section{Abundance Analysis}

\subsection{Features, Atomic Data, and Line Strengths}

Nearly all the atomic lines used in our analysis were selected from the extensive solar list of \markcite{T90} Th\'evenin (1990).  The vast majority are the allegedly clean ``case~a'' lines.  We first considered the effects of possible blends,  irrespective of whether a line was classified as a ``case~a'' feature, by inspecting the profiles 
from our high S/N solar proxy spectrum and the very high resolution solar flux atlas of \markcite{K84}Kurucz \etal (1984).  Because the extent of blending in our cool stars may be quite different, the Pleiades spectra were also inspected for obvious serious blending or problematic asymmetries.  Such features found in either the two Pleiads or the Sun were eliminated from consideration in the analysis.  Lines possibly 
significantly contaminated by telluric absorption or substantially in the wings of deep H$\alpha$ absorption were also excluded. 

The species and wavelengths of the features used here are listed in the first two columns of Table~2, and are taken from \markcite{T90} Th\'evenin (1990).  Lower excitation potential and oscillator strengths are from the same source, and are listed in the third and fourth columns.\footnote{Atomic data for the ${\lambda}8712,8717$ \ion{Mg}{1} lines are taken from \markcite{eag93}Edvardsson \etal (1993); that for the ${\lambda}6318,7691$ \ion{Mg}{1} lines is taken from \markcite{\ct90}Chang \& Tang (1990) and \markcite{Ket98}King \etal (1998).} These $g\/f$ values are solar oscillator strengths; i.e., they are derived by an inverse analysis of a specific solar spectrum via specific adopted solar abundances and a specific model atmosphere.  Because we can not adopt identical assumptions here (e.g., {\it stellar\/} model atmospheres that are internally consistent with the {\it solar\/} atmosphere used by Th\'evenin do not exist), a differential analysis with respect to the Sun is performed to yield abundances free from uncertainties in simply adopting $g\/f$ values and solar abundances.  Our derived abundances are thus normalized with a consistently determined solar abundance, and are $g\/f$-independent to the extent that we believe the solar oscillator strengths are plausible enough not to produce significant errors from differential curve-of-growth effects.
%\marginpar{Tab.~2}

\subsection{Equivalent Widths and Abundances} 

Equivalent widths of the selected lines in our stellar and solar spectra were measured using routines in the 1-d spectrum analysis package {\sf SPECTRE} (\markcite{fs}Fitzpatrick \& Sneden 1987).  To attempt determining yttrium abundances, the very weak solar $\lambda$6435 \ion{Y}{1} feature had to be 
measured from the Kurucz \etal (1984) solar flux atlas.  The  line strengths are reported in Table~2.  Not all lines could be measured in all stars due to differences in instrumental configuration, radial velocity shifts, cosmic ray hits, and noisy line profiles.  More uncertain measurements are marked with a colon.  

Abundances were derived from the equivalent widths using an updated version of the LTE analysis package {\sf MOOG} (\markcite{s73}Sneden 1973) and model atmospheres from the grids of \markcite{k92}Kurucz (1992).  The van der Waals broadening was handled in the manner of \markcite{u55}{\"U}nsold (1955) with 
an enhancement factor of 2.2, typical of the value found from and employed in solar and stellar analyses,  applied to all species.  The derived absolute solar abundance and normalized stellar abundance are given in Table~1 for each line. 

The ${\lambda}8712$ and ${\lambda}8717$ lines are the best Mg features in our spectra to derive Mg abundances.  The former is particularly well-blended in the Pleiades spectra with a neighboring \ion{Fe}{1} line.  Thus, Mg abundances were derived via spectral synthesis. The 8700~{\AA} line list was taken from the Kurucz CD-ROMs.  The Pleiades and solar spectra show synthetic absorption in the region to be lacking, since many observed features are not reproduced.  However, we do not believe this introduces systematic error into our Mg abundances because reliable continuum windows seem to exist, and because the absorption in the immediate Mg regions is dominated by absorption from a few stronger lines which are adjusted based on the
solar spectra.  The synthetic calculations were carried out in {\sf MOOG} with instrumental broadening modeled by convolution with a Gaussian; rotational broadening was also included using the published $v\sin{i}$ values. 

\subsection{Stellar Parameters and Uncertainties} 

\subsection*{3.3.1. Effective Temperatures and Microturbulent Velocity} 

Initial estimates of the stellar $T_{\rm eff}$ values were taken from the photometric estimates in JSFS and Soderblom \etal (1999).  Final values were determined spectroscopically by enforcing excitation
balance (i.e., no trend of [m/H] versus excitation potential) of the \ion{Fe}{1} lines---the most numerous species in our analysis.  Similarly, the microturbulent velocity was determined by the simultaneous 
requirement of no trend in the abundances with line strength; a solar value of 1.1~km~s$^{-1}$ was assumed.  

The \ion{Ni}{1} and \ion{Ti}{1} lines in our spectra provided useful checks on the parameters derived from the \ion{Fe}{1} lines.  The adequacy of these procedures is indicated in Figure~1, which shows the abundance 
vs.\ ${\chi}$ and reduced line strength for \ion{Fe}{1}, \ion{Ni}{1}, and \ion{Ti}{1} lines in HII~676 after adjustment of the individual [Ni/H] and [Ti/H] values to the mean [Fe/H] by addition of a constant. The derived parameters are listed in Table~1.  
%\marginpar{Fig.~1} 
\begin{figure}%1
\centerline{\epsfxsize=\linewidth\epsfbox{fig1.eps}}
\caption[]{The top panel shows the line-by-line normalized abundances of Fe (stars), Ni (open circles), and Ti (filled squares) as a function of lower excitation potential for the Pleiad HII~676. The individual Ni and Ti abundance values have been adjusted to the mean [Fe/H] by addition of a constant.  The bottom panel shows the line-by-line abundances {\it versus\/} reduced line strength.} 
\end{figure} 

By altering the assumed values of $T_{\rm eff}$ and (${\xi}$), trends can be introduced into the planes of Figure~1.  Confidence levels of the resulting correlations are calculated from the correlation coefficients, 
yielding uncertainties in the parameters.  We estimate the uncertainties in the microturbulence of ${\la}0.25$~km~s$^{-1}$.  Uncertainties in the spectroscopic $T_{\rm eff}$ values are ${\sim}75$~K for the Pleiads and ${\la}90$~K for the NGC~2264 stars; we conservatively adopt 90~K as the uncertainty for all of the latter. 

The derived $T_{\rm eff}$ values for the two Pleiads are in fine agreement (${\pm}25$~K) with the initial photometric estimates, suggesting no gross errors in the excitation analysis of these nearby ZAMS stars.  Similar agreement is seen for the probable NGC~2264 non-member J1088, which we suspect is an early-G interloping 
field dwarf.  In contrast, we find $T_{\rm eff}$ values for J407, 428, 680, and 682 that are 201, 84, 228, and 522~K larger than the photometric estimates.  On average, these differences are considerably larger than the 
errors in our spectroscopic determinations.  Likely causes are underestimates in reddening (systematic or for particular stars), deviations in color-$T_{\rm eff}$ relations for some PMS stars, inadequacy in the model 
atmospheres, and excitation departures.  Regardless, the differences underscore the uncertainty in PMS $T_{\rm eff}$ determinations (\markcite{K98}King 1998), which are critical to evaluating present day Li abundances from the youngest stars.  Given the $T_{\rm eff}$ revisions, we later present revised Li abundances in Table~5. 

