Abundance analysis of cool extreme helium star: LSS\,3378
Gajendra Pandey and Bacham E. Reddy$^1$
Bangalore 560034
India
$^1$Visiting Observer
Cerro-Tololo Inter-American Observatory
Abundance analysis of the cool extreme helium (EHe) star LSS\,3378 is presented. The abundance analysis is done using LTE line formation and LTE model atmospheres constructed for EHe stars. The atmosphere of LSS\,3378 shows evidence of H-burning, He-burning, and $s$-process nucleosynthesis. The derived abundances of iron-peak and $\alpha$-elements indicate absence of selective fractionation or any other processes that can distort chemical composition of these elements. Hence, the Fe abundance (log $\epsilon$(Fe) = 6.1) is adopted as an initial metallicity indicator. The measured abundances of LSS\,3378 are compared with those of R Coronae Borealis (RCB) stars and with rest of the EHe stars as a group.
lssms.tex
% mnsample.tex
%
% v1.2 released 5th September 1994 (M. Reed)
% v1.1 released 18th July 1994
% v1.0 released 28th January 1994

%\documentstyle[galley,epsf]{mn}
\documentstyle[galley,graphicx,epsf]{mn}
%\documentstyle[referee,epsf]{mn}
% If your system has the AMS fonts version 2.0 installed, MN.sty can be
% made to use them by uncommenting the line: %\AMStwofontstrue
%
% By doing this, you will be able to obtain upright Greek characters.
% e.g. \umu, \upi etc.  See the section on "Upright Greek characters" in
% this guide for further information.
%
% If you are using AMS 2.0 fonts, bold math letters/symbols are available
% at a larger range of sizes for NFSS release 1 and 2 (using \boldmath or
% preferably \bmath).

\newif\ifAMStwofonts
%\AMStwofontstrue

%%%%% AUTHORS - PLACE YOUR OWN MACROS HERE %%%%%



%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%

\ifoldfss
  %
  \newcommand{\rmn}[1] {{\rm #1}}
  \newcommand{\itl}[1] {{\it #1}}
  \newcommand{\bld}[1] {{\bf #1}}
  %
  \ifCUPmtlplainloaded \else
    \NewTextAlphabet{textbfit} {cmbxti10} {}
    \NewTextAlphabet{textbfss} {cmssbx10} {}
    \NewMathAlphabet{mathbfit} {cmbxti10} {} % for math mode
    \NewMathAlphabet{mathbfss} {cmssbx10} {} %  "   "    "
  \fi
  %
  \ifAMStwofonts
  %
    \ifCUPmtlplainloaded \else
      \NewSymbolFont{upmath} {eurm10}
      \NewSymbolFont{AMSa} {msam10}
      \NewMathSymbol{\upi}     {0}{upmath}{19}
      \NewMathSymbol{\umu}     {0}{upmath}{16}
      \NewMathSymbol{\upartial}{0}{upmath}{40}
      \NewMathSymbol{\leqslant}{3}{AMSa}{36}
      \NewMathSymbol{\geqslant}{3}{AMSa}{3E}
      \let\oldle=\le     \let\oldleq=\leq
      \let\oldge=\ge     \let\oldgeq=\geq
      \let\leq=\leqslant \let\le=\leqslant
      \let\geq=\geqslant \let\ge=\geqslant
    \fi
    %
  \fi
%
\fi % End of OFSS

\ifnfssone
  %
  \newmathalphabet{\mathit}
  \addtoversion{normal}{\mathit}{cmr}{m}{it}
  \addtoversion{bold}{\mathit}{cmr}{bx}{it}
  %
  \newcommand{\rmn}[1] {\mathrm{#1}}
  \newcommand{\itl}[1] {\mathit{#1}}
  \newcommand{\bld}[1] {\mathbf{#1}}
  %
  \def\textbfit{\protect\txtbfit}
  \def\textbfss{\protect\txtbfss}
  \long\def\txtbfit#1{{\fontfamily{cmr}\fontseries{bx}\fontshape{it}%
    \selectfont #1}}
  \long\def\txtbfss#1{{\fontfamily{cmss}\fontseries{bx}\fontshape{n}%
    \selectfont #1}}
  %
  \newmathalphabet{\mathbfit} % math mode version of \textbfit{..}
  \addtoversion{normal}{\mathbfit}{cmr}{bx}{it}
  \addtoversion{bold}{\mathbfit}{cmr}{bx}{it}
  %
  \newmathalphabet{\mathbfss} % math mode version of \textbfss{..}
  \addtoversion{normal}{\mathbfss}{cmss}{bx}{n}
  \addtoversion{bold}{\mathbfss}{cmss}{bx}{n}
  %
  \ifAMStwofonts
    %
    \ifCUPmtlplainloaded \else
      %
      % Make NFSS 1 use the extra sizes available for bold math italic and
      % bold math symbol. These definitions may already be loaded if your
      % NFSS format was built with fontdef.max.
      %
      \UseAMStwoboldmath
      %
      \makeatletter
      \new@mathgroup\upmath@group
      \define@mathgroup\mv@normal\upmath@group{eur}{m}{n}
      \define@mathgroup\mv@bold\upmath@group{eur}{b}{n}
      \edef\UPM{\hexnumber\upmath@group}
      %
      \new@mathgroup\amsa@group
      \define@mathgroup\mv@normal\amsa@group{msa}{m}{n}
      \define@mathgroup\mv@bold\amsa@group{msa}{m}{n}
      \edef\AMSa{\hexnumber\amsa@group}
      \makeatother
      %
      \mathchardef\upi="0\UPM19
      \mathchardef\umu="0\UPM16
      \mathchardef\upartial="0\UPM40
      \mathchardef\leqslant="3\AMSa36
      \mathchardef\geqslant="3\AMSa3E
      %
      \let\oldle=\le     \let\oldleq=\leq
      \let\oldge=\ge     \let\oldgeq=\geq
      \let\leq=\leqslant \let\le=\leqslant
      \let\geq=\geqslant \let\ge=\geqslant
      %
    \fi
  \fi
  %
\fi % End of NFSS release 1

\ifnfsstwo
  %
  \newcommand{\rmn}[1] {\mathrm{#1}}
  \newcommand{\itl}[1] {\mathit{#1}}
  \newcommand{\bld}[1] {\mathbf{#1}}
  %
  \def\textbfit{\protect\txtbfit}
  \def\textbfss{\protect\txtbfss}
  \long\def\txtbfit#1{{\fontfamily{cmr}\fontseries{bx}\fontshape{it}%
    \selectfont #1}}
  \long\def\txtbfss#1{{\fontfamily{cmss}\fontseries{bx}\fontshape{n}%
    \selectfont #1}}
  %
  \DeclareMathAlphabet{\mathbfit}{OT1}{cmr}{bx}{it}
  \SetMathAlphabet\mathbfit{bold}{OT1}{cmr}{bx}{it}
  \DeclareMathAlphabet{\mathbfss}{OT1}{cmss}{bx}{n}
  \SetMathAlphabet\mathbfss{bold}{OT1}{cmss}{bx}{n}
  %
  \ifAMStwofonts
    %
    \ifCUPmtlplainloaded \else
      \DeclareSymbolFont{UPM}{U}{eur}{m}{n}
      \SetSymbolFont{UPM}{bold}{U}{eur}{b}{n}
      \DeclareSymbolFont{AMSa}{U}{msa}{m}{n}
      \DeclareMathSymbol{\upi}{0}{UPM}{"19}
      \DeclareMathSymbol{\umu}{0}{UPM}{"16}
      \DeclareMathSymbol{\upartial}{0}{UPM}{"40}
      \DeclareMathSymbol{\leqslant}{3}{AMSa}{"36}
      \DeclareMathSymbol{\geqslant}{3}{AMSa}{"3E}
      %
      \let\oldle=\le     \let\oldleq=\leq
      \let\oldge=\ge     \let\oldgeq=\geq
      \let\leq=\leqslant \let\le=\leqslant
      \let\geq=\geqslant \let\ge=\geqslant
      %
    \fi
  \fi
  %
\fi % End of NFSS release 2

\ifCUPmtlplainloaded \else
  \ifAMStwofonts \else % If no AMS fonts
    \def\upi{\pi}
    \def\umu{\mu}
    \def\upartial{\partial}
  \fi
\fi

\begin{document}
\title{Abundance analysis of cool extreme helium star: LSS\,3378}
\author[G. Pandey]
       {Gajendra Pandey and Bacham E. Reddy$^1$\\
       Indian Institute of Astrophysics, Bangalore 560034, India\\
       $^1$Visiting Observer, Cerro-Tololo Inter-American Observatory}
\date{Accepted .
      Received ;
      in original form 2006 }

\pagerange{\pageref{firstpage}--\pageref{lastpage}}
\pubyear{2006}


\maketitle

\label{firstpage}

\begin{abstract}

Abundance analysis of the cool extreme helium (EHe) star LSS\,3378 is presented.
The abundance analysis is done using LTE line formation and
LTE model atmospheres constructed for EHe stars.

The atmosphere of LSS\,3378 shows evidence of H-burning, 
He-burning, and $s$-process nucleosynthesis. The derived abundances
of iron-peak and $\alpha$-elements indicate absence of selective fractionation
or any other processes that can distort chemical composition of these
elements. Hence, the Fe abundance (log $\epsilon$(Fe) = 6.1)
is adopted as an initial metallicity indicator.
The measured abundances of LSS\,3378 are
compared with those of R Coronae Borealis (RCB) stars and with rest of the EHe stars as a group.

\end{abstract}

\begin{keywords}
stars: abundances $-$ stars: AGB and post-AGB $-$ stars: chemically peculiar $-$ stars: evolution.
\end{keywords}


\section{Introduction}

The extreme helium stars (EHes) are supergiants with peculiar chemical
composition. The atmospheres of these supergiants with effective
temperature in the range 8000 -- 35000 K are devoid of hydrogen, and are
enriched with helium, carbon, and nitrogen with respect to the
atmospheres of normal main-sequence stars. Helium is the 
most abundant element in their atmospheres. There are about 21 known EHes.
There are 5 cool EHes, stars with effective temperatures 8000 to 13000 K, of
which 4 were analysed by Pandey et al. (2001). The remaining one cool EHe LSS\,3378
is analysed here.

Star No. 3378 (LSS\,3378) in the catalog of Stephenson and Sanduleak (1971)
was first identified as helium-rich B-type by Drilling (1973). Drilling
noticed the absence of Balmer absorption lines, presence of strong He\,{\sc i}
absorption lines, and the presence of C\,{\sc ii} line at 4267\AA\ in the
spectrum of LSS\,3378. Drilling also noted strong Mg\,{\sc ii} line at
4481\AA\ and Si\,{\sc ii} lines in absorption indicating a spectral class of
about B8 but, however, suggested somewhat earlier class than this because of the presence
of several weak O\,{\sc ii} lines.
Jeffery et al. (2001) reported the effective
temperature of LSS\,3378 to be about 10500 K by fixing $E_{B-V}$
using $IUE$ data and $UBV$ photometry. However, Jeffery et al. also reported that
the $IUE$ data for this star are very noisy.
Drilling et al. (1984) were the first to derive 
an effective temperature for LSS\,3378, obtaining 9400$\pm$500K, using the same 
procedure adopted by Jeffery et al.

R Coronae Borealis (RCB) stars, which are hydrogen-deficient F- and G-type supergiants, 
overlap in their effective temperatures with cooler EHes. With this abundance
analysis of LSS\,3378, we have the chemical composition of all the 5 known
cool EHes which are closely related to RCB stars. Our abundance analysis is based on
the high resolution optical spectrum.

\section{Observations}

The spectrum of LSS\,3378 was obtained on 2002 June 19 with the 4-m Blanco telescope
and the Cassegrain echelle spectrograph at CTIO in Chile. Three exposures of
25 minutes each, spectra covering the wavelength interval 5000 -- 8200\AA\ without gaps,
were recorded at a resolving power of R = 30,000.
The Image Reduction and Analysis Facility (IRAF) software packages were used
to reduce the recorded spectra.
The Th-Ar hollow cathode lamp provided lines for wavelength calibration.
A final spectrum was obtained by co-adding these three individual wavelength calibrated 
spectra from the three exposures. The maximum signal-to-noise (S/N) in the 
continuum (per pixel) of the co-added spectrum is between 200 and 250 at the 
middle of each echelle order.

The Na\,D lines in the spectrum of LSS 3378 are very strong and appear to
be saturated as seen in the spectra of other cool EHes: FQ\,Aqr, LS IV $-14^\circ109$,
BD $-1^\circ3438$, and LS IV $-1^\circ2$. These saturated Na\,D lines in the spectra of cool
EHes are certainly interstellar in origin. The Na\,D lines in the spectra of
cool EHes are stronger than in any other spectrum of normal star. Note that,
lines of ionized metals of the iron group are plentiful in the cool EHe's spectrum.
These lines are much stronger when compared with those observed in early A-type
and late B-type normal stars. This notable feature of the spectra of cool EHes
is attributable to the lower opacity in the atmosphere due to hydrogen deficiency.

%\begin{table*}
%\centering
%\begin{minipage}{90mm}
%\caption{Table of observations; the stars are ordered by their effective temperatures decreasing from top to bottom}
%\begin{tabular}{lrrrrr}
%\hline
%\multicolumn{1}{c}{Star}&\multicolumn{1}{c}{Date}&\multicolumn{1}{c}{R = }&\multicolumn{1}{c}{S/N}&
%\multicolumn{1}{c}{S/N} &\multicolumn{1}{c}{$V_{rad}$}\\
%&&\multicolumn{1}{c}{$\lambda/\Delta\lambda$}& \multicolumn{1}{c}{at 8727 \AA} 
%& \multicolumn{1}{c}{at 9850 \AA} & \multicolumn{1}{c}{km s$^{-1}$}\\
%\hline
%RCB stars:  &                &          &            &     & \\
%XX Cam     &   9 Oct 1997   &  60,000  &  ...       & 185 &11\\
%XX Cam     &  26 Jan 1998   &  60,000  & ...   & 120 &13\\
%XX Cam     &  17 Nov 2002  &  60,000  & ... & 186&11\\
%XX Cam     &  17 Nov 2002  &  60,000  & 278 & ...&...\\
%RY Sgr     &  22 June 1997  &  60,000  & ...  & 170 & $-$30\\
%RY Sgr     &  31 July 2002  &  120,000  & 300  & 200&$-$12\\
%UV Cas     &  24 July 1996  &  60,000  & ...   & 140 &$-$32\\
%UV Cas     &  15 Nov 2002  &  60,000  &  ...  & 120 &$-$32\\
%UV Cas     &  15 Nov 2002  &  60,000  & 178   & ...&...\\
%VZ Sgr     &  18 Aug 1999   &  60,000  & ...   & 80 &245\\
%R CrB      &  17 June 1995  &  60,000  &   340                & 180&20\\
%R CrB      &   7 Aug 1995  &  60,000  &    ...               & 110&22\\
%R CrB      &  20 June 1997  &  60,000  &   ...      & 185 &21\\
%R CrB      &  31 July 2002  &  120,000  & 400  & 200 &19\\
%SU Tau     &  15 Nov 2002   &  60,000  & ...   & 100    &43\\
%SU Tau     &  15 Nov 2002   &  60,000  & 180  & ...  &...\\
%V482 Cyg   &  24 July 1996  &  60,000  & ...   & 90  &$-$40\\
%GU Sgr     &  22 June 1997  &  60,000  & ...   & 95  &$-$45\\
%Standard star: &            &          &       &     &   \\
%$\gamma$ Cyg& 31 July 2002  &  120,000  &    520            & 230 &...\\
%\hline
%\end{tabular}
%%$^a$The stars are ordered based on their increasing effective temperatures\\
%%$^a$covers 8727\AA\ region\\
%%$^b$covers the 8727\AA\ and 9850\AA\ regions
%\end{minipage}
%\end{table*}

\section{Abundance analysis - method}

\subsection{Procedure}

For the abundance analysis of LSS\,3378, same procedure is followed as
described in Pandey et al. (2001, 2004, 2006). The analysis 
uses line-blanketed hydrogen-deficient model atmospheres computed 
by the code STERNE (Jeffery, Woolf \& Pollacco 2001).
STERNE model atmosphere was combined with the Armagh LTE code 
SPECTRUM (Jeffery, Woolf \& Pollacco 2001) to compute the equivalent width of a line
or a synthetic spectrum. In matching
a synthetic spectrum to an observed spectrum we include broadening
due to the instrumental profile, the microturbulent velocity $\xi$ and assign all
additional broadening, if any, to rotational broadening.
We use the standard rotational broadening
function $V(v\sin i,\beta)$ with
the limb darkening coefficient set at $\beta = 1.5$ (Jeffery, Woolf \& Pollacco 2001).
Observed unblended line profiles are used to obtain the projected rotational
velocity $v\sin i$. Synthetic line profile, including
the broadening due to instrumental profile, for
the adopted model atmosphere ($T_{\rm eff}$,$\log g, \xi$) and the abundance,
is found to be sharper than the observed. This extra broadening in the observed profile
is attributed to rotational broadening. 