\subsection*{3.3.2. Gravities and Model Metallicities}

The assumed overall metallicity of the model atmospheres is given in Table~1.  These are very close to the final determined Fe abundances, but the effects on the derived abundances of modest departures from the assumed model metallicity are not great.  We simply adopt a reasonable 0.15~dex uncertainty in the model atmosphere metallicity uncertainty when considering detailed errors. 

A gravity of log $g=4.60$ was assumed for both Pleiads on the basis of the binary star calibration shown in Figure~16.7 of \markcite{g92}Gray (1992) and theoretical isochrones.  We believe this gravity accurate 
certainly within 0.15~dex (taken as the uncertainty), and probably less.  For the warm NGC~2264 non-member J1088, a gravity of log $g=4.50$ was derived by forcing agreement between abundances derived from 
gravity-sensitive \ion{Fe}{2} lines and those derived from the \ion{Fe}{1} features; given the internal Fe abundance scatter and uncertainty in the other parameters, this estimate is accurate to within ${\pm}0.25$~dex.  Ideally, gravities for the other NGC~2264 stars would be similarly determined via ionization balance.  Unfortunately, most of the \ion{Fe}{2} lines in our wavelength region are high excitation lines immeasurably weak in our cool star spectra.  The 2 or 3 measurable lower excitation \ion{Fe}{2} lines we could identify were noticeably blended in the Pleiades spectra and deemed unsuitable. 

Thus, we have simply utilized the mass and radii values from the PMS theoretical models discussed in Soderblom \etal (1998) to estimate gravities for our NGC~2264 stars; these are given in Table~1. 
We have assigned a generous internal uncertainty of 0.50~dex to the theoretical gravities; though we believe the results to be more robust than this, this error also encompasses possible systematic errors.  The log~$g$ values are based on the photometric $T_{\rm eff}$ estimates adopted by Soderblom \etal (1998), and the detailed abundances given in Table~2 are calculated assuming no log~$g$ adjustment based on our refined spectroscopic $T_{\rm eff}$ estimates.  The assumed gravity does not affect the derived parameters,\footnote{Altering the assumed log~$g$ by a substantial 0.5~dex typically alters the derived 
microturbulence by ${\sim}0.1$~km/s and the inferred $T_{\rm eff}$ by 30--40~K.} but does impact the abundances directly.  The sensitivities to all parameters are given in Table~3 containing our final results.

To consider systematic effects, revised log~$g$ values were estimated from the PMS models assuming our hotter NGC~2264 $T_{\rm eff}$ values. The results are larger log~$g$ by 0.17, 0.09, 0.26, and 0.56~dex for
J407, J428, J680, and J682.  Even revisions of this magnitude do not seriously impact our results.  For example, the most extreme revision is ${\Delta}$[Si/H]$=+0.15$ dex for J682; changes with respect to Fe are smaller (even in this most extreme case ${\Delta}$[Si/Fe]${\la}+0.09$).  Regardless, in Table~4 we list final mean abundances assuming both the adopted and revised gravities so that the systematic effects can be assessed.  

\subsection{Blending and Line Identification Concerns} 
 
Differential analyses of stars with parameters differing from the solar temperature and pressure may be complicated by variations in the strength of a blending feature contaminating a given line.  E.g., a a 
low-excitation line of moderate strength in a cool star may be considerably weaker in the hotter Sun.  If such a feature was contaminated by an unrecognized high excitation feature, the relative contribution of the high excitation line to the blended line strength would be higher in the Sun; thus one would erringly derive an underabundance in the cool star. 

We have indeed observed such effects during our analysis.  The low-excitation \ion{V}{1} lines listed in Table~2 (and another \ion{V}{1} feature, $\lambda$6504.2, not listed) yield low abundances with respect to Fe in our cool stars, and a high abundance in the warm star J1088.  A line list from the Kurucz CD-Roms 
implicates high excitation solar contaminants as the cause---while the Kurucz $g\/f$ values suggest that neighboring high excitation features of \ion{Sc}{2}, \ion{Ti}{1}, \ion{Cr}{2}, and \ion{V}{1} have solar line strengths several orders of magnitude lower than the low-excitation $\lambda$6452.31 \ion{V}{1} feature, 
a ${\chi}=5.62$~eV \ion{Si}{1} line at $\lambda$6452.30 is of sufficient strength to affect the derived solar $g\/f$ values of Th\'evenin (1990) such that a ${\sim}0.4$ dex underestimate in [V/H] of our cool stars is the result. For the $\lambda$6605.92 \ion{V}{1} line, we find a neighboring ${\chi}=4.01$~eV \ion{Ti}{2} feature at 6605.90~{\AA} whose Kurucz $g\/f$ value (and a canonical solar Ti abundance) would result in a modified \ion{V}{1} solar $g\/f$ value that yields---more in line with expectations---[V/H]${\sim}0.0$ in, e.g., the cool Pleiads. 

Another example is provided by the low-excitation (${\chi}=0.50$~eV according to Th\'evenin, 0.00~eV according to the Kurucz list) $\lambda$6687.51 \ion{Y}{1} feature.  Whereas our $\lambda$6435 \ion{Y}{1} line gives a seemingly reasonable Pleiades abundance of [Y/H]$=+0.05$ for HII 676, the $\lambda$6687.51 line gives $-1.01$ in both HII 676 and 97.  Whereas the Th\'evenin and Kurucz $g\/f$ values for the former line agree to within a percent, the $g\/f$ values for the latter feature differ by one and a half orders of magnitude.  This difference is again symptomatic of an errant solar $g\/f$ value due to a blend not affecting the 
Pleiades stars.  Either of the neighboring ${\chi}=4.14$ and 4.88~eV \ion{Fe}{1} and \ion{Ca}{1} features at 6687.49 and 6687.72~{\AA} in the Kurucz line list provide the needed contamination to explain the discrepancy. 

We estimate that such blends or incomplete solar line identifications may seriously afflict ${\sim}25$\% of weak ($W{\la}15$~m\AA, say) low excitation solar lines.  Our results should not be greatly affected since 
most of the known cases we identified are limited to Th\'evenin's ``case~b'' features, which are not relied upon heavily here; an exception is the apparently clean ``case~a'' $\lambda$6793.6 \ion{Y}{1} line which we found 
to yield [Y/H]${\la}-1$ in our cool NGC~2264 stars.  Happily, miscreants can be efficiently identified and excluded by {\ }a)~their consistent effect on the abundances in different stars  {\ }b)~comparison of the solar $g\/f$ values with Kurucz $g\/f$ values, and {\ }c)~identifying specific blending features of sufficient wavelength and strength using the Kurucz line list. 

A similar but opposite effect may exist for high excitation lines in that low-excitation blends not affecting the solar $g\/f$ values or abundances may contaminate the cool Pleiads, resulting in derived overabundances.  This 
problem is more insidious since it is not necessarily confined to the known-to-be-blended ``case~b'' features.  For example, we find the (${\chi}=5.86$) $\lambda$6848.57 \ion{Si}{1} and (${\chi}=5.61$) $\lambda$7552.50 \ion{Ni}{1} features to be high excitation ``case~a'' lines which yield huge overabundances ([X/H]~$=+0.5-0.6$) in both Pleiades stars---apparently due to contamination by low-excitation blends not apparent in the solar spectrum.  While our estimate is not necessarily robust nor unbiased, our best guess is that perhaps as many as 15--20\% of the high excitation lines could be noticeably affected by low-excitation blends in our cool Pleiades dwarfs.  Fortunately, most of this minority fraction can be identified and excluded as described above.  

\section{Results}

Mean abundance ratios, [{\it X\/}/H] and [{\it X\/}/Fe], are listed for our Pleiades stars and NGC~2264 objects in Tables~3 and~4.  V~abundances are not included because they are deemed unreliable; the Y results should be regarded with some caution since they are derived from a single line and suffer from lingering blending concerns.  The tables list the mean {\it internal\/} uncertainty (${\sigma}_{\rm int}$) derived from the 
line-to-line scatter (assumed similar to the Fe or Ni results in those cases with few lines), followed by total mean uncertainties (${\sigma}_{\rm tot}$) which take into account the internal errors and the uncertainties in the parameters listed in Table~1.  The final four columns of Table~3 list representative sensitivities to the parameters for our cool stars; these are not applicable to the considerably hotter J1088. The NGC~2264 abundance ratios that result from upward gravity revisions, discussed above, according to our spectroscopic $T_{\rm eff}$ values are given in parentheses in Table~4. 
%\marginpar{Tab.~3}
%\marginpar{Tab.~4}

\section{Discussion}

\subsection{The Pleiades Results}

\subsection*{5.1.1. Pleiades Fe Abundance and Comparison with Previous Work}

The mean Fe abundance from our two cool Pleiads is [Fe/H]~$=+0.06$, with a total uncertainty of ${\sim}0.05$~dex.  Potential systematic errors (continuum normalization, blending, model atmosphere adequacy) are more difficult to quantify, but we do not believe them to be large.  The Fe analysis of two 
cool Hyades dwarfs by \markcite{kh96}King \& Hiltgen (1996) is very similar to the present one.  Their value of [Fe/H]~$=+0.11$ is in excellent agreement with the F- and G-star determinations of \markcite{c85}Cayrel \etal (1985) and \markcite{bf90}Boesgaard \& Friel (1990), who find mean values of [Fe/H]~$=+0.12$ 
and $+0.15$; this agreement is encouraging considering the 1500--2000~K range in $T_{\rm eff}$ covered by the studies, and the three different methods (photometric calibrations, H$\alpha$ profile fitting, \ion{Fe}{1} fine analysis) employed to determine $T_{\rm eff}$. 

Our Pleiades [Fe/H] falls between previous estimates derived from warmer stars.  \markcite{bbr88}Boesgaard \etal (1988) studied 17 Pleiades F-stars and determined a spectroscopic value of [Fe/H]~$=-0.03$.  \markcite{b89}Boesgaard (1989) selected the best eight stars and utilized new measurements of six Fe lines to find [Fe/H]~$=+0.02$.  \markcite{bf90}Boesgaard \& Friel (1990) analyzed new spectra for 12 of the same F-stars and estimated [Fe/H]~$=-0.02$.  \markcite{ccc88}Cayrel \etal (1988) spectroscopically determined Fe abundances in four Pleaides G dwarfs; their mean abundance is [Fe/H]~$=+0.13$. 

Our Pleiades K dwarf-based value of [Fe/H]~$=+0.06$ is intermediate to these F- and G-star estimates, and can not be said to conflict with either given the mutual uncertainties.  Placing 28 Pleiades F- and G-stars in the $m_1$ vs.\ $\beta$ plane, \markcite{e86}Eggen (1986) concluded that the Pleiades Fe abundance lies between [Fe/H] of 0.0 and 0.1.  His Stromgren calibration-based estimate of [Fe/H]~$=+0.08$ is in fine agreement with our spectroscopic value.  \markcite{N81}Nissen (1981) used Stromgren photometry and narrow band indices to find a mean [Fe/H]~$=-0.08 {\pm}0.04$ for 6 Pleiades F-stars, though two stars seem to have considerably lower abundances relative to the other four.  The difference between the Eggen and Nissen results is mildly surprising since the former uses the photometric metallicity calibration of the latter.  

\subsection*{5.1.2. Abundances of Other Elements}

Derived abundance ratios (with respect to Fe) of other elements in our two Pleiads and the NGC~2264 objects are plotted versus atomic number in Figure~2.  Fe, Mg, and Si are dominant contributors to free electrons in 
cool star photospheres.  The Mg and Si abundances of our Pleiads appear to scale with the mildly enhanced Fe.  From Table 3 we find mean values of [Mg/H]~$=+0.05$ and [Si/H]~$=+0.07$, again with a total uncertainty of ${\sim}0.05$ dex.  Ni appears to be underabundant with respect to iron; our results suggest [Ni/Fe]~$=-0.08{\pm}0.03$.  The similarity in the Ni and Fe line strengths, excitation distribution, ionization potential, and parameter sensitivities result in the small uncertainty in this ratio.  Because of this, we tentatively conclude the slight underabundance is real; indeed, a similar Ni underabundance is found in all 5~NGC~2264 stars too.  If true, the Pleiades ratio would lie on the extreme edge of the disk field star distribution of Edvardsson \etal (1993).  Excluding unknown systematic effects as the cause will require additional high quality data; 
particularly valuable would be Pleaids near the solar $T_{\rm eff}$.  
%\marginpar{Fig.~2}
\begin{figure*}%2
\centerline{\epsfxsize=4.5in\epsfbox{fig2.eps}}
\caption[]{Abundance ratios (with respect to Fe) are plotted versus atomic number for the two Pleiads and probable NGC~2264 non-member (left-hand column), and the NGC~2264 PMS members (righ-hand column).}  
\end{figure*}

Interestingly, the other light metals consistently show larger underabundances relative to Fe in both Pleiads.  In contrast to Ni, this is not seen in our NGC~2264 stars.  We find mean Pleiades values of 
[Cr/Fe]~$=-0.13{\pm}0.06$, [Ti/Fe]~$=-0.14{\pm}0.07$, 
[Ca/Fe]~$=-0.18{\pm}0.07$, and [Al/Fe]~$=-0.22{\pm}0.06$. Comparison with the numerous data from Edvardsson \etal (1993) indicates that these Pleiad's abundance ratios (particularly for Al and Ca) are remarkably low compared with field disk stars---especially when considered together.  Again, identification of unknown systematic effects could come from detailed analyses of high quality spectra of additional Pleaids. 

There are several possible explanations for these odd ratios.  First is Galactic chemical evolution.  \markcite{CL}Cunha \& Lambert (1994) find evidence of correlated B-star O and Si abundance variations between Orion subgroups and suggest that these arise from variations in enrichment of the nascent gas by Type~II supernovae.  This is not a likely explanation for our results given the differences between the 
$\alpha$-element (Type~II SN products) ratios---Mg and Si are more abundant than Ca and Ti.  No simple
signature of Galactic chemical evolution is reminiscent of our Pleiades abundance ratios.  
 
Another possibility is circumstellar/interstellar chemistry.  The elements Al, Ca, Ti, and perhaps Cr are often found to be considerably more depleted in the ISM than Fe, Mg, and Si.  This pattern qualitatively mimics our photospheric abundances.  Such a pattern might then be related to a disk/planet-star environment and/or photospheric accretion. Clearly, it would be of considerable interest to confirm the present results with additional data in more Pleiads.  

A third possibility is ionization-related effects.  The elements Mg, Si, Fe, and Ni have first ionization potentials ranging from 7.64 to 8.15~eV.  Those for Al, Ca, Ti, and Cr are lower---ranging from 5.99 to 6.76~eV.  Systematic errors in our analysis procedure that might lead to such an ionization-dependent effect are not apparent; moreover, this pattern is not seen in our NGC~2264 stars.  We can only speculate that such a pattern might be introduced by, e.g., model atmosphere uncertainties, the presence of convective inhomogeneities, chromospheric effects, or other NLTE effects.  Such effects (whatever their cause) might have important consequences for the absolute and star-to-star abundances of Li, which has an even lower ionization potential of 5.39~eV. 