The adopted $gf$-values for C, N, O, are from Wiese, Fuhr \& Deters (1996), and
for rest of the elements are from 
NIST database\footnote{http://physics.nist.gov/PhysRefData/ASD/lines\_form.html},
Kurucz's database\footnote{http://kurucz.harvard.edu}, Th\'{e}venin (1989, 1990),
and the compilations by R. E. Luck (private communication).
The Stark broadening and radiative broadening coefficients, if available, are mostly
taken from the Kurucz's database. 
The data  for
computing He\,{\sc i} profiles are the same as in Jeffery, Woolf \& Pollacco (2001).
Table A1 has the detailed line list used in our analysis.

\subsection{Atmospheric parameters}

The model atmospheres are defined by the effective temperature, the
surface gravity, and the chemical composition. The input composition of
He and C (the C/He ratio) is fully consistent with the C/He ratio derived from
the observed spectrum with that model; He and C abundances, particularly He,
dominate the continuous opacity. The input composition of rest of the elements
is solar with H/He fixed at 10$^{-4}$ by number.

The analysis involves the determination of effective temperature ($T_{\rm eff}$),
surface gravity ($\log g$), and microturbulent velocity ($\xi$) before estimating
the photospheric elemental abundances of the star. These parameters are determined 
from the line spectrum. The microturbulent velocity $\xi$ (in km s$^{-1}$) is first 
determined by the requirement that the abundance from
a set of lines of the same ion with similar excitation potentials be independent of 
a line's equivalent width.
 For an element represented in the
spectrum by two or more ions, imposition of ionization
equilibrium (i.e., the same abundance is required from lines of
different stages of ionization) defines a locus in the
($T_{\rm eff},\log g)$ plane.
 Different pairs of ions of a common element provide
loci of very similar slope in the ($T_{\rm eff},\log g)$ plane.

An indicator yielding a locus with a contrasting slope
in the ($T_{\rm eff},\log g)$ plane is required to break the
degeneracy presented by ionization equilibria.
A potential indicator is a He\,{\sc i} line.
For stars hotter than about 10,000 K,
the  He\,{\sc i} lines are less sensitive to $T_{\rm eff}$
than to $\log g$ on account of pressure broadening due to
the quadratic Stark effect.
The diffuse series lines are, in particular, useful because they are
less sensitive to the microturbulent velocity than the sharp lines.
A second indicator may be available: species represented
by lines spanning a range in excitation potential may serve as
a thermometer measuring $T_{\rm eff}$ with a weak dependence
on $\log g$.

\section{LSS\,3378 - abundance analysis results}
 
First to determine is the microturbulent velocity $\xi$. We adopt a model atmosphere
with $T_{\rm eff}$ = 10000 K, which is close to that found by Jeffery et al. (2001),
and adopt $\log g$ = 1.0
which is a fair assumption for these hydrogen-deficient supergiants.  $\xi$ is found
to be 4.5 and 7.5 km s$^{-1}$ from Fe\,{\sc ii} and C\,{\sc i} lines, respectively.
We adopt $\xi$ = 6 km s$^{-1}$ for abundance determination, and Figure 1 illustrates 
the method for obtaining $\xi$. Fe\,{\sc ii} lines used for determining $\xi$ were of 
similar lower excitation potential and so was the case for C\,{\sc i} lines. 

\begin{figure}
\epsfxsize=8truecm
\epsffile{lssfig1.ps}
\caption{Abundances from Fe\,{\sc ii} lines for LSS\,3378 versus their
reduced equivalent widths (log $W_{\lambda}/\lambda$).
A microturbulent velocity of $\xi \simeq 4.5$ km s$^{-1}$ is obtained from this
figure.}
\end{figure}

\begin{figure}
\epsfxsize=8truecm
\epsffile{lssfig2.ps}
\caption{Excitation equlilbrium for LSS\,3378 using Fe\,{\sc ii} lines, and
models with $\log g$ = 0.5.}
\end{figure}

\begin{figure}
\epsfxsize=8truecm
\epsffile{lssfig3.ps}
\caption{The $T_{\rm eff}$ vs $\log g$ plane for LSS\,3378. Loci satisfying
ionization equilibria are plotted -- see key on the figure.
The locus satisfying the excitation balance of Fe\,{\sc ii}
lines is shown by thick solid line.
The cross shows the adopted model atmosphere parameters.}
\end{figure}
 
\begin{figure}
\epsfxsize=8truecm
\epsffile{lssfig4.ps}
\caption{LSS\,3378's observed and synthesized He\,{\sc i} line profile
at 6678.15\AA. The He\,{\sc i} line profiles are synthesized using the model
$T_{\rm eff}$ = 10600 K and $\log g$ = 0.4, for different values of C/He - see key
on the figure.} 
\end{figure}

 The $T_{\rm eff}$ was estimated from Fe\,{\sc ii} lines spanning excitation 
potentials from 3 eV to 11 eV. For $\xi$ = 4.5 km s$^{-1}$,
models ($T_{\rm eff}$, $\log g$) were found which gave the same abundance independent
of excitation potential. The $T_{\rm eff}$ determined from Fe\,{\sc ii} lines
somewhat depends on the adopted surface gravity. 
Figure 2 illustrates the method for obtaining $T_{\rm eff}$.
Ionization equilibrium loci 
for C\,{\sc i}/C\,{\sc ii} and N\,{\sc i}/N\,{\sc ii} are shown in Figure 3.
These with the $T_{\rm eff}$ 
determined from the excitation equilibrium of Fe\,{\sc ii} lines, which is a function 
of surface gravity, gives the estimate of the
stellar parameters: $T_{\rm eff}$ = 10600$\pm$250 K, $\log g$ = 0.4$\pm$0.25 cgs, and
$\xi$ = 6$\pm$2 km s$^{-1}$ (Figure 3). Thus, the abundance analysis was conducted
for the model atmosphere (10600,0.4,6.0).
At this $T_{\rm eff}$ which is close to 11000 K, electron scattering (most of
the free electrons from photoionization of neutral helium), and photoionization of
neutral helium are the major sources of continuous opacity in the line forming
regions (Pandey et al. 2001).  Hence, for LSS\,3378 it is evident that helium controls
the continuum opacity, and the C/He ratio is directly determined 
from the measured equivalent widths of C\,{\sc i} and C\,{\sc ii}
lines observed in the spectrum of LSS\,3378.
The C\,{\sc i} lines give log $\epsilon$(C) = 9.53$\pm$0.19. Four C\,{\sc ii}
lines give 9.44$\pm$0.11. The abundances from C\,{\sc i} and C\,{\sc ii} lines
are in good agreement, and imply a C/He = 1\%.
At the $T_{\rm eff}$ of LSS\,3378, the He\,{\sc i} lines are mildly sensitive to the
C/He ratio. The C/He ratio may be derived by
fitting the He\,{\sc i} lines at 5048, 5876, and 6678\AA\ (Figure 4); the blending
Ne\,{\sc i} line at 6678\AA\ is taken into account.
These He\,{\sc i} lines give C/He = 0.8$\pm$0.1\%, where the uncertainty reflects only the scatter
of the three results.
LSS\,3378's final abundances, for key elements, derived for C/He = 1\% are 
summarized in Table 1; also given are solar abundances from Table 2 of Lodders (2003) 
for comparison.
The individual elemental abundances listed in Table 1 are given as
log $\epsilon(i)$, normalized such that log $\Sigma$$\mu_i \epsilon(i)$ = 12.15
where $\mu_i$ is the atomic weight of element $i$.
The lines used for abundance analysis including the mean abundance, and the
line-to-line scatter are given in Table A1. The abundance errors due to the
uncertainty in the adopted stellar parameters are given in Table 2.
The deduced $v\sin i$ is about 26 km s$^{-1}$.

\begin{table*}
\caption{Adopted key elemental abundances for LSS\,3378}
\begin{center}
%{\small
\begin{tabular}{lrcccccccccccccccc} \hline
Star & H & He & C & N & O & Mg & Al & Si & P & S & Ti & Cr & Mn & Fe & Y & Zr & Ba \\
\hline
LSS\,3378 & 7.18 & 11.50 & 9.46 & 8.26 & 9.29 & 5.97 & 5.98 & 6.62 & 4.85 & 6.52 & 4.29 & 4.47 &
4.57 & 6.11 & 2.85 & 3.51 & 2.49 \\
Sun       & 12.00& 10.98 & 8.46 & 7.90 & 8.76 & 7.62 & 6.54 & 7.61 & 5.54 & 7.26 & 5.00 & 5.72 & 
5.58 & 7.54 & 2.28 & 2.67 & 2.25 \\
\hline
\end{tabular}%}
\end{center}
\end{table*}

\section{Abundances}

%The derived abundances of LSS\,3378 are very similar to other EHes and
%this star belongs to the O-rich group of EHes (see Figures ...).
The derived abundances of the EHe LSS\,3378 are compared with the measured
abundances, taken from the literature, of the other EHes (see Pandey et al. 2006).
The photospheric composition of LSS\,3378 reveals that the surface material
is contaminated by the products of H-burning, and He-burning reactions, as observed
for most of the EHes.

LSS\,3378's abundance ratios: Cr/Fe, Mn/Fe, S/Fe, Si/Fe, and Ti/Fe, are
as expected for that in metal-poor normal and unevolved stars except for the low
Mg/Fe ratio. The abundances
of Fe, Cr, and Mn$-$the iron peak elements, and S, Si, and Ti$-$the 
$\alpha$-elements represent the initial metallicity of LSS\,3378 as these elements are 
expected to be unaffected by H- and the He-burning, and
attendant nuclear reactions.
We choose Fe to be the indicator of the initial metallicty in
LSS\,3378 for spectroscopic convenience.

LSS\,3378's derived elemental abundances that are affected by evolution
are of H, C, N, O, Ne, Y, Zr, and Ba. Note that, the Ne abunadnaces from 
Ne\,{\sc i} lines are affected by
departures from LTE, hence, not discussed here. However, the neon LTE abunadnces
from Ne\,{\sc i} lines are listed in Table A1. \\
{\it Hydrogen}$-$H abundance log $\epsilon$(H) is about 7.2 that fits the suggested trend
of increasing H with increasing $T_{\rm eff}$ for all the EHes with C/He ratio of
about 1\%, the exception being the hottest EHe 
LS IV $+6^\circ2$ (see Pandey et al. 2006).\\
{\it Carbon}$-$The C/He ratio is 0.0083, the mean C/He ratio from 15 EHes, that excludes
the two EHes HD\,144941 and V652\,Her with much lower C/He ratio, is 0.0066.\\
{\it Nitrogen}$-$Nitrogen is enriched above its initial abundance expected based
on the Fe abundance. The observed N abundance is the result of complete
conversion of the initial C, N, and O to N via H-burning CN-cycle and the ON-cycles.\\
{\it Oxygen}$-$Oxygen abundance relative to Fe is overabundant by about 1.5 dex.
LSS\,3378 fits to the oxygen-rich group of EHes (see Pandey et al. 2006).\\
{\it Yttrium, Zr, {\rm and} Ba}$-$Relative to iron, Y, Zr, and Ba are overabundant with
respect to solar by about a factor of 80 (1.9 dex), 150 (2.2 dex), and 40 (1.6 dex),
respectively. Note that, the measured abundances are are based on one or two lines.
The observed Y, Zr, and Ba overabundances are attributed to contamination of the
atmosphere by $s$-process products.

\section{Conclusions}

The abundances of the cool EHe star LSS\,3378 are measured. The measured C/He ratio
is about 1\% similar to most of the EHes. LSS\,3378 fits the oxygen-rich group
of EHes discussed by Pandey et al. (2006). The measured abundances of Y, Zr, and Ba
suggest that in the EHe star LSS\,3378, $s$-process nucleosynthesis did occur in its
earlier evolution. With this analysis, a total of four EHes: LSS\,3378, PV\,Tel,
V1920\,Cyg, and LSE\,78, show a strong enhancement of Y and Zr attributable to an
$s$-process (Pandey et al. 2004, 2006).

An interesting similarity is suggested by the $s$-process abundances in LSS\,3378
and those of the RCB stars. The light $s$-process elements are more enhanced than the
heavy $s$-process elements in RCB stars (Asplund et al. 2000; Rao \& Lambert 2003;
Vanture, Zucker, \& Wallerstein 1999; Bond, Luck, \& Newman 1979). 
Similar enhancement of light $s$-process elements (Y, Zr) over heavy $s$-process 
element (Ba) is observed in LSS\,3378.

Enrichment of $s$-process elements is not expected for the EHes resulting from
a merger of a He with a C-O white dwarf as discussed by Pandey et al. (2004, 2006).
However, synthesis by neutrons via the $s$-process may occur during the merger and
needs to be explored.
 
To further improve the chemical analysis, non-LTE calculations should be performed
for key elements like N, O, Ne, Si, S, and Fe, and in particular neon.

 
\begin{table}
%\centering
%\begin{minipage}{75mm}
\caption{Abundance errors due to uncertainties in the stellar parameters: $\Delta$$T_{\rm eff}$,
$\Delta$$\log g$, and $\Delta$$\xi$. The abundance error due to $\Delta$$T_{\rm eff}$ is the difference
in abundances derived from the adopted model ($T_{\rm eff}$, $\log g$, $\xi$) and 
a model ($T_{\rm eff}$+$\Delta$$T_{\rm eff}$, $\log g$, $\xi$). The abundance error due to
$\Delta$$\log g$ is the difference in abundances derived from the adopted
model ($T_{\rm eff}$, $\log g$, $\xi$) and a model ($T_{\rm eff}$, $\log g$+$\Delta$$\log g$, $\xi$).
The abundance error due to $\Delta$$\xi$ is the difference in abundances derived from the adopted
model ($T_{\rm eff}$, $\log g$, $\xi$) and a model ($T_{\rm eff}$, $\log g$, $\xi$+$\Delta$$\xi$).}
\begin{tabular}{lccc} \hline
\multicolumn{1}{l}{Species}&\multicolumn{1}{c}{$\Delta$$T_{\rm eff}$ = +300}&
\multicolumn{1}{c}{$\Delta$$\log g$ = +0.25}& 
\multicolumn{1}{c}{$\Delta$$\xi$ = +2.0} \\
\multicolumn{1}{c}{}& \multicolumn{1}{c}{[K]} &
\multicolumn{1}{c}{[cgs]} &\multicolumn{1}{c}{[km s$^{-1}$]}\\
\hline
H\,{\sc i} & $-$0.19 & +0.22 & +0.09\\
C\,{\sc i} & $-$0.18 & +0.18 & +0.11\\
C\,{\sc ii} & +0.02 & +0.02 & +0.13\\
N\,{\sc i} & $-$0.18 & +0.18 & +0.12\\
N\,{\sc ii} & +0.04 & $-$0.02 & +0.09\\
O\,{\sc i} & $-$0.17 & +0.20 & +0.24\\
Ne\,{\sc i} & $-$0.04 & +0.11 & +0.41\\
Na\,{\sc i} & $-$0.17 & +0.17 & +0.09\\
Mg\,{\sc ii} & $-$0.22 & +0.18 & +0.28\\
Al\,{\sc ii} & $-$0.12 & +0.11 & +0.17\\
Al\,{\sc iii} & +0.04 & $-$0.05 & +0.09\\
Si\,{\sc ii} & $-$0.15 & +0.16 & +0.26\\
P\,{\sc ii} & $-$0.07 & +0.03 & +0.05\\
S\,{\sc ii} & 0.00 & 0.00 & +0.20\\
%Sc\,{\sc ii} & $-$0.26 & +0.18 & +0.01\\
Ti\,{\sc ii} & $-$0.27 & +0.18 & +0.03\\
Cr\,{\sc ii} & $-$0.35 & +0.17 & $-$0.05\\
Mn\,{\sc ii} & $-$0.22 & +0.17 & +0.03\\
Fe\,{\sc ii} & $-$0.25 & +0.15 & +0.11\\
Y\,{\sc ii} & $-$0.27 & +0.19 & +0.01\\
Zr\,{\sc ii} & $-$0.28 & +0.19 & $-$0.01\\
Ba\,{\sc ii} & $-$0.25 & +0.19 & +0.01\\
\hline
\end{tabular}
%\end{minipage}
\end{table}


The observations presented here were obtained at CTIO, National
Optical Astronomy Observatories (NOAO), which is operated by the Association of
Universities for Research in Astronomy Inc. (AURA) under a cooperative agreement
with the National Science Foundation, USA. We thank the referee Tony Lynas-Gray
for drawing our attention to the strong Na\,D profiles in LSS\,3378's spectrum.

\begin{thebibliography}{99}

\bibitem{} Asplund, M., Gustafsson, B., Lambert, D. L., Rao, N. K., 2000, A\&A, 353, 287
\bibitem{} Bond, H. E., Luck, R. E., Newman, M. J., 1979, ApJ, 233, 205 
\bibitem{} Drilling, J. S., 1973, ApJ, 179, L31
\bibitem{} Drilling, J. S., Sch\"{o}nberner, D., Heber, U., Lynas-Gray, A. E., 1984, ApJ, 278, 224
\bibitem{} Jeffery, C. S., Starling, R. L. C., Hill, P. W., Pollacco, D. L., 2001, MNRAS, 321, 111
\bibitem{} Jeffery, C. S., Woolf, V. M., Pollacco, D. L., 2001, A\&A, 376, 497
\bibitem{} Pandey, G., Rao, N. K., Lambert, D. L., Jeffery, C. S., Asplund, M., 2001, MNRAS, 324, 937
\bibitem{} Pandey, G., Lambert, D. L., Rao, N. K., Jeffery, C. S., 2004, ApJ, 602, L113
\bibitem{} Pandey, G., Lambert, D. L., Jeffery, C. S., Rao, N. K., 2006, ApJ, 638, 454
\bibitem{} Rao, N. K., Lambert, D. L., 2003, PASP, 115, 1304 
\bibitem{} Stephenson, C. B., Sanduleak, N., 1971, Publ. Warner \& Swasey Obs. 1, 1
\bibitem{} Th\'{e}venin, F., 1989, A\&AS, 77, 137
\bibitem{} Th\'{e}venin, F., 1990, A\&AS, 82, 179
\bibitem{} Vanture, A. D., Zucker, D., Wallerstein, G., 1999, ApJ, 514, 932 
\bibitem{} Wiese, W. L., Fuhr, J. R., Deters, T. M., 1996, {\it J. Phys. Chem. Ref. Data}, Monograph No. 7

\end{thebibliography}


\appendix

\section{Lines used for abundance analysis}

The lines used for the abundance analysis of LSS\,3378
are given in Table A1. The lower excitation potential ($\chi_L$),
$gf$-value, measured equivalent width ({\em $W_{\lambda}$}) and the 
abundance (log $\epsilon$) derived for each line is also
listed. The abundances are derived using $C/He$ = 1.0\% model for 
LSS\,3378.

\begin{table}
\caption{Lines used to derive elemental abundances for LSS\,3378}
\begin{tabular}{lrrrrl}
\hline
Ion&&&{\em W$_{\lambda}$}&&\\
$\lambda$ (\AA)& $\chi$ (eV)& log $gf$& (m\AA)& log $\epsilon$$^a$
& Ref.$^b$\\
\hline
H\,{\sc i}&&&&&\\
%No. of lines=1&&&&&\\
   6562.8&    10.20&   0.71&    599&    7.18& Jeffery \\
\hline
He\,{\sc i}&&&&&\\
%No. of lines=3&&&&&\\
% 3888.65&   19.73&   -0.70&    368&   11.59& Jeffery\\
% 3926.53&   21.13&  -1.65&    185&   11.75& Jeffery\\
% 3964.73&   20.61&   -1.30&    210&   11.39& Jeffery\\
% 4120.81&   20.87&  -1.52&    200&   11.58& Jeffery\\
% 4471.48&   20.87&   0.04&    375&   11.33& Jeffery\\
% 4713.14&   20.87&  -1.07&    267&   11.73& Jeffery\\
  5047.74&   21.22&  $-$1.60&   Synth$^c$&11.54   &Jeffery \\
  5875.62&   20.96&   0.74&   Synth&11.54   & Jeffery \\
  6678.15&   21.22&   0.33&   Synth&11.54   & Jeffery \\
\hline
C\,{\sc i}&&&&&\\
%No. of lines=30&&&&&\\
5052.18 & 7.69 & $-$1.30 &252 &   9.26 &  Wiese \\
5380.34 & 7.69 & $-$1.62 &193 &   9.13 &  Wiese \\
5540.76 & 8.64 & $-$2.40 & 81 &   9.59 &  Wiese \\
5817.70 & 8.85 & $-$2.86 & 19 &   9.41 &  Wiese \\
%5850.24 & 8.77 & $-$2.70 & 36 &   9.51 &  Wiese \\
%5963.99 & 8.65 & $-$2.70 & 81 &   9.90 &  Wiese \\
6001.12 & 8.64 & $-$2.05 &164 &   9.83 &  Wiese \\
6002.98 & 8.65 & $-$2.16 & 83 &   9.37 &  Wiese \\
6010.68 & 8.64 & $-$1.96 &166 &   9.75 &  Wiese \\
6014.83 & 8.64 & $-$1.59 &215 &   9.72 &  Wiese \\
6078.39 & 8.85 & $-$2.27 & 74 &   9.51 &  Wiese \\
6113.15 & 8.85 & $-$2.63 & 66 &   9.79 &  Wiese \\
6115.84 & 8.85 & $-$2.52 & 85 &   9.85 &  Wiese \\
6120.81 & 8.85 & $-$2.40 & 84 &   9.72 &  Wiese \\
%6337.18 & 8.77 & $-$2.30 & 59 &   9.36 &  Wiese \\
%6342.32 & 8.77 & $-$2.10 & 74 &   9.29 &  Wiese \\
6413.55 & 8.77 & $-$2.00 &153 &   9.78 &  Wiese \\
6568.71 & 9.00 & $-$2.16 & 93 &   9.62 &  Wiese \\
6587.61 & 8.54 & $-$1.00 &268 &   9.43 &  Wiese \\
6591.46 & 8.85 & $-$2.40 & 54 &   9.47 &  Wiese \\
6595.24 & 8.85 & $-$2.40 & 35 &   9.24 &  Wiese \\
6611.35 & 8.85 & $-$1.85 &123 &   9.47 &  Wiese \\
6671.85 & 8.85 & $-$1.66 &166 &   9.55 &  Wiese \\
6688.79 & 8.85 & $-$2.16 & 95 &   9.56 &  Wiese \\
6711.32 & 8.54 & $-$2.70 & 50 &   9.57 &  Wiese \\
7108.93 & 8.64 & $-$1.59 &182 &   9.48 &  Wiese \\
7111.47 & 8.64 & $-$1.09 &252 &   9.46 &  Wiese \\
7113.18 & 8.65 & $-$0.77 &288 &   9.38 &  Wiese \\
7476.18 & 8.77 & $-$1.57 &193 &   9.60 &  Wiese \\
7685.19 & 8.77 & $-$1.52 &182 &   9.48 &  Wiese \\
7832.64 & 8.85 & $-$1.82 &123 &   9.43 &  Wiese \\
7860.88 & 8.85 & $-$1.15 &241 &   9.48 &  Wiese \\
%\end{tabular}
%}
%\end{center}
%\newpage
%\begin{center}
%{\bf Table 4.6} (continued)\\
%\vspace{0.5cm}
%{\it Lines used to derive elemental abundances for FQ Aqr}\\
%\vspace{0.5cm}
%\noindent
%{\scriptsize
%\begin{tabular}{lccccl}
%\hline \hline
%$\lambda$(\AA~)& $\chi$(eV)& log gf& W(m\AA~)& Abundance$^{a}$& Reference$^{b}$\\
%\hline
\hline
Mean:&&&&9.53$\pm$0.18&\\
\hline
C\,{\sc ii}&&&&&\\
%No. of lines=2&&&&&\\
5537.61 & 19.50 & $-$1.79 & 42 & 9.57 &  Wiese \\
5891.60 & 18.05 & $-$0.44 &204 & 9.34 &  Wiese\\
6578.05 & 14.45 & $-$0.03 &848 & 9.36 &  Wiese\\
6582.88 & 14.45 & $-$0.33 &728 & 9.48 &  Wiese\\
\hline
Mean:&&&& 9.44$\pm$0.11&\\
\hline
\end{tabular}
%\medskip
%$^a$Normalized such that log$\Sigma \mu_{i} \epsilon(i)$ = 12.15\\
%$^b$Sources of $gf$-values
\end{table}
\begin{table}
\contcaption{}
\begin{tabular}{lrrrrl}
\hline
Ion&&&{\em W$_{\lambda}$}&&\\
$\lambda$ (\AA)& $\chi$ (eV)& log $gf$&  (m\AA)& log $\epsilon$$^a$
& Ref.$^b$\\
\hline
N\,{\sc i}&&&&&\\
%No. of lines=5&&&&&\\
6008.46 & 11.60 & $-$1.11 & 75  & 8.30  & Wiese \\
6644.96 & 11.76 & $-$0.86 & 63  & 8.03  & Wiese \\
7406.22 & 12.01 & $-$0.74 & 68  & 8.08  & Wiese \\
7423.64 & 10.33 & $-$0.71 &237  & 8.32  & Wiese \\
7442.30 & 10.33 & $-$0.38 &302  & 8.40  & Wiese \\
\hline
Mean:&&&&8.23$\pm$0.16&\\
\hline
N\,{\sc ii}&&&&&\\
%No. of lines=2&&&&&\\
5045.10 & 18.48 & $-$0.41 &  79  &  8.37 & Wiese  \\  
5666.63 & 18.47 & $-$0.05 &  86  &  8.14 & Wiese  \\
5679.56 & 18.48 &  0.25 & 118  &  8.17 & Wiese  \\
5686.21 & 18.47 & $-$0.55 &  38  &  8.05 & Wiese  \\
5710.77 & 18.48 & $-$0.52 &  79  &  8.55 & Wiese  \\
\hline
Mean:&&&&8.26$\pm$0.20&\\
\hline
O\,{\sc i}&&&&&\\
%No. of lines=8&&&&&\\
4968.79 & 10.74 & $-$1.26  &  233   &  9.35 &   Wiese \\
5330.74 & 10.74 & $-$0.88  &  302   &  9.35 &   Wiese \\
5436.86 & 10.74 & $-$1.40  &  208   &  9.25 &   Wiese \\
5555.00 & 10.99 & $-$1.80  &  149   &  9.35 &   Wiese \\
6158.19 & 10.74 & $-$0.30  &  401   &  9.41 &   Wiese \\
7473.24 & 14.12 & $-$0.37  &   86   &  9.00 &   Wiese \\
\hline
Mean:&&&&9.29$\pm$0.15&\\
\hline
Ne\,{\sc i}&&&&&\\
%No. of lines=14&&&&&\\
%5748.30 & 18.56 & $-$1.02  &  119   &  8.90  &  NIST \\
%5764.42 & 18.56 & $-$0.31  &  200   &  8.98  &  NIST \\
5852.49 & 16.85 & $-$0.46  &  302   &  9.21  &  NIST \\
5881.90 & 16.62 & $-$0.75  &  280   &  9.20  &  NIST \\
%5944.83 & 16.62 & $-$0.52  &  319   &  9.30  &  NIST \\
%5965.47 & 18.73 & $-$1.13  &  133   &  9.25  &  NIST \\
6030.00 & 16.67 & $-$1.04  &  231   &  9.08  &  NIST \\
6074.34 & 16.67 & $-$0.48  &  338   &  9.43  &  NIST \\
6143.06 & 16.62 & 0.10   &  505   &  9.93  &  NIST \\
6163.59 & 16.72 & $-$0.60  &  315   &  9.39  &  NIST \\
6217.28 & 16.62 & $-$0.96  &  247   &  9.11  &  NIST \\
6266.50 & 16.72 & $-$0.36  &  404   &  9.78  &  NIST \\
6334.43 & 16.62 & $-$0.32  &  419   &  9.76  &  NIST \\
6382.99 & 16.67 & $-$0.23  &  446   &  9.85  &  NIST \\
6402.25 & 16.62 & 0.35   &  584   &  9.74  &  NIST \\
6506.53 & 16.67 &$-$0.02   &  485   &  9.80  &  NIST \\
6598.95 & 16.85 & $-$0.34  &  357   &  9.55  &  NIST \\
7032.41 & 16.62 & $-$0.25  &  449   &  9.84  &  NIST \\
%\end{tabular}
%}
%\end{center}
% 
%\newpage
% 
%\begin{center}
%{\bf Table} (continued)\\
%\vspace{0.5cm}
%{\it Lines used to derive elemental abundances for FQ Aqr}\\
%\vspace{0.5cm}
%\noindent
%{\scriptsize
%\begin{tabular}{lccccl}
%\hline
%$\lambda$(\AA~)& $\chi$(eV)& log gf& W(m\AA~)& Abundance$^{a}$& Reference$^{b}$\\
%\hline
\hline
Mean:&&&&9.55$\pm$0.30&\\
\hline
Na\,{\sc i}&&&&&\\
%No. of lines=2&&&&&\\
5682.63 & 2.10 & $-$0.70  &  32  &  6.40  &  NIST \\
6160.75 & 2.10 & $-$1.23  &  29  &  6.89  &  NIST \\
8194.82 & 2.10 &  0.51  & 257  &  6.78  &  NIST \\
\hline
Mean:&&&&6.69$\pm$0.26&\\
\hline
Mg\,{\sc ii}&&&&&\\
%No. of lines=6&&&&&\\
7877.05 & 10.00 & 0.39  & 246  & 5.89  &  NIST \\  
7896.37 & 10.00 & 0.65  & 310  & 6.04  &  NIST \\
\hline
Mean:&&&&5.97$\pm$0.11&\\
\hline
\end{tabular}
%\medskip
%$^a$Normalized such that log$\Sigma \mu_{i} \epsilon(i)$ = 12.15\\
%$^b$Sources of $gf$-values
\end{table}
\begin{table}
\contcaption{}
\begin{tabular}{lrrrrl}
\hline
Ion&&&{\em W$_{\lambda}$}&&\\
$\lambda$ (\AA)& $\chi$ (eV)& log $gf$& (m\AA)& log $\epsilon$$^a$
& Ref.$^b$\\
\hline
Al\,{\sc ii}&&&&&\\
%No. of lines=4&&&&&\\
5593.23 & 13.26 & 0.41 &   95  &  5.19 & NIST  \\ 
6226.13 & 13.07 & 0.05 &   86  &  5.39 & NIST  \\
6231.72 & 13.07 & 0.40 &  171  &  5.71 & NIST  \\
6823.40 & 13.07 &$-$0.14 &  102  &  5.72 & NIST  \\
6837.08 & 13.08 & 0.08 &  122  &  5.66 & NIST  \\
7042.05 & 11.32 & 0.35 &  393  &  6.52 & NIST  \\
%7063.64 & 11.32 &$-$0.35 &  289  &  6.43:& NIST  \\
\hline
Mean:&&&&5.70$\pm$0.45&\\
\hline
Al\,{\sc iii}&&&&&\\
5696.47 & 15.64&  0.24  &  85  &  6.04 &  NIST  \\  
5722.65 & 15.64& $-$0.07  &  57  &  6.03 &  NIST  \\
\hline
Mean:&&&&6.04$\pm$0.01&\\
\hline
Si\,{\sc ii}&&&&&\\
%No. of lines=6&&&&&\\
5055.98 & 10.07 &  0.44 & 490  &  6.86 &  NIST \\   
5957.56 & 10.07 & $-$0.35 & 304  &  6.72 &  NIST \\
5978.93 & 10.07 & $-$0.06 & 329  &  6.59 &  NIST \\
6829.83 & 12.88 & $-$0.27 &  69  &  6.31 &  NIST \\
\hline
Mean:&&&&6.62$\pm$0.23&\\
\hline
P\,{\sc ii}&&&&&\\
%No. of lines=2&&&&&\\
5499.69 & 10.80 & $-$0.30  &  38  & 4.69 & NIST \\  
6024.13 & 10.76 & 0.14   & 111  & 4.94 & NIST \\
6034.04 & 10.74 & $-$0.22  &  71  & 4.95 & NIST \\
6165.57 & 10.80 & $-$0.34  &  45  & 4.83 & NIST \\
\hline
Mean:&&&&4.85$\pm$0.12&\\
\hline
S\,{\sc ii}&&&&&\\
%No. of lines=3&&&&&\\
5009.52  &13.62 & $-$0.28  & 169 & 6.81  & NIST \\   
5014.00  &14.07 & 0.10   & 173 & 6.70  & NIST \\
5103.29  &13.67 & $-$0.11  & 108 & 6.11  & NIST \\
5320.70  &15.07 & 0.49   & 151 & 6.56  & NIST \\
5428.64  &13.58 & $-$0.13  & 176 & 6.72  & NIST \\
%5432.74 & 13.62&  0.26  &  292&  7.31 &  NIST \\
%5453.79 & 13.67&  0.48  &  326&  7.34 &  NIST \\
5473.60  &13.58 & $-$0.18  & 154 & 6.57  & NIST \\
%5509.65 & 13.62&  $-$0.14 &  207&  7.03 &  NIST \\
5555.97  &13.62 & $-$0.99  &  45 & 6.32  & NIST \\
5606.09  &13.73 & 0.31   & 194 & 6.54  & NIST \\
5616.61  &13.66 & $-$0.64  &  88 & 6.46  & NIST \\
5659.96  &13.68 & $-$0.05  & 170 & 6.62  & NIST \\
5664.76  &13.66 & $-$0.25  & 111 & 6.29  & NIST \\
\hline
Mean:&&&&6.52$\pm$0.21&\\
\hline
%Sc\,{\sc ii}&&&&&\\
%No. of lines=7&&&&&\\
%5684.20 & 1.51 & $-$1.25 &23.0 & 3.58 &  Luck \\
%\hline
Ti\,{\sc ii}&&&&&\\
5072.28 & 3.12 & $-$0.75 &   36.0  &  4.18 & NIST \\  
5188.68 & 1.58 & $-$1.21 &  102.0  &  4.39 & NIST \\
\hline
Mean:&&&&4.29$\pm$0.15&\\
\hline
Cr\,{\sc ii}&&&&&\\
5246.77 & 3.71 & $-$2.48 &  18.0 &  4.42 &  NIST \\   
%5274.96 & 4.07 & $-$1.29 & 170.0 &  4.76 &  Kurucz \\
5313.56 & 4.07 & $-$1.65 &  90.0 &  4.62 &  NIST \\
%5334.87 & 4.07 & $-$1.56 & 135.0 &  4.82 &  Kurucz \\
5420.92 & 3.76 & $-$2.58 &  20.0 &  4.37 &  NIST \\
%5478.37 & 4.18 & $-$1.91 &  68.0 &  4.76 &  Kurucz \\
\hline
Mean:&&&&4.47$\pm$0.13&\\
\hline
\end{tabular}
%\medskip
%$^a$Normalized such that log$\Sigma \mu_{i} \epsilon(i)$ = 12.15\\
%$^b$Sources of $gf$-values
\end{table}
\begin{table}
\contcaption{}
\begin{tabular}{lrrrrl}
\hline
Ion&&&{\em W$_{\lambda}$}&&\\
$\lambda$ (\AA)& $\chi$ (eV)& log $gf$& (m\AA)& log $\epsilon$$^a$
& Ref.$^b$\\
\hline
Mn\,{\sc ii}&&&&&\\
%No. of lines=3&&&&&\\
5297.06&  9.86 & 0.87  &  47.0 &  4.67  & Kurucz \\   
6122.45& 10.18 & 0.95  &  57.0 &  4.85  & Kurucz \\
7415.81&  3.71 &$-$2.20  &  40.0 &  4.37  & Kurucz \\
7432.30&  3.71& $-$2.50  &  21.0 &  4.37  & Kurucz \\
\hline
Mean:&&&&4.57$\pm$0.24&\\
\hline
Fe\,{\sc ii}&&&&&\\
%No. of lines=59&&&&&\\
4993.36 & 2.81 & $-$3.65   &  113.0  &  6.32  & NIST  \\  
%5081.90 & 10.38&  $-$0.59  &   87.0  &  7.04  & NIST  \\
5169.03 & 2.89 & $-$0.87   &  487.0  &  6.28  & NIST  \\
5197.58 & 3.23 & $-$2.10   &  235.0  &  5.77  & NIST  \\
5234.63 & 3.22 & $-$2.05   &  278.0  &  6.02  & NIST  \\
5247.95 & 10.53&  0.63   &  121.0  &  6.14  & NIST  \\
5254.93 & 3.23 & $-$3.00   &   97.0  &  5.77  & Kurucz \\
5272.40 & 5.96 & $-$2.05   &   82.0  &  6.18  & NIST  \\
5276.00 & 3.20 & $-$1.96   &  242.0  &  5.64  & NIST  \\
5284.11 & 2.89 & $-$3.19   &  110.0  &  5.86  & NIST  \\
%5325.55 & 3.22 & $-$2.52   &  105.0  &  5.34  & NIST  \\
5362.87 & 3.20 & $-$2.70   &  192.0  &  6.03  & Kurucz  \\
5525.13 & 3.27 & $-$4.61   &   24.0  &  6.64  & Kurucz  \\
5534.85 & 3.25 & $-$2.93   &  195.0  &  6.29  & NIST  \\
5591.37 & 3.27 & $-$4.69   &   12.0  &  6.40  & Kurucz  \\
5627.50 & 3.39 & $-$4.36   &   33.0  &  6.62  & NIST  \\
5952.51 & 5.96 & $-$2.05   &   61.0  &  5.97  & Kurucz  \\
5961.71 & 10.68&  0.69   &  162.0  &  6.41  & Kurucz  \\
5991.38 & 3.15 & $-$3.74   &   82.0  &  6.32  & NIST  \\
6147.74 & 3.89 & $-$2.70   &  123.0  &  5.94  & Kurucz  \\
6149.26 & 3.89 & $-$2.92   &  118.0  &  6.13  & NIST  \\
6175.15 & 6.22 & $-$2.00   &   64.0  &  6.08  & Kurucz  \\
6179.38 & 5.57 & $-$2.81   &   42.0  &  6.32  & NIST  \\
6247.56 & 3.89 & $-$2.52   &  199.0  &  6.21  & NIST  \\
6331.96 & 6.22 & $-$1.96   &   48.0  &  5.87  & Kurucz  \\
6446.41 & 6.22 & $-$2.16   &   38.0  &  5.97  & NIST  \\
7462.41 & 3.89 & $-$2.70   &   90.0  &  5.68  & Kurucz \\
\hline
Mean:&&&&6.11$\pm$0.27&\\
\hline
Y\,{\sc ii}&&&&&\\
%No. of lines=3&&&&&\\
5205.73 & 1.03 & $-$0.34  &   28.0  &  2.85  & Luck \\
\hline
Zr\,{\sc ii}&&&&&\\
4962.29 & 0.97 & $-$1.69  &    6.0  &  3.42 & Thevenin \\  
5350.09 & 1.83 & $-$0.93  &   19.0  &  3.60 & Thevenin \\
\hline
Mean:&&&&3.51$\pm$0.13&\\
\hline
Ba\,{\sc ii}&&&&&\\
6496.90 & 0.60 & $-$0.37  &   14.0  &  2.49 & Luck \\
\hline
\end{tabular}
\medskip
$^a$Normalized such that log$\Sigma \mu_{i} \epsilon(i)$ = 12.15\\
\medskip
$^b$Sources of $gf$-values\\
\medskip
$^c$Spectrum synthesis
\end{table}
\begin{table}
%\contcaption{}
References:
\medskip
\begin{tabular}{lll}
Jeffery && Jeffery, Woolf, \& Pollacco (2001)\\
Kurucz && Kurucz's database\\
Luck && Compilations by R. E. Luck \\
NIST&& NIST database\\
Thevenin && Th\'{e}venin (1989, 1990)\\
Wiese    && Wiese, Fuhr, \& Deters (1996)\\
\hline
\end{tabular}
\end{table}