To distinguish between the depletion and ionization possibilities, we present abundances of K in Table~5.  K is useful in this regard since its interstellar depletion is comparable to or lower than that of Fe, Mg, and Si, but its ionization potential (4.34~eV) is even lower than that of Al, Ca, Ti, and Cr.  These abundances are derived from the strong \ion{K}{1} resonance line at 7699~\AA; uncertainty in these abundances may arise from the lack of hyperfine structure information and NLTE effects.  Excluding these, the total uncertainties in the individual K abundances are estimated to be ${\sim}0.15$ dex.  The mean [K/H] ratio, $-0.33{\pm}0.10$, is subsolar.  Our suspicion is that lack of hyperfine structure and NLTE effects comparable to those of the analogous $\lambda$6707 \ion{Li}{1} feature may lead to an {\it over\/}estimate of this value (perhaps by a few tenths of a dex). 
%\marginpar{Tab.~5}

The K abundances suggest an abundance pattern associated with ionization potential rather than interstellar depletion.  If so, the similarities in the \ion{K}{1} and \ion{Li}{1} transitions, atomic data, and line strengths
may mean that Li abundances could be afflicted too.  More interesting is the possibility of star-to-star abundance variations.  If the source of an ionization-dependent effect varies from star-to-star in the Pleiades, might this introduce related scatter in their Li (and K) abundances?  This interesting question merits study of additional Pleiades stars and illustrates the potential utility of detailed abundances of numerous elements (not just Fe).

A fourth possibility is a temperature error ($T_{\rm eff}$ 100~K too low, say) modestly larger than the estimated uncertainty.  While such an error does not affect the mean [Fe/H] abundance, it does result in a suggestive concordance of the [Ni,Cr,Ti,Ca,Si,Mg/Fe] ratios at a value of $-0.07$ to within 0.03--0.04~dex.  However, the [Al/Fe] and [K/Fe] ratios remain significantly lower at $-0.18$ and $-0.30$ respectively.  In this case, plotting the revised abundance ratios in the various planes of Edvardsson \etal (1993) indicates the Pleiads remain consistently distinct from nearby field dwarfs.  

Moderate $T_{\rm eff}$ adjustments are unable to remove the relative Al and K deficiencies and the differences with field dwarf abundance ratios that we find.  Larger $T_{\rm eff}$ adjustments also fail to achieve simultaneous concordance.  E.g., an upward adjustment of a few hundred degrees would alleviate the subsolar [Al/Fe] and [K/Fe] ratios, but also introduce serious anomalies---[Ca,Ti/Fe] would approach $+0.2$ while the other $\alpha$-element ratios [Mg,Si/Fe] would approach $-0.25$ due to their large and opposite 
temperature sensitivities.  We note that $T_{\rm eff}$ adjustments are even more problematic in explaining an analogous abundance dichotomy in one of our NGC~2264 stars below.  This does not mean that adjustments larger than our uncertainties are not at all possible---that is difficult to argue against definitively; rather, apparent curiosities remain after such adjustments. 
 
\subsection*{5.1.3. Comment on Intra-Cluster Abundance Differences}  

It is impossible to place meaningful limits on the internal Pleiades abundance scatter from two stars, but a couple comments are made.  First, the agreement between our HII~676 and 97 abundances is outstanding.  All the ratios agree to well within the uncertainties alone.  The agreement seems considerably better than expected for $1\sigma$ uncertainties, but it is not obvious that these have been overestimated. 

Second, \markcite{cds}Cayrel de Strobel (1990) suggested that the Pleaides F- and G-star data indicate real star-to-star Fe variations.  We do not believe this to be the case though.  The Pleiades F-star Fe abundances of Boesgaard \& Friel (1990) evince a scatter of 0.06--0.07~dex, which is actually smaller than their typical per star error of 0.09--0.10~dex.  Similarly, the scatter in the G-star Fe abundances of Cayrel \etal (1988) is 0.10~dex; again this is actually smaller than the typical ${\sim}0.15$ dex per star uncertainty listed in their Table~1. 

Differing [Fe/H] values would alter a star's position in the cluster HR diagram.  Consider the case of HII~296 and 1776.  Cayrel \etal (1988) determined [Fe/H]~$=+0.26$ and $+0.02$ for these two Pleiads.  The semi-empirical and theoretical ZAMS fiducial and luminosity-metallicity sensitivity from \markcite{vp89}VandenBerg \& Poll (1989) and the two stars' $(B-V)$ indices predict a 0.62~V magnitude difference between the two stars if they have the same metallicity.  If their true metallicities do differ by 0.24~dex, the predicted V~magnitude difference is 0.25~mag.  The observed magnitude difference is 0.55 magnitudes, suggesting only a 0.05~dex [Fe/H] difference at most.\footnote{Unless the helium abundance of HII~296 was a very large $Y{\sim}0.37$ compared to the canonical solar value of $Y{\sim}0.27$.} Evidence for an internal ``metallicity'' scatter in the Pleiades remains wanting.  

\subsection*{5.1.4. Abundances and the Pleiades Distance Modulus}

Pinsonneault \etal (1998) have considered the 0.3~mag discrepancy between the Hipparcos 
parallax-based Pleiades distance modulus and that inferred from cluster main sequence fitting.  They suggest the source of the discrepancy to be systematic errors at the 1~mas level in the parallaxes.  As they note, another possibility is that the Pleiades metallicity is really a factor of ${\sim}2$ {\it lower\/} than previously believed.  Our K-star results suggest that this is not the case.  Our Fe, Mg, and Si---the dominant electron donors in the stellar photosphere---abundances are slightly larger (by 0.06~dex) than solar. 

\subsection*{5.1.5. Li Abundance Scatter in Cool Pleiades Dwarfs}

Our two Pleiads' Li abundances presented in Table~5 (where we have remeasured the JSFS line strengths as part of this work) suggests a difference of ${\sim}0.65$~dex.  Their $T_{\rm eff}$ differs by 100--120~K, as inferred from both the photometry and excitation balance.  From Figure~4 of JSFS we estimate the slope of the Li abundance decline with decreasing $T_{\rm eff}$, and find that the $T_{\rm eff}$ difference is consistent with only 0.2~dex of the Li abundance difference; thus,  a real abundance difference of ${\sim}0.4$~dex remains. 

The source of such scatter in cool open cluster dwarfs is a long-standing problem of great interest and importance.  Numerous mechanisms have been proposed to explain this scatter.  An additional one---perhaps unpalatable but also unexplored---is the presence of intra-cluster abundance variations. As emphasized by Swenson \etal (1994), Li burning in cool dwarf models is a sensitive function of ``metallicity'', which (in this context) should not be identified with Fe abundance alone.  We find no abundance differences between HII~97 and 676 that are significant with respect to the errors.  However, it is also impossible to exclude small metal abundance differences of, e.g., 0.03~dex.  As can be gauged from Figure~3 of CDP95, such a small difference can lead to large ${\sim}0.4$~dex differences in standard model Li depletion at the age and temperature of our Pleiads. 

Thus, very slight metal abundance differences could explain some or all of the Li abundance difference between HII~97 and 676.  Exploring this possibility further would benefit from analysis of the slowly rotating Pleiads more amenable to accurate abundance analysis near 5100~K evincing a Li abundance spread ${\ga}1$~dex.  At this $T_{\rm eff}$, the model predictions suggest that the observed Li abundance spread would require 
metallicity differences of ${\sim}0.25$~dex.  Variations of this size can be easily detected or excluded.  The [Fe/H] scatter of ${\la}0.1$~dex inferred from combining our results with those of Boesgaard \& Friel (1990) 
and Cayrel \etal (1988) argues, albeit indirectly, against a metal abundance scatter of this size as responsible for the Li spread.  

As noted above, our abundances seemingly evince an ionization potential-related pattern.  If real, this could have consequences for the Li abundances.  Moreover, if such an effect varies from star-to-star, this may introduce scatter in the Li abundances.  Relatedly, we note that Pleiades Li abundances and K~line strengths show scatter significantly larger than measurement errors, and that the star-to-star Li and K differences are correlated (\markcite{KKP}King \etal 2000).  This possibility could be explored with, e.g., additional Al determinations in Pleaids with a range of $T_{\rm eff}$.  If the Al underabundance is a real signature of Galactic evolution, one expects a constant low abundance.  If more speculative possibilities (circumstellar/interstellar chemistry or ionization) are correct, then Al might show star-to-star variations; of particular interest would be the extent of scatter above and below ${\sim}5200$~K, the $T_{\rm eff}$ below which large variations in Pleiades Li become apparent. 

\subsection{The NGC 2264 Results} 

\subsection*{5.2.1. J407, J428, and J680}

The mean Fe abundance for J407, 428, and 680 from Table~4 is [Fe/H]~$=-0.18$ with a scatter of 
0.08--0.09~dex, which is comparable to the estimated total uncertainties; the value assuming the revised gravities based on our spectroscopic $T_{\rm eff}$ values is little changed at $-0.16$.  Like our Pleiads, these stars (and the other cool member J682) all evince a small Ni deficiency---[Ni/Fe]~$=-0.08{\pm}0.03$; again this is independent of choice of log $g$.  

The abundance ratios with respect to Fe of all other elements in all three stars, and thus their mean, are solar (i.e., to within a couple hundredths of a dex, [{\it X\/}/Fe]~$=0.$) within the uncertainties; again this is independent of the adopted log $g$.  Consistently low Al, Ti, Ca abundances seen in our Pleiads are not evident for these NGC~2264 PMS stars. 

J428 has an interesting pattern of abundances ([{\it X\/}/H]) with respect to J407 and J680.  All the elements in J428 are consistently ${\sim}0.15$~dex lower than in the other two stars.  The lower abundance also 
is seen for K, which is 0.3~dex lower in J428 relative to J407 and J680. The difference then, does not seem to be due to chance.  Table~4 shows that the difference does not depend on possible choices of log $g$.  The parameter sensitivities given in Table~3 indicate that simple errors in $T_{\rm eff}$ or log $g$ can not simultaneously account for a uniform underabundance in all the elements.   E.g., witness the difference in magnitude of $T_{\rm eff}$ sensitivity of Fe compared to Ti and Ca.  The necessary underestimate in log $g$ (${\sim}1.25$~dex) to bring the Fe abundance into agreement with J407 and J680 is unrealistic.  

A more realistic explanation is a plausible overestimate of the microturbulence which would lower all the abundances.  This effect would  be exacerbated as our iterative analysis would then adopt an errantly low model atmosphere metallicity which would lower all the abundances slightly further still.  Another possibility is binarity.  On the basis of our single epoch modest S/N spectrum it is difficult to exclude a blending SB2 component.  The effects of dilution could errantly lower our measured equivalent widths, resulting in a derived underabundance.  If so, though, we believe the effect to be considerably less than the maximum ${\sim}0.3$~dex effect expected from an identical unblended companion.  This is consistent with a mildly subsolar metal-abundance.

At present, our best estimate from these stars is [Fe/H]~$=-0.15$, a mild Ni deficiency [Ni/Fe]~$=-0.08$ similar to that seen in our Pleiads, and solar ratios (with respect to Fe) of the remaining elements. 

\subsection*{5.2.2. The Probable Non-Member J1088}

Soderblom \etal (1999) ascribed ambiguous but probable non-membership to J1088 on the basis of its H$\alpha$ absorption, discrepant radial velocity, slightly low Li abundance, and color-magnitude diagram position.  This classification seems confirmed by the abundances.  All abundances except Mg are slightly supersolar, with [Fe/H]~$=+0.07$.  Additionally, the spectroscopic gravity of log $g=4.50$ is rather high, and more consistent with a near-ZAMS classification; indeed, the Li abundance and gravity are even consistent with Hyades age values.  

\subsection*{5.2.3. Abundances in J682}

As seen in Figure~2, curious results are found for J682.  Our derived Fe abundance of 
[Fe/H]~$=-0.67{\pm}0.10$ is starkly lower than the other three cool members.  Like the others, Ni shows a small deficiency of [Ni/Fe]~$=-0.11{\pm}0.11$.  The Si and Mg ratios with respect to Fe are solar within the uncertainties.  However, the abundances of the remaining elements (Cr, Ti, Ca, and Al) are significantly higher at [{\it X\/}/H]~${\sim}-0.29$.  This dichotomy is analogous to that seen in our Pleaids, except more extreme and in the opposite sense with lower ionization potential elements showing higher abundances.  Table~4 indicates that this is true regardless of which (substantially different) log $g$ value is adopted in the analysis.  

Like our Pleiads, the source of the difference does not seem to reside in parameter errors.  While the $T_{\rm eff}$ sensitivity of Fe, Ni, Si, and Mg is the same and opposite to that of Cr, Ti, Ca, and Al, no constant $T_{\rm eff}$ offset can bring all the abundances into agreement.  Additionally, any $T_{\rm eff}$ adjustment would have to be implausibly large (${\sim}10^3$~K) to bring just the Fe abundance into agreement with the other three NGC~2264 stars. The uncertainty estimates from the excitation balance analysis suggest considerably smaller values (${\pm}90$~K).  

The photospheric abundance differences thus qualitatively mimic a pattern opposite to interstellar abundances, but also could be interpreted as depending on ionization potential.  Comparing the J682 K abundance in Table~5 with those of J407 and J680 suggests that the [K/H] and [K/Fe] ratio is more aligned with Cr, Ti, Ca, and Al rather than Fe, Ni, Mg, and Si; such alignment was inferred in our Pleiads too.  

Whatever the cause of these intra-photospheric and star-to-star differences, they raise concern about the reliability of LTE analyses with standard model atmospheres in PMS and active ZAMS stars.  A particularly important question is how such deficiencies affect the derived absolute and star-to-star Li abundances in young clusters, which provide critical information concerning stellar physics and Galactic chemical evolution.  Additional abundance determinations from high S/N and resolution spectroscopy of numerous elements in more young cluster stars spanning a large range (${\ga}1000$~K) in $T_{\rm eff}$ are needed to investigate these important issues more substantively.  

\subsection*{5.2.4. Revised Li Abundances} 

The revised Li abundances of our cool NGC~2264 stars presented in Table~5 are higher than those presented in Soderblom \etal (1999) due to our higher spectroscopic $T_{\rm eff}$ estimates.  They noted evidence for star-to-star scatter in NGC~2264 Li, but believed it likely attributable to uncertainties associated with the photometric temperatures.  However, the spectroscopic $T_{\rm eff}$-based abundances seem to either exhibit a 0.20--0.25~dex scatter about a mean (NLTE) value of 3.19, or show a precipitous 0.4~dex drop over $\sim300$~K in $T_{\rm eff}$.  If the former, then a conservative 0.15~dex estimate of the Li abundance uncertainty suggests a modest 0.15--0.20~dex real scatter.  Regardless, either explanation would conflict with standard stellar models of uniform metal abundance, which are unable to produce such scatter or yield such a steep Li decline with $T_{\rm eff}$ at the young age (${\sim}5$~Myr) of NGC~2264.  