\end{document}




\subsection{The [C\,{\sc i}] lines}

%An accurate wavelength and $gf$-value are required
% for each line included in spectrum synthesis. 
The standard reference for the C\,{\sc i} spectrum (Moore 1993)
gives a  predicted wavelength of
8727.126 \AA\ for the [C\,{\sc i}] $2p^2$ $^1$D$_2$ -- $2p^2$ $^1$S$_0$
transition and 
 9850.264 \AA\ for the [C\,{\sc i}]
$2p^2$ $^3$P$_2$ -- $2p^2$ $^1$D$_2$  transition. 
%as computed from energy levels derived from combinations of
%ultraviolet lines between
% levels of the ground configuration and a common upper excited
%level.
% A plausible uncertainty in the predicted position of the line is $\pm$0.02 cm$^{-1}$
%or $\pm$0.019\AA.
 Wavelengths of the forbidden carbon lines have not been measured directly
from laboratory sources.


The predicted wavelength of the 8727 \AA\ line
 is confirmed by the wavelength of the absorption line in the solar
spectrum (Allende Prieto, Lambert, \& Asplund 2002) and the emission line in planetary
nebulae (Liu et al. 1995).
For the 8727 \AA\ line,
we adopt $\log gf$ = $-$8.14 (Galavis, Mendoza \& Zeippen 1997).
 The lower excitation potential is 1.264 eV.
 There is a blending Fe\,{\sc i} line for which we adopt the
$\log gf$ = $-$4.4 (Allende Prieto, Lambert, \& Asplund 2002), but
this line is not a significant
contributor to the spectrum of a RCB star.

For the 9850 \AA\ line, we adopt the predicted wavelength.
 There is supporting evidence
from our spectra for this wavelength.
 The 2002 July spectra of
RY Sgr show the 9850 \AA\ and 8727 \AA\ [C\,{\sc i}] lines in emission
(see Figures 1 and 2). Adopting 8727.126 \AA\ as the rest
wavelength for the latter emission and assuming the two lines
have the same velocity, we find the astrophysical wavelength
of  the 9850 \AA\ emission is close to the predicted wavelength.
Liu et al. (1995) report, however,
 a rest wavelength of 9850.36 $\pm$ 0.01\AA\ from detections
of the line in four planetary nebulae. 
The 0.1\AA\ or 3km s$^{-1}$
 difference between the wavelength predicted from
energy levels and that observed from planetary nebulae does not
seriously impact derivation of the carbon abundance from the
line in the RCB spectra.
The transition probability for the 9850 \AA\ line is taken from
Galavis, Mendoza \& Zeippen (1997) who give A = 2.23 $\times 10^{-4}$ s$^{-1}$ or
$\log gf = -10.79$ with an uncertainty of about 0.04 dex.   
The lower excitation potential is a mere 0.005 eV.

Two additional lines complete multiplet 1F.
The stronger line with a $gf$-value about 3 times smaller than for the 9850 \AA\ line is
at 9824.30 $\pm$ 0.05 \AA\ (Liu et al. 1995) and
irretrievably blended in the cooler RCB (GU Sgr, V482 Cyg, SU Tau, and R CrB) spectra.
%, being sandwiched between N\,{\sc i} and O\,{\sc i} lines
In the spectra of VZ Sgr, UV Cas, and XX Cam, 
the 9824.30 \AA\ line is absent, but this is as expected from the S/N and
the estimated strength of the observed 9850 \AA\ line.
The third [C\,{\sc i}] line at 9808.3 \AA\ is expected to be 4000
times weaker than the 9850 \AA\ line, and  undetectable.
(The multiplet 2F $2p^2$ $^3$P$_{1}$ - $2p^2$ $^1$S$_{0}$ transition
at 4621.57 \AA\ with $\log gf$ = $-$11.12 is almost certainly blended as the
blue spectral region of a RCB is rich in strong lines.)



\subsection{Overview of the Spectra}

The region around 9850 \AA\ is shown in Figure 1 where the
stars are ordered by decreasing temperature from top to bottom.
The spectrum of the  normal yellow supergiant $\gamma$ Cyg is shown at the
bottom. In the spectra of the hottest stars, the strongest
lines are from C\,{\sc i} and N\,{\sc i} transitions with weaker
lines of Si\,{\sc i}, Fe\,{\sc i}, and Fe\,{\sc ii}. The [C\,{\sc i}]
line falls in the red wing of a Fe\,{\sc ii} line. Spectra of the
coolest stars GU Sgr and V482 Cyg show additional lines which we identify as 
high rotational lines of the 1-0 band of the CN molecule's Red System
(Davis \& Phillips 1963). These CN lines are prominent in the
spectra of GU Sgr and V482 Cyg and traceable in SU Tau, R CrB and
RY Sgr. Two CN lines bracket the [C\,{\sc i}] line.

Inspection of Figure 1 shows several additional features of
interest. As anticipated, the
strength of a given C\,{\sc i} line is unchanged along the
temperature sequence from XX Cam at the high temperature
end to GU Sgr at the low temperature end. Also anticipated
is the very large width of all absorption lines in the
RCB spectra. Compare the widths with those of lines in $\gamma$ Cyg,
which itself is commonly referred to as showing broad lines.
A surprise  is
the appearance of emission in the RY Sgr spectrum near the
wavelength of the [C\,{\sc i}] line. The illustrated spectrum is from
31 July 2002 in which the 8727 \AA\ line is also in
emission. A spectrum from 22 June 1997 also shows emission
but to the red of the wavelength of the [C\,{\sc i}] line.
The possibility of emission contaminating the
absorption line is an unfortunate complication.

\begin{figure}
\epsfxsize=8truecm
\epsffile{fig1.ps}
\caption{The spectra of the RCB stars in the region of [C\,{\sc i}] 9850.264\AA\
line, which is indicated by a vertical solid line, are plotted with their effective temperatures 
increasing from bottom to
top. The spectrum of $\gamma$ Cyg is also plotted. The positions of
a number of other atomic and some CN lines are also
marked and the vertical dashed lines indicate the Fe\,{\sc ii} features. The relative
intensity scale for V482 Cyg, SU Tau, VZ Sgr, UV Cas, and RY Sgr is same as for RCB.}
\end{figure}

\begin{figure}
\epsfxsize=8truecm
\epsffile{fig2.ps}
\caption{The spectra of RCB stars in the region of [C\,{\sc i}] 8727.126\AA\
line, which is indicated by a vertical solid line, are plotted with their effective temperatures
increasing from bottom to top.
The spectrum of $\gamma$ Cyg is also plotted. The positions of the key lines
are also marked.}
\end{figure}

%\subsection{The  8727 \AA\ Region}

%The remarks made about the 9850 \AA\ region apply in large measure to
%the region around the 8727 \AA\ [C\,{\sc i}] line, which  is
%stronger than the 9850 \AA\ line 
% (Figure 2).
 The [C\,{\sc i}] 8727 \AA\ line (Figure 2)
is slightly blended, as expected, with a Si\,{\sc i} line. Emission
is very pronounced in the spectrum of RY Sgr. The central line depths
of the [C\,{\sc i}] run from about 16 \% for SU Tau to
36 \% for RY Sgr with no obvious systematic trend with the
effective temperature.

\section{Spectrum Synthesis - Procedure}

The fact that the [C\,{\sc i}] lines are blended with other
lines requires that
the carbon abundance be extracted  by
spectrum synthesis. 
%This technique requires a
%reliable model atmosphere and atomic data for the offending
%blends. 
In addition to synthesizing the [C\,{\sc i}] line and its
blends, we use lines of C\,{\sc i}, N\,{\sc i}, Si\,{\sc i},
Fe\,{\sc i} and Fe\,{\sc ii} to obtain abundances and compare with
values given by Asplund et al. (2000). 
%from their analysis of lines
%at shorter wavelengths.

Synthetic spectra are generated using the H-deficient model
atmospheres computed by Asplund et al. (1997) and the Uppsala
spectrum synthesis code BSYNRUN. Models corresponding to a C/He ratio of
1 \% by number were chosen. This composition
is equivalent to a carbon abundance $\log\epsilon$(C) = 9.54
on the scale $\Sigma\mu_i\epsilon_i = 12.15$, where
$\mu_i$ is the mean atomic weight of element $i$. Stellar parameters
derived by Asplund et al. (2000) 
%in their comprehensive abundance analysis 
were adopted as our initial values. Synthetic
spectra  were convolved
with a Gaussian profile to account for the combined broadening of
stellar lines by the atmospheric macroturbulent velocity, rotational
broadening, and the instrumental profile.

\subsection{The 9850 \AA\ Region}


The 9850 \AA\ [C\,{\sc i}] line
may be blended with Fe\,{\sc ii}  and CN lines. The Fe\,{\sc ii}
line  at 9849.74 \AA\ has a lower excitation potential of 6.729 eV.
To obtain the $gf$-value of the line, we use the spectrum of $\gamma$
Cyg where the Fe\,{\sc ii} line is present. With a model
atmosphere constructed for the stellar parameters and iron abundance
given by Luck \& Lambert (1981), we find $\log gf$ = $-$2.60. The
basic data for the CN 1-0 lines were taken from Kurucz's
line list (http://kurucz.harvard.edu) and
checked against the SCAN tape (J{\o}rgensen \& Larsson 1990)
and a recipe given by Bakker \& Lambert (1998). A dissociation
energy of  7.65 eV was adopted for the molecule (Bauschlicher, Langhoff \&  Taylor
1988). Unblended CN lines away from the 9850 \AA\ feature
were fitted by adjusting the N abundance, and this abundance
used in computing the CN lines blended with the [C\,{\sc i}]
line.
An estimate of the N abundance was obtained not only from the
CN lines but also from N\,{\sc i} lines near the [C\,{\sc i}] lines.
The $gf$-values for the N\,{\sc i} lines were taken from 
Wiese, Fuhr \& Deters (1996) compilation. 

%For a direct comparison of the [C\,{\sc i}] line with the C\,{\sc i}
%lines, we include two C\,{\sc i} lines in the synthetic spectrum
%around 9850 \AA. 
The C\,{\sc i} lines at 9852.27 \AA,  a
blend of two lines, and 9859.15 \AA\
(Figure 1) have also been used for estimating the C abundance. 
Their $gf$-values are determined by inversion of their
equivalent widths  in the $\gamma$ Cyg spectrum using 
%the model previously mentioned and 
the C abundance derived
from a set of weak C\,{\sc i} lines at shorter wavelengths. 
%and  the
%8727 \AA\ [C\,{\sc i}] line which is clearly seen in
%Figure 2. The 9850 \AA\ [C\,{\sc i}] line is not detected in
%our spectrum of $\gamma$ Cyg.


\subsection{The 8727 \AA\ Region}

The red wing of the  [C\,{\sc i}] line is blended with a Si\,{\sc i}
line (Figure 2). The $gf$-value of the latter line is
derived from its equivalent width in the solar spectrum. The
$gf$-value of the N\,{\sc i} line to the red of the Si\,{\sc i}
line is taken from Wiese, Fuhr \& Deters (1996). Spectrum synthesis of
this region was confined to a 6 \AA\ window around the
[C\,{\sc i}] line.

%\begin{table*}
%\centering
%\begin{minipage}{150mm}
%\caption{Photospheric C abundances of C, N, and Fe, for
%the analysed R\,CrBs. (Uncertain abundances are marked with :)}
%\begin{tabular}{lcccccccc} \hline
%\multicolumn{1}{c}{}& \multicolumn{6}{c}{log$\epsilon(el)$} \\ \cline{2-9}
%\multicolumn{1}{c}{}&\multicolumn{1}{c}{GU~Sgr}&\multicolumn{1}{c}{V482~Cyg}
%&\multicolumn{1}{c}{SU~Tau}&\multicolumn{1}{c}{R~CrB}&
%\multicolumn{1}{c}{VZ~Sgr}& \multicolumn{1}{c}{UV~Cas}&\multicolumn{1}{c}{RY~Sgr}&\multicolumn{1}{c}{XX~Cam}\\ \hline
%log$\epsilon$(C)&&&&&&&&\\
%$[\rm C\,{\sc i}]$ 8727.126\AA&...&...&8.2$\pm$0.3&8.4$\pm$0.3&...&8.3$\pm$0.3&...&8.4$\pm$0.3\\
%$[\rm C\,{\sc i}]$ 9850.36\AA&9.0:&$<$9.2:&8.9$\pm$0.2&9.1$\pm$0.2&$<$8.7&$<$8.6$\pm$0.2&...&$<$8.8$\pm$0.2\\
%log$\epsilon$(C)&&&&&&&&\\
%C\,{\sc i} 9852\AA&9.4$\pm$0.1&9.6$\pm$0.1&9.0$\pm$0.1&9.2$\pm$0.1&9.1$\pm$0.1&9.0$\pm$0.1&9.2$\pm$0.1&8.9$\pm$0.1\\
%C\,{\sc i} 9859.15\AA&9.2$\pm$0.1&9.6$\pm$0.1&9.0$\pm$0.1&9.2$\pm$0.1&9.1$\pm$0.1&9.0$\pm$0.1&9.3$\pm$0.1&8.8$\pm$0.1\\
%Asplund et al. (2000)&8.8$\pm$0.3&8.9$\pm$0.2&8.8$\pm$0.2&9.2$\pm$0.2&9.1$\pm$0.3&9.2$\pm$0.4&8.9$\pm$0.2&9.0$\pm$0.4\\
%log$\epsilon$(N)&&&&&&&&\\
%N\,{\sc i} 8728.90\AA&...&...&7.9$\pm$0.2&8.1$\pm$0.2&...&8.5$\pm$0.2&...&8.7$\pm$0.2\\
%N\,{\sc i} 9863.34\AA&9.1$\pm$0.2&9.0$\pm$0.2&7.8$\pm$0.2&8.2$\pm$0.2&7.6$\pm$0.2&8.6$\pm$0.2&8.6$\pm$0.2&8.9$\pm$0.2\\
%N\,{\sc i} 9872.15\AA&9.1:$\pm$0.2&9.0:$\pm$0.2&...&8.2$\pm$0.2&7.6$\pm$0.2&8.6$\pm$0.2&8.1:$\pm$0.2&9.0$\pm$0.2\\
%Asplund et al. (2000)&8.7$\pm$0.5&8.8$\pm$0.5&8.5$\pm$0.1&8.4$\pm$0.2&7.6$\pm$0.2&8.5$\pm$0.3&8.5$\pm$0.2&8.9$\pm$0.3\\
%log$\epsilon$(Si)&&&&&&&&\\
%Si\,{\sc i} 8728.01\AA&...&...&6.3$\pm$0.2&6.9$\pm$0.2&...&7.1$\pm$0.2&...&6.8$\pm$0.2\\
%Asplund et al. (2000)&7.2$\pm$0.2&7.2$\pm$0.1&6.7$\pm$0.1&7.2$\pm$0.2&7.3$\pm$0.2&7.4$\pm$0.2&7.3$\pm$0.2&7.1$\pm$0.3\\
%log$\epsilon$(Fe)&&&&&&&&\\
%Fe\,{\sc i} 9861.74\AA&6.8:$\pm$0.2&6.7$\pm$0.2&6.1$\pm$0.2&6.3$\pm$0.2&...&6.8$\pm$0.2&7.0$\pm$0.2&6.6$\pm$0.2\\
%Fe\,{\sc ii} 9849.74\AA&7.1:$\pm$0.2&...&6.1$\pm$0.2&6.5$\pm$0.2&...&6.8$\pm$0.2&...&6.8$\pm$0.2\\
%Asplund et al. (2000)&6.3$\pm$0.3&6.7$\pm$0.3&6.1$\pm$0.3&6.5$\pm$0.1&5.8$\pm$0.2&6.9$\pm$0.2&6.7$\pm$0.2&6.8$\pm$0.2\\
%\hline
%\end{tabular}
%\end{minipage}
%\end{table*}


%\begin{table}
%\centering
%\begin{minipage}{65mm}
%\caption{Abundance changes due to the uncertainty in the adopted
%$T_{\rm eff}$ and log$g$ of the analysed R\,CrBs. The uncertainty ($\Delta$)
%in $T_{\rm eff}$ and log$g$ are shown first. The corresponding abundance changes $\Delta$(log$\epsilon$) in the derived 
%C and Fe abundances (using 9850\AA\ observed profile) are given for two adopted model atmospheres.The uncertainty
%in the stellar parameters and the corresponding abundance changes are defined by $\Delta$ = perturbed $-$ adopted.}
%\begin{tabular}{lccc}
%\hline
%\multicolumn{1}{c}{} & \multicolumn{1}{c}{$\Delta$$T_{\rm eff}$} & 
%\multicolumn{1}{c}{$\Delta$log$g$} & \multicolumn{1}{c}{$\Delta$$\xi$}\\
%\multicolumn{1}{c}{} & \multicolumn{1}{c}{$\pm$250} &\multicolumn{1}{c}{$\pm$0.5} &\multicolumn{1}{c}{$\pm$1.0}\\ 
%\multicolumn{1}{c}{} & \multicolumn{1}{c}{(K)} &\multicolumn{1}{c}{(cgs units)} &\multicolumn{1}{c}{(km s$^{-1}$)}\\
%\hline
%\multicolumn{1}{c}{} & \multicolumn{1}{c}{$\Delta$(log$\epsilon$)} &
%\multicolumn{1}{c}{$\Delta$(log$\epsilon$)} &\multicolumn{1}{c}{$\Delta$(log$\epsilon$)}\\
%\hline
%log$\epsilon$(C)$^a$   & $\pm$0.2  &    0.0& 0.0\\
%log$\epsilon$(C)$^b$   & $\pm$0.2  &    0.0& 0.0\\
%log$\epsilon$(Fe)$^a$  & $\mp$0.1  &   $\pm$0.2& 0.0\\
%log$\epsilon$(Fe)$^b$  & $\mp$0.1  &   $\pm$0.2& 0.0\\
%\hline
%\end{tabular}
%$^a$Adopted $T_{\rm eff}$=6750K, log$g$=0.50(cgs units), and $\xi$=7km s$^{-1}$\\
%$^b$Adopted $T_{\rm eff}$=7250K, log$g$=0.75(cgs units), and $\xi$=9km s$^{-1}$\\
%\end{minipage}
%\end{table}

\section{Spectrum Synthesis - Results}

%In the Introduction, we noted the `carbon problem'
%(Asplund et al. 2000). The C\,{\sc i} lines in RCB spectra are
%much weaker than predicted. Quantitatively, the effect corresponds to
%a factor of four reduction in the product of a line's $gf$-value and
%the C abundance. 
The [C\,{\sc i}] lines were sought in order to
shed new light on `carbon problem'. Asplund et al.
suggested that the 9850 \AA\ [C\,{\sc i}]
line was not subject to the carbon problem. Their analysis was
flawed in that the contributions of the blending Fe\,{\sc ii}
and CN line were not considered.

In addition to fitting the [C\,{\sc i}] lines, we analysed other lines
in order to check the abundances given by Asplund
et al. (2000). To within the uncertainties of the analyses and with the
exception of V482 Cyg (see below), we confirm the result for
carbon using C\,{\sc i} lines in the 9850 \AA\ region, the
N abundance using N\,{\sc i} lines in the 9850 \AA\ and 8727 \AA\ regions,
the Si abundance from Si\,{\sc i} lines in the  9850 \AA\ and 8727 \AA\ regions, and the Fe
abundance from Fe\,{\sc i} and Fe\,{\sc ii} lines in the 9850 \AA\ region.
The CN 1-0 lines are also used to check the N abundance.
In these atmospheres, molecules are trace species whose formation
does not affect the partial pressures of the constituent
atoms. Then,
consideration of the ratio of line to continuous opacity shows that 
the strength of a CN line should be dependent on the N abundance
 but  independent
of the C abundance.