\subsection{Presence of the \protect\ion{O}{1} 7774~\AA\ Triplet} 

Figure~3 shows our Keck spectra of HII~676 and J407 in the 7774~\AA\ \ion{O}{1} triplet region.  The expected positions of the \ion{O}{1} triplet lines are marked.  To within ${\sim}0.05$~\AA, each position contains a feature which extends below the noisy pseudocontinuum, and has significant breadth.  Moreover, as commonly observed and expected from reasonably well-determined $gf$ values, the three features decline in strength from blue to red.  The spectrum of HII~97 is similar in all respects.  None of the other cool NGC~2264 stars studied here conclusively demonstrates the presence of the triplet, though the strongest bluest (7771.95~\AA) component of the triplet may be marginally detected. 
%\marginpar{Fig.~3}
\begin{figure*}%3
\centerline{\epsfxsize=4.5in\epsfbox{fig3.eps}}
\caption[]{The spectrum of the ${\lambda}7774$ \ion{O}{1} triplet region in our cool Pleiad HII~676, the cool PMS star NGC~2264 J407, and the older field K~dwarf GL~241.  The expected positions of the $\lambda$7774 \ion{O}{1} triplet are marked.  Each of these positions apparently corresponds to a real feature of significant strength and breadth in our spectrum of each star.  The decreasing line strength from the blue line to the red line is also the behavior expected of the \ion{O}{1} triplet.} 
\end{figure*}

The presence of the lines in our two Pleiads and in J407 is remarkable.  For solar [O/H], these particular features should be vanishingly weak in our spectra.  As summarized in Table~5, however, the measured line strengths suggest a Pleiad O~abundance enhanced by 0.85~dex over solar.  The J407 abundance, [O/H]~$=+0.43$, also represents a large overabundance. Figure~3 also contains a high S/N (${\sim}150$) spectrum of the K6 field dwarf GL~241, which was obtained with the same instrumental configuration as part of DF's dissertation work.  This inactive star has no measurable Li feature, suggesting a much greater age (a Gyr or more) than our NGC~2264 and Pleiades objects.  The spectroscopic $T_{\rm eff}$ (4555~K) and Fe abundance (${\sim}-0.10$) are similar to our cluster stars (\markcite{F98}Fischer 1998).  As can be seen, the \ion{O}{1} triplet is evident in the older cool field dwarf.  The line strengths suggest an abundance [O/H]~${\sim}+0.2$ (Table~5; calculated assuming log g$=4.60$ and ${\xi}=1.2$~km/s).  NLTE effects are believed to affect abundances derived from the \ion{O}{1} triplet in solar-type stars.  However, at the solar $T_{\rm eff}$, gravity, and O~abundance the NLTE corrections are small (${\sim}-0.06$~dex, \markcite{T94}Takeda 1994), and {\it decrease\/} with decreasing $T_{\rm eff}$.  Thus, the value in the K dwarf is surprising---suggesting either a genuine overabundance or shortcomings in the NLTE calculations. 

Because the larger enhancements in our cluster stars seem exceedingly unlikely---particularly in light of the F- and G-star results of \markcite{k93}King (1993)---we can only speculate as to the cause.  One possibility is additional unconsidered NLTE effects due to the presence of a strong chromosphere in our young stars.  Recently, \markcite{yt95}Takeda (1995) has shown that including the chromospheric temperature rise as part of model atmospheres can lead to significant NLTE effects on the formation of the \ion{O}{1} lines {\it in the Sun\/}.  Unfortunately, detailed NLTE line formation computations including the presence of chromospheric layers in model atmospheres appropriate for our cool stars have not been carried out.  If such a mechanism is the root cause, it may be interesting that the overabundance in the older Pleiades is larger than the much younger NGC~2264 star.  

The \ion{O}{1} results demonstrate that LTE analyses with standard model atmospheres are inadequate for some transitions in some PMS and ZAMS stars.  While such a general claim is not particularly surprising, detailed abundances of a variety of elements in numerous young cluster stars are sorely lacking. Thus, the extent to which related or unrelated systematic effects are important factors to consider in such stars remains poorly known.  Such effects may have to be better understood before rigorously reliable absolute and relative abundances of important elements like Li can be derived with full confidence in young stars. 

\section{Summary} 

We derive spectroscopic parameters and abundances of several elements in 2~cool Pleiades dwarfs, 4~cool NGC~2264 PMS stars, and a probable NGC~2264 non-member from high resolution, moderate S/N Keck/HIRES spectra.  The NGC~2264 abundances are some of the first for elements other than Fe and Li 
in PMS stars.  Our modest-sized sample was selected to minimize potential sources of error and uncertainty.  Particular attention was paid to blending concerns.  

Our Pleiades Fe abundance, [Fe/H]~$=+0.06{\pm}0.05$, is between spectroscopic values from F- and G-stars, and in satisfactory agreement with photometric estimates.  The concordance of spectroscopic values is quite reasonable given the range in $T_{\rm eff}$, the differing model atmospheres employed, and the different means (photometric, Balmer profiles, \ion{Fe}{1} fine analysis) by which $T_{\rm eff}$ values have been estimated.  Our abundance does not resolve the 0.3~mag discrepancy in the Pleiades distance modulus derived from main-sequence fitting and {\it Hipparcos\/}parallaxes.  Mg and Si also show similar mild enhancements, but we find a small Ni deficiency ([Ni/Fe]~$=-0.08{\pm}0.03$) in both Pleiads and all our NGC~2264 stars; given the internal uncertainties and the similarity in the parameter sensitivities of Ni and Fe, we tentatively conclude the difference is real. 

The Pleiades ratios of Cr/Ti/Ca/Al are significantly lower ([{\it X\/}/Fe]${\sim}-0.17$); this pattern mimics that of ISM abundances.  While the sensitivity of the Cr/Ti/Ca/Al abundances to $T_{\rm eff}$ is opposite to that of Mg, Si, Fe, and Ni, no simple adjustment in $T_{\rm eff}$ can bring all the results into simultaneous concordance. The low K abundances suggest the pattern is more closely associated with ionization potential; such effects might be related to the Pleiades Li scatter, but this is unclear.  Various evidence suggests real 
intra-cluster [Fe/H] scatter is ${\la}0.05$--0.10 dex; however, a small 0.03~dex [Fe/H] difference can explain the 0.4~dex Li abundance difference between our two cool Pleiads via the metallicity sensitivity of standard PMS Li burning. Investigating intra-cluster metal abundance differences as the source of Pleiades Li abundance scatter would best be accomplished by analyses of low $v\sin i$ Pleiads near 5100~K, where larger and easily detectable metal abundance differences are required to explain the Li scatter.

Three NGC~2264 members provide our best estimate of [Fe/H]~$=-0.15$; except for Ni, the abundance ratios of the remaining elements are solar in these stars. The abundances in J428 are consistently 0.15~dex lower than in J407 and J680. Whether this is a real difference or due to plausible errors in the derived microturbulence, assumed model atmosphere metallicity, and/or spectral dilution from a SB2 component is unclear.  
 
Mildly supersolar abundances for another NGC 2264 star (J1088) are consistent with its probable 
non-membership (Soderblom \etal 1999).  The fourth NGC~2264 member (J682) exhibits curiously low abundances ([{\it X\/}/H]~${\sim}-0.7$) of Mg/Si/Fe/Ni, but its Cr/Ti/Ca/Al abundances are ${\sim}0.3$~dex larger.  This dichotomy is in the opposite sense to the Pleiads'.  The star's K~abundance is also aligned with its Cr/Ti/Ca/Al values, again suggesting a relation to ionization potential.  