  Following remarks about $\gamma$ Cyg,
 we discuss the RCB
stars in order of decreasing effective temperature.
It should be noted that the C abundances given here for the RCB
stars are not self-consistent -- the model atmospheres and their 
continuous absorption were calculated assuming a logarithmic
C abundance of 9.54. If the same abundance is adopted
for the calculations of the  carbon lines -- permitted and forbidden --
the predicted  and observed equivalent widths
are in disagreement. Thus, the C abundance derived here for the
RCB stars is basically a measure of the inconsistency, which we
term the carbon problem. 

\begin{figure}
\epsfxsize=8truecm
\epsffile{fig3.ps}
\caption{Observed (solid line) and synthetic
[C\,{\sc i}] 9850.264\AA\ line profiles of RCBs.
Synthetic profiles, including the blends, are shown for each star for three different
abundances. The maximum abundance of $\log\epsilon$(C) = 9.5, the model
input abundance is shown by the long dashes and in no case does this fit the
observed feature. The minimum abundance of $\log\epsilon$(C) = 7.5 shows in
effect the predicted spectrum for no contribution from the [C\,{\sc i}]
line (see dot-dash line). The best-fitting synthetic spectrum
is shown by the line made of short dashes and the corresponding
abundance is indicated in each panel.}
\end{figure}

\subsection{$\gamma$ Cyg}

Our analysis of permitted and forbidden C\,{\sc i} lines uses
a MARCS model atmosphere (Gustafsson et al. 1975)
for the atmospheric parameters derived by Luck \& Lambert  (1981).
The 8727 \AA\ [C\,{\sc i}] line (Figure 2) yields the abundance
$\log\epsilon$(C) = 7.88.
The 9850 \AA\ [{C\,{\sc i}] line is not detectable and the upper limit to
its equivalent width corresponds to $\log\epsilon$(C) $\leq$ 8.18.
 A selection of 14 C\,{\sc i} lines from 5380 \AA\ to 8873 \AA\ with
$gf$-values from Wiese, Fuhr \& Deters (1996) give the abundance
$\log\epsilon$(C) = 7.94 $\pm$ 0.14 (std. deviation). These lines range
in equivalent width from 6 m\AA\ to 82 m\AA. 
%The carbon abundance is sub-solar because the first dredge-up has
%brought CN-cycled material into the atmosphere. The N abundance
%is correspondingly supra-solar (Luck \& Lambert 1981).
For $\gamma$ Cyg, it is seen that the C\,{\sc i} lines and the
8727 \AA\ forbidden line return the same abundance within the errors.
This result is in sharp contrast to the  results
from the RCB stars.

%The C\,{\sc i} lines in the 9850 \AA\ region cannot be used 
%in the abundance analysis because
%their $gf$-values have not been determined by either theoretical or
%experimental means. These lines, however, are used to provide an
%`astrophysical' estimate of the $gf$-value for lines used in the
%construction of synthetic spectra for the RCB stars. 



\subsection{XX Cam}


 The 9850 \AA\ feature is dominated by
the Fe\,{\sc ii} line such that we can set only an
upper limit to the [C\,{\sc i}] contribution and the associated
C abundance of $\log\epsilon$(C) $\leq$ 8.8 (Figure 3). 
The 8727 \AA\ line (Figure 4) is fit with an abundance
$\log\epsilon$(C) = 8.4.  The limit from the 9850 \AA\ line
to the C abundance is consistent with that from the 8727 \AA\ line.
 The two C\,{\sc i}
lines near 9850 \AA\ give an abundance
consistent with the C abundance of $\log\epsilon$(C) = 9.0 $\pm$ 0.4
 from Asplund et al. who used 
 a collection of C\,{\sc i} lines and Opacity Project $gf$-values.


Asplund et al. put  the uncertainty in $T_{\rm eff}$  at about
$\pm$ 250 K which corresponds to 
corrections to abundances derived from   C\,{\sc i} and
[C\,{\sc i}] lines of  0.05, and 0.2 dex, respectively.
The uncertainty in $\log g$ $\simeq \pm$ 0.5 provides only
minor corrections to the abundances. An additional
potential source of uncertainty is the adopted microturbulent
velocity of $\xi = 9.0$ km s$^{-1}$ (Asplund et al. 2000). A change of
$\xi$ by $\pm$ 2 km s$^{-1}$ changes the C  abundance by 0.2 dex
from the C\,{\sc i} lines and by only 0.05 dex from the [C\,{\sc i}]
8727 \AA\ line. These error estimates do not allow the 0.6 dex
difference in the C\,{\sc i} and [C\,{\sc i}] abundances. 
%to be
%attributed to uncertainties arising from the adopted atmospheric
%parameters. This conclusion holds for the other stars
%where the 8727 \AA\ is observed.

%The original carbon problem uncovered by Asplund et al. is confirmed,
%the C\,{\sc i} lines give an abundance 0.5 dex less than
%the input abundance of the adopted model. 
The carbon problem for the
[C\,{\sc i}] line is even more severe as seen for
8727 \AA\ [C\,{\sc i}] line which gives an abundance 0.6 less than the
C\,{\sc i} lines and 1.1 dex less than the input abundance.

\begin{figure}
\epsfxsize=8truecm
\epsffile{fig4.ps}
\caption{Observed [C\,{\sc i}] 8727.126\AA\ line profile of XX Cam (solid lines).
 Synthetic spectra are shown for four C abundances, as shown on the figure.}
\end{figure}

\subsection{RY Sgr}

The 31 July 2002 spectrum  shows
emission in the blue wing of the 8727 \AA\ forbidden
line.  Peak emission occurs  at the velocity of $-$29 km s$^{-1}$.
This is
blue-shifted  by about 17 km s$^{-1}$ from the photospheric
velocity of $-$12 km s$^{-1}$. The emission line has a width (FWHM) of about
19 km s$^{-1}$. A spectrum  taken the same
night shows emission at 9850 \AA\ (Figure 1).
Emission at [C\,{\sc i}] may be a common occurrence for this star.
Our 1997 June 22 spectrum shows 9850 \AA\ in emission at the
radial velocity of $-$10.5 km s$^{-1}$ equivalent to a red-shift of
19.5 km s$^{-1}$ relative to the photosphere.
This spectrum does not
include the region around the 8727 \AA\ line.

RY Sgr exhibits a pulsation with an amplitude of about 35 km s$^{-1}$
(Lawson, Cottrell \& Clark 1991). The mean or systemic velocity was given as $-$21 km s$^{-1}$ by
Lawson, Cottrell \& Clark (1991). Relative to this velocity the emission at 8727 \AA\
in the 2002 spectrum is blue-shifted by about 8 km s$^{-1}$. 
Perhaps, the emission at  maximum light  is related to the
fact that RY Sgr is a  large-amplitude radial velocity
pulsator.  
%At certain phases the absorption lines
%are split into two components (Cottrell \& Lambert 1982)
%suggesting passage of a shock through the atmosphere. This
%shock might  be the site of emission. 
There is no hint of emission in our
spectra of the other RCBs, even in V482 Cyg observed 
below maximum light.



By inspection, the 9850 \AA\ blend of Fe\,{\sc ii} and [C\,{\sc i}]
lines is very similar to that of XX Cam and, hence, the forbidden
line is a minor contributor to the absorption feature.
 Synthesis of the 8727 \AA\ line (Figure 5)
suggests that emission overlays the absorption line and one might
suppose that an abundance $\log\epsilon$(C) $\sim$ 9.2 is a lower
limit but this  assumes that the emission
comes from a region exterior to the regions in which the absorption
line is formed. This is not necessarily a valid assumption 
for RY Sgr with its large pulsation and evidence for an internal
shock. We decline to quote a carbon abundance. But, a clue might be
derived from RY Sgr observations that the emission effects both [C\,{\sc i}]
lines. The estimations from the red absorption wing (assumed to be 
uneffected by emissions) of both these [C\,{\sc i}] lines suggests the 
same carbon abundance of about 9.2 which agrees with that from 
C\,{\sc i} lines.

%Inspection of Figure 1 shows that the CN lines are present. 
With the
adopted model, the CN lines return a N abundance about 0.6 dex higher than
from the N\,{\sc i} lines used here and by Asplund et al. This 
difference between N abundances can be eliminated by lowering the adopted
$T_{\rm eff}$ by about 200 K, or by a combination of a smaller $T_{\rm eff}$
change accompanied by a reduction of $\log g$. These adjustments are
within the error limits considered reasonable  by Asplund et al. 

\begin{figure}
\epsfxsize=8truecm
\epsffile{fig5.ps}
\caption{Observed [C\,{\sc i}] 8727.126\AA\ line profile of RY Sgr (solid lines).
 Synthetic spectra are shown for four C abundances, as shown on the figure.}
\end{figure}

\subsection{UV Cas}

The 8727 \AA\ line is reasonably well fit with the abundance
$\log\epsilon$(C) = 8.3 (Figure 6). Two observations 
show no convincing evidence for the 9850 \AA\ [C\,{\sc i}]
line, say $\log\epsilon$(C) $\leq$ 8.7 (Figure 3). The two
C\,{\sc i} lines in the 9850 \AA\ region suggest an abundance close to
$\log\epsilon$ = 9.0, a value only 0.2 dex less than that obtained by
Asplund et al. (2000). The CN lines are not detectable in our spectra.
UV Cas joins XX Cam in showing a larger carbon
problem for the 8727 \AA\ line than for the C\,{\sc i} lines.

\begin{figure}
\epsfxsize=8truecm
\epsffile{fig6.ps}
\caption{Observed [C\,{\sc i}] 8727.126\AA\ line profile of UV Cas (solid lines).
 Synthetic spectra are shown for four C abundances, as shown on the figure.}
\end{figure}

\subsection{VZ Sgr}

This RCB star was classified as a `minority' star by Lambert \&
Rao (1994), i.e., the metal lines are weak.
 This is
evident from Figure 1 which shows that the Fe\,{\sc i} and Fe\,{\sc ii} 
lines are absent from VZ Sgr's spectrum. Unfortunately,  the region around
8727 \AA\ has not yet been observed. 
%The 9850 \AA\ forbidden carbon line is not detected. 
Synthesis (Figure 3) suggest an upper limit
$\log\epsilon$(C) $\leq$ 8.7 for [C\,{\sc i}] 9850 \AA\  line, a limit clearly below the input
abundance of 9.5. The two C\,{\sc i} lines give 
$\log\epsilon$(C) = 9.1. Asplund et al. gave $\log\epsilon$(C) = 
8.8 $\pm$ 0.3.

\subsection{R CrB}

At 8727 \AA\, the synthesis corresponding to $\log\epsilon$(C) $\simeq$ 8.4
provides a fair fit to the observed high-resolution profile of 2002 July 31 (Figure 7).
A lower resolution (R = 60000) spectrum of 1995 June 17 provides
the same abundance. In neither spectra is there  a hint that
emission has distorted the absorption profile.

The 9850 \AA\ syntheses  include
the CN lines blending with the [C\,{\sc i}] and Fe\,{\sc ii}
feature. On the high-resolution spectrum of
2002 July 31, 
the CN 1-0 lines are blue-shifted by  3.6 km s$^{-1}$ relative to the
high-excitation lines of C\,{\sc i} and N\,{\sc i}.
Since the [C\,{\sc i}] line  and CN lines are all of low
excitation, we assume  
that they have a common  blueshift.
 Synthetic spectra (Figure 3) indicate
an abundance $\log\epsilon$(C) = 9.1. Examination of our library
of 9850 \AA\ spectra of R CrB (e.g., Rao \& Lambert 1997; Rao et al.
1999) shows that the velocity shift between CN and the C\,{\sc i} and
N\,{\sc i} varies during the star's pulsation and this shift
ranges from $-$3.0 to $+$4.0 km s$^{-1}$.
Syntheses taking into the account the velocity shift return the
same abundance from four different spectra.
If the blue-shift is neglected, the fit to the 9850 \AA\ feature
is less satisfactory but the derived carbon abundance is little
affected.
Asplund et al.'s atmospheric parameters
 are used in all cases.
%, even though Rao \& Lambert suggested these parameters change slightly during the pulsation. 
The C\,{\sc i} lines near 9850 \AA\ are well fitted with a
similar abundance ($\log\epsilon$(C) = 9.2), an abundance  equal to that
offered by Asplund et al. 

%As in the case of XX Cam,  the [C\,{\sc i}] lines add a new
%dimension to the carbon problem,
The 9850 \AA\ line confirms the
abundance obtained from C\,{\sc i} lines, but the 8727 \AA\
line gives an abundance 0.7 dex below that from the 9850 \AA\ line or
1.1 dex below the input carbon abundance of the model.

The N\,{\sc i} lines in the 8727 \AA\ and 9850 \AA\ regions give
the abundance $\log\epsilon$(N) = 8.2, a value consistent
with the result of 8.4 $\pm$ 0.2 given by Asplund et al. 
The observed CN 1-0
lines are well matched with the abundance $\log\epsilon$(N) = 8.2.
This may be fortuitous agreement because the sensitivity of the
CN line strengths to a change of the atmospheric parameters,
especially to $T_{\rm eff}$, is high. Note that a change of 250 K changes
the required abundance by 0.5 dex.

\begin{figure}
\epsfxsize=8truecm
\epsffile{fig7.ps}
\caption{Observed [C\,{\sc i}] 8727.126\AA\ line profile of R CrB (solid lines).
 Synthetic spectra are shown for four C abundances, as shown on the figure.}
\end{figure}


\subsection{SU Tau}

%Synthesis of the 15 November 2002 spectrum shows that 
The abundance $\log\epsilon$(C) = 8.2
provides an excellent fit to the observed [C\,{\sc i}] line
at 8727 \AA\ (Figure 8).
%The 9850 \AA\ feature is dominated
%by the contribution from the [C\,{\sc i}] line. 
A synthetic spectrum with the abundance
$\log\epsilon$(C) = 8.9 gives an excellent fit to the observed
9850 \AA\ feature (Figure 3).
For this fit, the  CN lines are blue-shifted by 6 km s$^{-1}$ from the
high excitation atomic lines, and this shift is assumed to apply to the
[C\,{\sc i}] line also, but the derived abundance is not
critically influenced by the shift. The abundance from the 9850 \AA\ 
[C\,{\sc i}] line is consistent with that derived by us from the
 the 9850 \AA\ region's C\,{\sc i} lines and the C\,{\sc i} lines
used by Asplund et al. In contrast and consistent with the results for
XX Cam and R CrB, the 8727 \AA\ line  gives a markedly
lower carbon abundance.  

The N abundance reported by Asplund et al. was based on just
two
N\,{\sc i} lines. Examination of our superior spectra provides a
more accurate equivalent width for one line. The other line is
not present on our spectrum. We identify a third line.
We adopt
Wiese et al.'s $gf$-value and derive the N abundance
$\log\epsilon$(N) = 7.9. We suggest that this is a more reliable
estimate than Asplund et al.'s value of 8.5. 
The N abundance estimate from the CN lines is 0.5 dex lower.
 The difference is erasable with
minor adjustments to the adopted atmospheric parameters, such as
an increase of $T_{\rm eff}$ by only 140 K.

\begin{figure}
\epsfxsize=8truecm
\epsffile{fig8.ps}
\caption{Observed [C\,{\sc i}] 8727.126\AA\ line profile of SU Tau (solid lines).
 Synthetic spectra are shown for four C abundances, as shown on the figure.}
\end{figure}


\subsection{V482 Cyg}

%Our spectrum includes the 9850 \AA\ but not the 8727 \AA\ region. 
There is a hint of emission at the 9850 \AA\
[C\,{\sc i}] line (Figure 3). Examination of the three
individual exposures shows that this
`emission' occurs  in only one exposure that being the one with the 
lowest signal-to-noise ratio. We, therefore, consider the emission
to be an artefact. A fit to the absorption feature suggests an abundance
upper limit $\log\epsilon$(C) $\leq$ 8.9.

In conflict with this upper limit,
the C abundance of  $\log\epsilon$(C) = 9.5$\pm$0.3 is derived
from C\,{\sc i} lines in our spectrum. Asplund et al. 
obtained $\log\epsilon$(C) = 8.9 from C\,{\sc i} lines.
Examination of our spectrum shows that C\,{\sc i} lines
across the spectrum are systematically stronger than reported
by Asplund et al. This strengthening and the adoption of the
same model as Asplund et al. necessarily results in a higher carbon
abundance. Recall that the star was two magnitudes below maximum
light. We suppose that the atmospheric structure differed from
that prevailing at maximum light. The temperature gradient
may have been steeper in decline than at maximum light. A very
similar strengthening of C\,{\sc i} (and other) lines was
noted by Rao et al. (1999) in spectra of R CrB taken early in 
a deep decline.  