A 0.15--0.20~dex scatter or steep decline with $T_{\rm eff}$ is present in the NGC 2264 Li abundances rederived using our $T_{\rm eff}$ values; neither can be accommodated by extant stellar models.  Finally, we note the surprising presence of the ${\lambda}7774$ \ion{O}{1} triplet in our Pleaids, one of the cool NGC~2264 stars, and the K6 field dwarf GL~241.  The inferred LTE O~abundances are enhanced over solar by 0.23--0.85~dex, suggesting that even NLTE calculations of the \ion{O}{1} triplet are incomplete and perhaps implicating the importance of an overlying chromosphere on line formation in cool young stars.  The results demonstrate the potential utility of cluster abundances besides Fe and Li in addressing fundamental issuses concerning stellar evolution and systematic errors in the analysis of cool PMS and ZAMS stars with the assumptions of standard model photospheres and LTE.

\acknow

The observations were made at the W.~M.\ Keck Observatory, which is operated as a scientific partnership between the California Institute of Technology and the University of California, and made possible by the generous financial support of the W.~M.\ Keck Foundation. 

\begin{references}

\reference{bf90} Boesgaard, A.~M., Friel, E.~D. 1990, \apj, 351, 467 

\reference{bbr88} Boesgaard, A.~M., Budge, K.~G., \& Ramsay, M.~E. 1988, \apj, 327, 389

\reference{b89} Boesgaard, A.~M. 1989, \apj, 336, 798 

\reference{c85} Cayrel, R., Cayrel de Strobel, G., \& Campbell, B. 1985, \aap, 146, 249

\reference{ccc88} Cayrel, R., Cayrel de Strobel, G., \& Campbell, B. 1988, in The Impact of Very High S/N Spectroscopy on Stellar Physics, eds.\ G.~Cayrel de Strobel and M.~Spite, (Dordrecht:  Kluwer), p.~449

\reference{cds} Cayrel de Strobel, G. 1990, MmSAI, 61, 613 

\reference{cdp95} Chaboyer, B., Demarque, P., \& Pinsonneault, M.~H. 1995, \apj, 441, 876 (CDP95)

\reference{ct90} Chang, T.~N.\ \& Tang, X. 1990, JQSRT, 43, 207 

\reference{CL} Cunha, K.\ \& Lambert, D.~L. 1994, \apj, 426, 170
 
\reference{dgg92} Demarque, P., Green, E.~M., \& Guenther, D.~B. 1992, \aj, 103, 151

\reference{e86} Eggen, O.~J. 1986, \pasp, 98, 755

\reference{eag93} Edvardsson, B., Andersen, J., Gustafsson, B., Lambert, D. L., Nissen, P.~E., \& Tomkin, J. 1993, \aap, 275, 101

\reference{F98} Fischer, D. 1998, Ph.~D.~Dissertation, University of California Santa Cruz 

\reference{fs} Fitzpatrick, M.~J., \& Sneden, C. 1987, BAAS, 19, 1129

\reference{g92} Gray, D.~F. 1992, in The Observation and Analysis of Stellar Photospheres, Cambridge: Cambridge University Press 

\reference{j96} Jones, B.~F., Shetrone, M., Fischer, D., \& Soderblom, D.~R. 1996, \aj, 112, 186 (JSFS)

\reference{Ket98} King, J.~R., Stephens, A., Boesgaard, A.~M., \& Deliyannis, C.~P. 1998, \aj, 115, 666 

\reference{K98} King, J.~R. 1998, \aj, 116, 254

\reference{kh96} King, J.~R.\ \& Hiltgen, D.~D. 1996, \aj, 112, 2650

\reference{KKP} King, J., Krishnamurthi, A., \& Pinsonneault, M.~H. 2000, \aj, in press 

\reference{k93} King, J.~R. 1993, Ph.~D.~Dissertation, University of Hawaii

\reference{k92} Kurucz, R.~L. 1992, private communication

\reference{k84} Kurucz, R.~L., Furenlid, I., Brault, J., \& Testerman, L. 1984, Solar Flux Atlas from 296 to 1300~nm, NSO Atlas No.~1

\reference{N81} Nissen, P.~E. 1981, \aap, 97, 145

\reference{psskh} Pinsonneault, M.~H., Stauffer, J., Soderblom, D.~R., King, J.~R., \& Hanson, R.~B. 1998, \apj, 504, 170 

\reference{rapps} Randich, S., Aharpour, N., Pallavicini, R., Prosser, C.~F., \& Stauffer, J.~R. 1997, \aap, 323, 86

\reference{s73} Sneden, C. 1973, \apj, 184, 839 
\raggedbottom\newpage

\reference{skhjfs} Soderblom, D.~R., King, J.~R., Hanson, R.~B., Jones, B. F., Fischer, D., Stauffer, J.~R., \& Pinsonneault, M.~H. 1998, \apj, 504, 192 

\reference{S99} Soderblom, D.~R., King, J.~R., Siess, L., Jones, B.~F., \& Fischer, D. 1999, \aj, 118, 1301 

\reference{sfira} Swenson, F.~J., Faulkner, J., Iglesias, C.~A., Rogers, F. J., \& Alexander, D.~R. 1994, \apj, 422, L79

\reference{T94} Takeda, Y. 1994, \pasj, 46, 53

\reference{yt95} Takeda, Y. 1995, \pasj, 47, 463

\reference{T90} Th\'evenin, F. 1990, A\&AS, 82, 179

\reference{taa97} Twarog, B.~A., Ashman, K.~M., \& Anthony-Twarog, B.~J. 1997, \aj, 114, 2556

\reference{u55} {\"U}nsold, A. 1955, Physik der Sternatmospharen, (Berlin:  Springer Verlag)

\reference{vp89} VandenBerg, D.~A.\ \& Poll, H.~E. 1989, \aj, 98, 1451 

\end{references}

\clearpage
\onecolumn

\begin{deluxetable}{lcccc}
\tablewidth{4in}
\tablecaption{Model Atmosphere Parameters and Adopted Uncertainties} 
\tablehead{
\colhead{Star} & \colhead{$T_{\rm eff}$} & \colhead{log $g$} & \colhead{${\xi}$} & \colhead{{[m/H]}}\\ 
\colhead{ } & \colhead{K} & \colhead{cgs} & \colhead{km s$^{-1}$} & \colhead{ }} 
\startdata
\ion{H}{2} 676 & 4410${\pm}75$ & 4.60${\pm}0.15$ & 1.00${\pm}0.25$ & $+0.05{\pm}0.15$ \nl   
\ion{H}{2} 97  & 4525${\pm}75$ & 4.60${\pm}0.15$ & 0.90${\pm}0.25$ & $+0.05{\pm}0.15$ \nl 
J407           & 4660${\pm}90$ & 3.83${\pm}0.50$ & 1.40${\pm}0.25$ & $-0.10{\pm}0.15$ \nl 
J428           & 4300${\pm}90$ & 3.91${\pm}0.50$ & 2.05${\pm}0.25$ & $-0.25{\pm}0.15$ \nl     
J680           & 4250${\pm}90$ & 4.10${\pm}0.50$ & 1.95${\pm}0.25$ & $-0.10{\pm}0.15$ \nl 
J682           & 4570${\pm}90$ & 3.45${\pm}0.50$ & 2.15${\pm}0.25$ & $-0.70{\pm}0.15$ \nl 
J1088          & 5875${\pm}90$ & 4.50${\pm}0.25$ & 1.55${\pm}0.25$ & $+0.05{\pm}0.15$ \nl 
Sun            & 5770 & 4.44 & 1.10 & $+0.00$ \nl
\enddata
\end{deluxetable}