The 
%9850 \AA\ 
N\,{\sc i} lines give an abundance 0.8 dex
higher than that from the CN lines. 
Erasure of the N\,{\sc i} -- CN discrepancy requires
a different choice of atmospheric parameters. The possibilities
range from raising $T_{\rm eff}$ by 400 K at the adopted
$\log g$ to lowering $\log g$ by about 0.8 dex at the adopted
$T_{\rm eff}$. These are not unacceptable changes given that
the atmosphere may have been perturbed.

\subsection{GU Sgr}

%Our observations cover the 9850 \AA\ but not the 8727 \AA\ region.  
The blend containing the 9850 \AA\ [C\,{\sc i}] line
is contaminated with CN lines which are strong in this spectrum.
Our syntheses
suggests an upper limit $\log\epsilon$(C) $\leq$ 9.0,
a value consistent with Asplund et al.'s value of 8.8. 


Analysis of the CN lines gives a N abundance of
$\log\epsilon$(N) = 8.1. Since the N\,{\sc i} lines in the
9850 \AA\ region are blended, we compare the CN-based abundance
with the result $\log\epsilon$(N) = 8.7 $\pm$ 0.5 given by
Asplund et al. The difference may be removed by a modest
adjustment of the adopted $T_{\rm eff}$ (say, a 250 K increase
at the adopted $\log g$).



\subsection{Summary of the Carbon Abundances}

%\begin{table*}
%\centering
%\begin{minipage}{105mm}
%\caption{Photospheric C abundances from C\,{\sc i} and [C\,{\sc i}] lines}
%\begin{tabular}{lccccccc} \hline
%\multicolumn{1}{c}{}& \multicolumn{3}{c}{Model$^a$}&  \multicolumn{1}{c}{}&\multicolumn{3}{c}{log$\epsilon$(C)} \\
%\cline{2-4} \cline{6-8}
%\multicolumn{1}{c}{Star}& \multicolumn{1}{c}{$T_{\rm eff}$} & \multicolumn{1}{c}{log$g$} & \multicolumn{1}{c}{$\xi$}&
%\multicolumn{1}{c}{}&\multicolumn{1}{c}{C\,{\sc i}$^a$} &\multicolumn{1}{c}{$[\rm C\,{\sc i}]$} &\multicolumn{1}{c}{$[\rm C\,{\sc i}]$}\\
%\multicolumn{1}{c}{}& \multicolumn{1}{c}{K}& \multicolumn{1}{c}{cm s$^{-2}$}&\multicolumn{1}{c}{km s$^{-1}$}&\multicolumn{1}{c}{}&
%\multicolumn{1}{c}{}&\multicolumn{1}{c}{9850\AA}&\multicolumn{1}{c}{8727\AA}\\ \hline
%XX~Cam & 7250 & 0.75 & 9.0 & & 9.0 & $<$8.8 & 8.4\\
%RY~Sgr & 7250 & 0.75 & 6.0 & & 8.9 & ... & ...    \\
%UV~Cas & 7250 & 0.50 & 7.0 & & 9.2 & $<$8.7 & 8.3 \\
%VZ~Sgr & 7000 & 0.50 & 8.0 & & 9.1 & $<$8.7 & ... \\
%R~CrB  & 6750 & 0.50 & 7.0 & & 9.2 & 9.1 & 8.4 \\
%SU~Tau & 6500 & 0.50 & 7.0 & & 8.8 & 8.9 & 8.2 \\
%V482~Cyg& 6500 & 0.50 &4.0 & & 8.9 & $<$8.9 & ...\\
%GU~Sgr & 6250 & 0.50 & 7.0 & & 8.8 &$<$9.0 & ... \\
%$\gamma$ Cyg$^b$ & 6100 & 0.55 & 3.5 & & 7.9 & $<$8.2 & 7.9 \\
%\hline
%\end{tabular}
%$^a$Asplund et al. (2000) for R~CrB stars; Luck \& Lambert (1981) for $\gamma$ Cyg\\
%$^b$see Section 5.1 of the text
%%$^a$covers 8727\AA\ region\\
%%$^b$covers the 8727\AA\ and 9850\AA\ regions
%\end{minipage}
%\end{table*}
\begin{table*}
\centering
\begin{minipage}{105mm}
\caption{Photospheric C abundances from C\,{\sc i} and [C\,{\sc i}] lines}
\begin{tabular}{lcccccccc} \hline
\multicolumn{1}{c}{}& \multicolumn{3}{c}{Model$^a$}&  \multicolumn{1}{c}{}&\multicolumn{4}{c}{log$\epsilon$(C)} \\ 
\cline{2-4} \cline{6-9}
\multicolumn{1}{c}{Star}& \multicolumn{1}{c}{$T_{\rm eff}$} & \multicolumn{1}{c}{log$g$} & \multicolumn{1}{c}{$\xi$}&
\multicolumn{1}{c}{}&\multicolumn{1}{c}{C\,{\sc i}$^a$} &\multicolumn{1}{c}{C\,{\sc i}$^b$} &\multicolumn{1}{c}{$[\rm C\,{\sc i}]$} 
&\multicolumn{1}{c}{$[\rm C\,{\sc i}]$}\\
\multicolumn{1}{c}{}& \multicolumn{1}{c}{K}& \multicolumn{1}{c}{cm s$^{-2}$}&\multicolumn{1}{c}{km s$^{-1}$}&\multicolumn{1}{c}{}&
\multicolumn{1}{c}{} &\multicolumn{1}{c}{} &\multicolumn{1}{c}{9850\AA}&\multicolumn{1}{c}{8727\AA}\\ \hline
RCB stars: &  &      &     & &     &     &        & \\
XX~Cam & 7250 & 0.75 & 9.0 & & 9.0 & 8.9 & $<$8.8 & 8.4\\
RY~Sgr & 7250 & 0.75 & 6.0 & & 8.9 & 9.3 & ... & ...    \\
UV~Cas & 7250 & 0.50 & 7.0 & & 9.2 & 9.0 & $<$8.7 & 8.3 \\
VZ~Sgr & 7000 & 0.50 & 8.0 & & 9.1 & 9.1 & $<$8.7 & ... \\
R~CrB  & 6750 & 0.50 & 7.0 & & 9.2 & 9.2 & 9.1 & 8.4 \\
SU~Tau & 6500 & 0.50 & 7.0 & & 8.8 & 9.0 & 8.9 & 8.2 \\
V482~Cyg& 6500 & 0.50 &4.0 & & 8.9 & 9.6 & $<$8.9 & ...\\
GU~Sgr & 6250 & 0.50 & 7.0 & & 8.8 & 9.3 & $<$9.0 & ... \\
Standard star: &  &      &     & &     &     &        & \\
$\gamma$ Cyg$^c$ & 6100 & 0.55 & 3.5 & & 7.9 & ... & $<$8.2 & 7.9 \\
\hline
\end{tabular}
$^a$Asplund et al. (2000) for R~CrB stars, the expected C abundance is 9.54 when C/He of 1\% models are used; Luck \& Lambert (1981) for $\gamma$ Cyg\\
$^b$based on two C\,{\sc i} lines near 9850\AA\\
$^c$see Section 5.1 of the text
%$^a$covers 8727\AA\ region\\
%$^b$covers the 8727\AA\ and 9850\AA\ regions
\end{minipage}
\end{table*}

Our analyses use the appropriate MARCS model for the
atmospheric parameters recommended by Asplund et al. Throughout
models constructed for C/He = 1\% by number are used. This input
abundance corresponds to a carbon abundance $\log\epsilon$(C) = 9.5.
In Table 2, we summarize the carbon abundances derived from
the two [C\,{\sc i}] lines,  
and the C\,{\sc i}-based abundance given by
Asplund et al. illustrating the carbon problem.
%Across the table there is a carbon problem, in that the tabulated
%abundances are systematically less than this input abundance.

Except for V482 Cyg, the carbon abundance derived from
the two 9850 \AA\ region  C\,{\sc i} lines
 is in good agreement with that obtained
by Asplund et al. 
%from a larger sample of C\,{\sc i} lines at
%shorter wavelengths. 
%Recall that V482 Cyg was observed below
%maximum light.

%Detection of the 9850 \AA\ forbidden line for R CrB and SU Tau
%leads to a carbon abundance consistent with that from the
%C\,{\sc i} lines.
%Even for stars where 
%an upper limit to the carbon abundance is provided from the forbidden
%line  at
%9850 \AA\, the limit is clearly below the input abundance 
Table 2 shows 
that a carbon problem extends to 9850 \AA\ forbidden line.
This conclusion is contrary to that reached by Asplund et al.
from analysis of equivalent widths of the 9850 \AA\ line.
The explanation
is that Asplund et al. did not recognize 
%in their lower resolution spectra 
that the [C\,{\sc i}] line was a
blend. In the case of V482 Cyg observed two magnitudes below
maximum light, the carbon problem has vanished for the C\,{\sc i}
lines but not for the 9850 \AA\ [C\,{\sc i}] line. This is the
only case where the C\,{\sc i} lines and 9850 \AA\ differ in the
abundance they provide. 

Surprisingly, the 8727 \AA\ [C\,{\sc i}] line offers further
information on
the carbon problem. For each of the four stars for
which we have observed the 8727 \AA\  line, the derived
abundance is less than that from the C\,{\sc i} lines. 
In the case of R CrB and SU Tau, 
%the two stars for which a carbon
%abundance is obtained from both forbidden lines,
 the 8727 \AA\ line gives an abundance 0.7 dex
less than  that from the 9850 \AA\  line. 
That the two forbidden lines give very different abundances is
especially puzzling if these lines are formed, as expected, in or close
to LTE and, as might be supposed, are a product of  the stellar
photosphere, i.e., a region with the temperature decreasing
monotonically outward. 
An unidentified atomic line superimposed on the 8727 \AA\
forbidden line would mean that we have overestimated
the carbon abundance from that line. 
In this scenario, the carbon problem is more severe
for both of the [C\,{\sc i}] than for the C\,{\sc i} lines.
However, it is difficult to find a carrier for a line which
is strong in R CrB and SU Tau but weak in other RCB stars and
$\gamma$ Cyg.





%\begin{table}
%\centering
%\begin{minipage}{55mm}
%\caption{Carbon abundances for R CrB for various model atmospheres}
%\begin{tabular}{cccccc} \hline
%\multicolumn{1}{c}{}&\multicolumn{1}{c}{}&\multicolumn{1}{c}{}& \multicolumn{3}{c}{log$\epsilon$(C)} \\
%\multicolumn{3}{c}{Model}& \multicolumn{1}{c}{C\,{\sc i}} &
%\multicolumn{1}{c}{$[\rm C\,{\sc i}]$} &\multicolumn{1}{c}{$[\rm C\,{\sc i}]$}\\
%\cline{1-3}
%\multicolumn{3}{c}{$T_{\rm eff}$,log$g$,$\xi$}& \multicolumn{1}{c}{}&
%\multicolumn{1}{c}{9850\AA}&\multicolumn{1}{c}{8727\AA}\\ \hline
%&6750,0.5,7&&9.2&9.1&8.4\\
%&6500,0.5,7&&9.2&8.9&8.2\\
%&7000,0.5,7&&9.3&9.3&8.6\\
%&6750,0.0,7&&9.3&9.1&8.4\\
%&6750,1.0,7&&9.1&9.1&8.4\\
%&6750,1.0,9&&8.9&9.1&8.4\\
%&6750,1.0,5&&9.3&9.1&8.4\\
%\hline
%\end{tabular}
%\end{minipage}
%\end{table}

The sensitivities of the permitted and forbidden lines to
the choice of model atmosphere are different. To illustrate
these sensitivities, we give in Table 3 the abundances derived from
R CrB's lines for a series of model atmospheres centred on 
Asplund et al.'s choice of $(T_{\rm eff}, \log g, \xi)$
= (6750, 0.5, 7). The carbon abundance from C\,{\sc i} lines is,
as expected from the insensitivity of the ratio of line to continuous
opacity to physical conditions, almost independent of the choice of
model. The abundance is also insensitive to the choice of the
microturbulent velocity $\xi$ for a weak line but not for a strong
line. In Table 3, the mean C abundance derived from weak and strong
permitted lines is given. The abundance from the forbidden
lines is insensitive to the choice of the surface gravity but
dependent on $T_{\rm eff}$. We note that the change $\pm$ 250 K in
$T_{\rm eff}$ leads to
a change in the carbon abundance of $\pm$ 0.2 dex. There is no
dependence on $\xi$. The $T_{\rm eff}$ and $\log g$ sensitivities
are not very different for the hottest stars (XX Cam and RY Sgr)
and the coolest (GU Sgr). 

To increase the carbon abundance from the 9850 \AA\ [C\,{\sc i}]
line to the input abundance ($\log\epsilon$(C) = 9.5), requires
that the $T_{\rm eff}$ of R CrB and SU Tau be raised about 
500 K, and 750 K, respectively. These increases are not
only outside the bounds considered acceptable by Asplund
et al., but they do not remove the carbon problem for the
permitted carbon lines. In addition,  they introduce
a  discrepancy between the
N abundance from the N\,{\sc i}  and CN lines. Also note that
the higher temperatures do not eliminate the abundance difference
of 0.7 dex from the 9850 \AA\ and 8727 \AA\ forbidden lines
of R CrB and SU Tau. In short, the [C\,{\sc i}] lines are
part of a now enlarged carbon problem.     

\begin{table*}
\centering
\begin{minipage}{55mm}
\caption{Carbon abundances for R CrB for various model atmospheres}
\begin{tabular}{cccccc} \hline
\multicolumn{1}{c}{}&\multicolumn{1}{c}{}&\multicolumn{1}{c}{}& \multicolumn{3}{c}{log$\epsilon$(C)} \\
\multicolumn{3}{c}{Model}& \multicolumn{1}{c}{C\,{\sc i}} &
\multicolumn{1}{c}{$[\rm C\,{\sc i}]$} &\multicolumn{1}{c}{$[\rm C\,{\sc i}]$}\\
\cline{1-3}
\multicolumn{3}{c}{$T_{\rm eff}$,log$g$,$\xi$}& \multicolumn{1}{c}{}&
\multicolumn{1}{c}{9850\AA}&\multicolumn{1}{c}{8727\AA}\\ \hline
&6750,0.5,7&&9.2&9.1&8.4\\
&6500,0.5,7&&9.2&8.9&8.2\\
&7000,0.5,7&&9.3&9.3&8.6\\
&6750,0.0,7&&9.3&9.1&8.4\\
&6750,1.0,7&&9.1&9.1&8.4\\
&6750,1.0,9&&8.9&9.1&8.4\\
&6750,1.0,5&&9.3&9.1&8.4\\
\hline
\end{tabular}
\end{minipage}
\end{table*}

%\begin{table}
%\centering
%\begin{minipage}{55mm}
%\caption{Carbon abundances for R CrB for various model atmospheres}
%\begin{tabular}{cccccc} \hline
%\multicolumn{1}{c}{}&\multicolumn{1}{c}{}&\multicolumn{1}{c}{}& \multicolumn{3}{c}{log$\epsilon$(C)} \\
%\multicolumn{3}{c}{Model}& \multicolumn{1}{c}{C\,{\sc i}} &
%\multicolumn{1}{c}{$[\rm C\,{\sc i}]$} &\multicolumn{1}{c}{$[\rm C\,{\sc i}]$}\\
%\cline{1-3}
%\multicolumn{3}{c}{$T_{\rm eff}$,log$g$,$\xi$}& \multicolumn{1}{c}{}&
%\multicolumn{1}{c}{9850\AA}&\multicolumn{1}{c}{8727\AA}\\ \hline
%&6750,0.5,7&&9.2&9.1&8.4\\
%&6500,0.5,7&&9.2&8.9&8.2\\
%&7000,0.5,7&&9.3&9.3&8.6\\
%&6750,0.0,7&&9.3&9.1&8.4\\
%&6750,1.0,7&&9.1&9.1&8.4\\
%&6750,1.0,9&&8.9&9.1&8.4\\
%&6750,1.0,5&&9.3&9.1&8.4\\
%\hline
%\end{tabular}
%\end{minipage}
%\end{table}


\section{Discussion}

Possible solutions to the carbon problem presented by the
C\,{\sc i} lines were 
discussed by Asplund et al. (2000 --see also Gustafsson \& Asplund 1996).
 Two proposed
solutions were suggested by Asplund et al.
as worthy of further consideration: the $gf$-values
of the C\,{\sc i} lines are in error, or theoretical model
atmospheres are a misrepresentation of the temperature
structure of real RCB stars. In addition to  commenting
on these solutions as they affect the forbidden lines,
we examine
 departures from LTE as they affect  
line formation. Finally, we discuss hand-crafted model atmospheres
including models
with a chromosphere (a temperature rise in the outer layers).