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\begin{deluxetable}{llccccccccc} 
%\scriptsize
\tablenum{3}
\tablecaption{Final Pleiades Abundances and Parameter Sensitivities} 
\tablehead{
\colhead{Ratio} & \colhead{\ion{H}{2} 676} & \colhead{${\sigma}_{\rm int}$} & \colhead{${\sigma}_{\rm tot}$} & \colhead{\ion{H}{2} 97} & \colhead{${\sigma}_{\rm int}$} & \colhead{${\sigma}_{\rm tot}$} & \colhead{${\Delta}T{\rm eff}$} & \colhead{${\Delta}\log{g}$} & \colhead{${\Delta}{\xi}$} & \colhead{${\Delta}{\rm [m/H]}$}\\ 
\colhead{ } & \colhead{ } & \colhead{dex} & \colhead{dex} & \colhead{ } & \colhead{dex} & \colhead{dex} & \colhead{$^{+300}_{-200}$ K} & \colhead{${\pm}0.25$ dex} & \colhead{${\pm}0.5$ km/s} & \colhead{${\pm}0.1$ dex}} 
\startdata
[Fe/H]  & $+0.04$ & 0.01 & 0.07 & $+0.08$ & 0.02 & 0.08 & $^{-0.01}_{+0.06}$ & ${\pm}0.03$ & $^{-0.08}_{+0.06}$ & ${\pm}0.04$\nl 
        & & & & & & & & \nl
[Ni/H]  & $-0.01$ & 0.02 & 0.08 & $-0.03$ & 0.02 & 0.08 & $^{-0.04}_{+0.08}$ & ${\pm}0.04$ & $^{-0.08}_{+0.06}$ & ${\pm}0.04$\nl
[Ni/Fe] & $-0.05$ &      & 0.03 & $-0.11$ &      & 0.03 & & & & \nl 
        & & & & & & & & \nl
[Cr/H]  & $-0.05$ & 0.02 & 0.07 & $-0.09$ & 0.06 & 0.09 & $^{+0.15}_{-0.02}$ & ${\pm}0.00$ & $^{-0.05}_{+0.04}$ & ${\pm}0.03$\nl
[Cr/Fe] & $-0.09$ &      & 0.05 & $-0.17$ &      & 0.08 & & & & \nl 
        & & & & & & & & \nl
[Ti/H]  & $-0.09$ & 0.02 & 0.09 & $-0.07$ & 0.04 & 0.10 & $^{+0.32}_{-0.15}$ & ${\pm}0.00$ & $^{-0.11}_{+0.09}$ & ${\pm}0.02$\nl
[Ti/Fe] & $-0.13$ &      & 0.09 & $-0.15$ &      & 0.10 & & & & \nl 
        & & & & & & & & \nl
[Ca/H]  & $-0.14$ & 0.05 & 0.09 & $-0.09$ & 0.05 & 0.09 & $^{+0.29}_{-0.17}$ & ${\pm}0.04$ & $^{-0.06}_{+0.05}$ & ${\pm}0.02$\nl
[Ca/Fe] & $-0.18$ &      & 0.10 & $-0.17$ &      & 0.10 & & & & \nl
        & & & & & & & & \nl
[Si/H]  & $+0.10$ & 0.05 & 0.10 & $+0.03$ & 0.04 & 0.10 & $^{-0.25}_{+0.23}$ & ${\pm}0.07$ & $^{-0.01}_{+0.01}$ & ${\pm}0.02$\nl
[Si/Fe] & $+0.06$ &      & 0.09 & $-0.05$ &      & 0.09 & & & & \nl
        & & & & & & & & \nl 
[Mg/H]  & $+0.07$ & \llap{$\sim$}0.08 & 0.11 & $+0.03$ & \llap{$\sim$}0.06 & 0.09 & $^{-0.13}_{+0.16}$ & ${\pm}0.03$ & $^{-0.00}_{+0.01}$ & ${\pm}0.03$\nl 
[Mg/Fe] & $+0.03$ &      & 0.10 & $-0.05$ &      & 0.08 & & & & \nl
        & & & & & & & & \nl 
[Al/H]  & $-0.17$ & 0.02 & 0.04 & $-0.15$ & 0.05 & 0.06 & $^{+0.13}_{-0.05}$ & ${\pm}0.00$ & $^{-0.03}_{+0.01}$ & ${\pm}0.01$\nl
[Al/Fe] & $-0.21$ &      & 0.07 & $-0.23$ &      & 0.09 & & & & \nl
        & & & & & & & & \nl 
[Y/H]   & $+0.05$ & \llap{$\sim$}0.10 & 0.16 & \nodata & \nodata & \nodata & $^{+0.45}_{-0.27}$ & ${\pm}0.00$ & $^{-0.14}_{+0.15}$ & ${\pm}0.01$\nl 
[Y/Fe]  & $+0.01$ &              & 0.16 & \nodata & \nodata & \nodata & & & & \nl
        & & & & & & & & \nl
\enddata
\end{deluxetable}

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\def\en{\enspace}
\begin{deluxetable}{lccccccc}
%\scriptsize
\tablenum{5}
\tablecaption{Potassium, Revised Lithium, and Oxygen Abundances} 
\tablehead{
\colhead{Star} & \colhead{EW(K)} & \colhead{[K/H]} & \colhead{EW(Li)} & \colhead{log $N$(Li)} & \colhead{log $N$(Li)} & \colhead{EW(\ion{O}{1})} & \colhead{[O/H]}\\ 
\colhead{ } & \colhead{m{\AA}} & \colhead{ } & \colhead{m{\AA}} & \colhead{LTE} & \colhead{NLTE\tablenotemark{a}} & \colhead{m{\AA}} & \colhead{ }} 
\startdata
\ion{H}{2} 676 & 486 & $-0.37$ & \en23 & \llap{+}0.29 & \llap{+}0.68 & 18.9/13.5/10.0 & $+0.89$/$+0.81$/$+0.82$\nl 
\ion{H}{2} 97  & 447 & $-0.30$ & \en78 & \llap{+}1.04 & \llap{+}1.32 & & \nl 
GL 241         &     &         & & & & 10.8/8.3/5.8 & $+0.23$/$+0.23$/$+0.22$ \nl
J407           & 287 & $-0.24$ & 476 & 3.55 & 3.39 & 26.6/23.6/16.0 & $+0.42$/$+0.49$/$+0.39$\nl 
J428           & 394 & $-0.56$ & 477 & 3.00 & 2.90 & \nodata & \nodata \nl  
J680           & 605 & $-0.25$ & 561 & 3.16 & 3.08 & \nodata & \nodata \nl 
J682           & 305 & $-0.38$ & 507 & 3.50 & 3.38 & \nodata & \nodata \nl 
J1088          & 192 & $+0.29$ & 127 & 2.93 & 2.85 & note\tablenotemark{b} & note\tablenotemark{c} \nl 
Sun            & 154 & \en5.31\tablenotemark{d} & & & & note\tablenotemark{b} & note\tablenotemark{b}\nl 
\enddata
\tablenotetext{a}{Calculated as in Soderblom {\it et~al.\/} (1999); for the Pleiades, $T_{\rm eff}=4500$, log $g=4.5$, [m/H]$=0.0$ is assumed}
\tablenotetext{b}{See Table 2}
\tablenotetext{c}{See Tables 2 and 3} 
\tablenotetext{d}{Absolute solar abundance, log $N$(K), derived with log $gf=-0.16$}
\end{deluxetable}

\end{document}