\subsection{The $gf$-values}

The $gf$-values of the [C\,{\sc i}] lines are known accurately from
theoretical calculations. 
%Differences at the 10 \% level may be
%debated but the large and different decreases in the $gf$-values
%of the 8727 \AA\ and 9850 \AA\ lines that are  needed
%to eliminate the carbon problem are outside the realm of
%possibility. 
Confirmation and extension of the carbon problem from the forbidden
lines strongly suggests that the $gf$-values of the permitted
lines can  not be primarily responsible for the problem.

\subsection{ Non-LTE Effects}

Non-LTE effects in neutral carbon atoms in RCB atmospheres
were evaluated by Asplund \& Ryde (1996). Their  results
indicate that departures from LTE are confined to shallow
optical depths. Effects on the C\,{\sc i} lines are slight
because a  line and the continuum
 are formed deep in the atmosphere between
levels in the
carbon atom with small and very similar departures from LTE.
Gustafsson \& Asplund (1996) give the correction to LTE
abundance from C\,{\sc I} lines
 as less than 0.1 dex and suggest that $+$0.02 dex is a typical
value. Thus, Non-LTE effects can not account for the
0.6 dex typical carbon problem arising from the permitted lines.
Departure
coefficients given by Asplund \& Ryde show that a non-LTE  correction
to the LTE abundance from a forbidden line  is also very small
and cannot be the source of the line's carbon problem.
These non-LTE calculations adopt a model computed in
LTE. 

Carbon problems may be reduced or even eliminated by invoking
a chromosphere (see below).
A temperature rise at the top of the photosphere
is possible even in the absence of mechanical energy deposition.
This non-LTE effect  sometimes called the Cayrel mechanism
(Cayrel 1963) is discussed by Mihalas (1978).  To assess the
temperature rise for a RCB star due to the Cayrel mechanism,
non-LTE effects on the populations of the carbon atom were
included in the computation of the model atmosphere.
These non-LTE effects on the carbon level-populations were calculated 
using MULTI (Carlsson 1986), with the modelled carbon atom 
described in Asplund \& Ryde (1996). Realistic background fluxes, computed with 
the MARCS model atmosphere code, were included. Typical departure coefficients 
(b=n/n(LTE)) where n is the occupation number density) of about b = 1 for the
lower C\,{\sc i}
 levels, b = 1.4 for intermediate ones and b = 0.6 for the upper ones 
were obtained. These figures were found to be characteristic of the depth 
interval $- 4 < \tau_{\mathrm{Ross}} < - 3$, with much smaller opposite effects 
further
in. These departures from LTE should mainly affect the opacities. C\,{\sc i} is also 
an
important electron contributor, but our calculations show only an 
underionization
of less than 10\% at shallow depths.

We then brought in these results into the MARCS program
in order to include their effect  on  the
continuous opacities. This procedure is not fully self-consistent - the models 
are still LTE models - but the main effect of the departures on the opacities 
are modelled in a reasonable way, by essentially diminishing the 
C\,{\sc i} opacities 
longwards of 6000 \AA, to 0.6 times the LTE value, and increasing them 
shortwards of that wavelength to 1.4 times the LTE value. This was applied for 
the
uppermost atmosphere with a decrease corresponding to our non-LTE results at 
greater depth.
A Cayrel effect resulted, as expected, with a heating of the upper layers, but 
only by
about 1-6 K, the maximum occuring at $\tau_{\mathrm{Ross}} = 0.01$. 

Thus, the resultant Cayrel effect is small for C\,{\sc i} in RCB stars with a 
temperature increase less than 10 K at the top of the photosphere. This 
micro-chromosphere does not alleviate the carbon problem! The small Cayrel 
effect is partly
linked to the fact that carbon atoms are 
 not totally dominating the scene. In fact at 
these relevant depths photoionization of carbon atoms
only contributes typically 50\% of the opacity. So, even drastic  
changes (at least reductions) to the carbon opacity
 only lead to moderate changes of fluxes 
and of the energy balance. Another circumstance is that the ultraviolet
flux in the 
surface layers is so weak that it does not seem to matter very much and is unlike 
the situation for hot stars and the H\,{\sc i} opacity, only a tiny fraction of 
the flux of typical RCB stars comes at energies higher than the ionization 
threshold, or even at energies able to photoionize from lower excited C\,{\sc i}
levels.

\subsection{Hand-crafted Models}

The fact that the carbon problems defined by the C\,{\sc i}
and the [C\,{\sc i}]
lines are very similar from star-to-star across the sample
of analysed RCB stars implies that the problems' resolution cannot depend
very sensitively
on a star's individual characteristics.
Departures of the atmospheric structure from that predicted
by a standard (MARCS) model must be very similar across the
sample of stars. If departures are attributable to deposition of
mechanical energy, the flux of mechanical energy cannot vary
widely across the sample. 

That the atmospheric structure is partly the culprit is suggested
by the observation that a RCB in a decline shows a reduced
carbon problem. The case of V482 Cyg was discussed in
Section 5.8. There, the carbon problem for the C\,{\sc i} lines
vanished. Note that the C\,{\sc i} lines are noticeably stronger in our
spectrum than in spectra taken at maximum light.
 The problem remains for the 9850 \AA\ forbidden line, but
this might be the result of overlying emission, as occurred in the
case of RY Sgr. (Unfortunately, the 8727 \AA\ line was not on the
recorded spectrum.) A similar strengthening of C\,{\sc i} lines
leading to elimination of the carbon problem but with retention
of the problem for the [C\,{\sc i}] 9850 \AA\ line
was seen by Rao et al. (1999) for R CrB in decline.

\subsubsection{The photospheric temperature gradient}

The carbon problem may be related to an incorrect representation by the
MARCS RCB models of the
temperature gradient through the region of line formation; 
a steeper/shallower gradient produces a stronger/weaker
 line. In this connection, it
may be noted that  Lambert \& Rao's (1994) abundance analysis
used unblanketed model atmospheres (Sch\"{o}nberner 1975) and
found  fair agreement between the observed and calculated
equivalent widths of the C\,{\sc i} lines. Note that the small carbon problem of
0.2 dex  might be attributed to a combination of
identifiable errors. The primary factor responsible for 
the negligible carbon problem
is that Sch\"{o}nberner's
unblanketed models have a shallower temperature gradient in the
line forming region than the
MARCS (line blanketed) RCB models (Gustafsson \& Asplund 1996).

 Unblanketed models were
 calculated with the MARCS code. 
Our abundance analysis for R CrB using unblanketed MARCS models
shows that the carbon problems presented by
permitted and forbidden C\,{\sc i} lines are similar for
unblanketed and blanketed MARCS models. 
 The carbon abundance derived from [C\,{\sc i}] 9850 \AA\
line using unblanketed  models again agrees that from the 
 C\,{\sc i} lines. The large difference between the carbon abundance
derived from the [C\,{\sc i}] 9850 \AA\ line
 and the [C\,{\sc i}] 8727 \AA\ line using  
the  blanketed models is unaltered by using the unblanketed
models. The difference between the unblanketed MARCS and Sch\"{o}nberner's
models reflects differing  temperature gradients of the
models.
  

Asplund et al. describe hand-crafted adjustments to the line blanketed models that
reduce the photospheric temperature gradient and alleviate the carbon problem
presented by the C\,{\sc i} lines.
A shallower gradient (with respect to a blanketed MARCS model) 
 was qualitatively justified as arising from  
deposition of mechanical energy whose source may be the
sub-photospheric convection zone.
Examples of  temperature modified models, which alleviate the
carbon problem from C\,{\sc i} lines, were kindly provided by
Martin Asplund (private communication). Our analyses using these
models show that the carbon problem vanishes for the
9850 \AA\ [C\,{\sc i}] line as it does for the C\,{\sc i} lines, but
the 
8727 \AA\ [C\,{\sc i}] line continues to present a problem;
 the carbon abundance from 
8727 \AA\ [C\,{\sc i}] line is about 0.7 dex less than that from 
9850 \AA\ [C\,{\sc i}] line, a difference found also for 
unmodified line blanketed and unblanketed models.

\subsubsection{A chromosphere}

Addition of a temperature rise (`a chromosphere')
at the top of the photosphere may also modify the predicted
strengths of the  absorption lines in the emergent spectrum.
The chromosphere must be crafted 
by hand without theoretical guidance. 




\begin{figure}
\epsfxsize=8truecm
\epsffile{fig9.ps}
\caption{Three $T$ $-$ $\tau_{\mathrm Ross}$ structures are shown - see key at the top
of the figure.  The solid line shows a MARCS model for  $T_{\rm eff}$ = 6750
and log$g$ = 0.5. The modified   $T$ $-$ $\tau$ structure with a heated plateau
devised by Asplund et al. (2000) to eliminate the carbon problem for
C\,{\sc i} lines is shown by the dotted line. Our model with a chromosphere
is shown by the dashed line.}
\end{figure}



Experiments were made by adding a chromosphere to a MARCS line-blanketed
 model.
Figure 9 shows one experiment appropriate for R CrB itself.
A temperature rise 
is introduced 
at $\log\tau_{\mathrm{Ross}} = -$1.6.
 A flat temperature profile of 7000 K in the interval
$\log\tau_{\mathrm{Ross}} \le -$1.6 and a LTE spectrum synthesis
% (with $n_e$ $\approx$ 2 $\times$ 10$^{11}$ cm$^{-3}$)
reproduce the profiles of the observed
 [C\,{\sc i}] 9850 \AA\ and [C\,{\sc i}] 8727 \AA\ absorption lines. 
The two lines now return the same C abundance which is equal to
the input abundance, or, the carbon problem vanishes for these lines. 
The  weaker
observed C\,{\sc i} lines are also reproduced with the input abundance.
Very strong C\,{\sc i} lines are predicted, as expected for an LTE
calculation, to have
emission cores.
The emission cores might be reduced
if the non-LTE effects were taken into consideration.
Quite obviously, the chosen chromosphere is not a unique solution. If the onset of
the sharp temperature rise is placed as in Figure 9, a chromospheric
temperature of 6750 to 7000 K  provides acceptable
solutions to the carbon problems. Similar chromospheres remove the
carbon problems for the other RCBs. The chromospheric
temperature may scale with a star's effective temperature, but the range of
sampled effective temperatures is small. 

Addition of a chromosphere not only changes the predicted
strengths of the C\,{\sc i} and [C\,{\sc i}] lines  but also the
predicted strengths of other lines and, therefore, the derived
abundances. High-excitation lines such as the N\,{\sc i} line at 8729 \AA\
are so weakened by chromospheric emission that the  observed 
absorption line in XX Cam and R\,CrB cannot be fit whatever the
adopted nitrogen abundance. 
Although
careful crafting of the chromospheric temperature profile may
alleviate this difficulty, it seems likely that a non-LTE
synthesis may be necessary. One may need to adopt
spherical rather than a plane-parallel geometry. The chromosphere, if
it exists, is likely to be highly-structured and a far  cry from
the homogeneous plane-parallel (or spherical) layers assumed
for the models. 



Although the artifice of a chromosphere serves to reduce or even
eliminate the carbon problems, its presence presents another
puzzle.
The energy flux needed to produce the extra heating
suggested here is substantial. Following Asplund et al. (2000),
we find it to be on the order of 10 -- 15 \% of the total
stellar flux. For comparison, the chromospheric--coronal heating of
normal late-type stars is limited to about 1 \% of the total
bolometric flux. A clue to the problem may lie in the very large
widths of RCB lines -- compare the widths of lines (Figure 1 and 2)
in $\gamma$ Cyg and the RCBs. These large widths imply a very turbulent
atmosphere and dissipation of energy. The possibility of an external 
radiation field generating very large widths of RCB lines is to
be explored.


Emission lines appear when a RCB goes into decline. The
emitting gas is not necessarily to be identified with the
hand-crafted chromosphere. Emission from a shell around the
star would dilute the photospheric absorption lines. 
Emission in the 9850 \AA\ and 8727 \AA\ lines might
weaken the absorption lines such that the emission is not
seen but the weakened absorption lines  return
different carbon abundances. The [C\,{\sc i}] lines
are seen in emission in spectra taken in decline (Rao \& Lambert
1993; Rao et al. 1999).
The measured fluxes in decline are possibly lower
limits to the fluxes at maximum light as the dust cloud
responsible for the decline presumably obscures parts of the
emitting shell.
 At maximum light, the cores of the strongest low
excitation absorption lines of singly-ionized metals appear
distorted by emission (Lambert, Rao \& Giridhar 1990; Rao et al. 1999).
It is just such lines which are prominent in the emission
line spectrum.  
 The emission line spectrum is of low
excitation and does not include the C\,{\sc i} lines.
Therefore, the emitting shell which may contribute to the
[C\,{\sc i}] problem is not a player in the C\,{\sc i}
problem.
In this context, it is relevant to recall the case of V482 Cyg
observed below maximum light where
the carbon problem is absent for the
C\,{\sc i} lines but not for the [C\,{\sc i}] 9850 \AA\
line. Our speculation is that the emitting shell fills in the forbidden
line and the photospheric temperature gradient is shallower than 
at maximum light.





\section{Concluding remarks}

The intriguing carbon problem presented by the
RCB stars 
%and their C\,{\sc i} lines 
was discovered 
by Asplund et al. (2000) in an application of MARCS RCB
models to analysis of C\,{\sc i} lines. Our contribution to the problem has been
to present and analyse observations of the [C\,{\sc i}] 8727 \AA\
and 9850 \AA\ lines, also with the MARCS models.
 The 9850 \AA\ line
presents a problem of similar magnitude (about 0.6 dex) 
%to that provided by the 
as C\,{\sc i} lines,
but the carbon problem presented
by the 8727 \AA\ line is more severe by about 0.7 dex than the original
problem (about 0.6 dex).
%discovered by Asplund et al. from
%their analyses of C\,{\sc i} lines. The original carbon problem has been
%given a new dimension.
 

The fact that the carbon problems defined by the C\,{\sc i}
and the [C\,{\sc i}]
lines are similar from star-to-star across the sample
of analysed RCB stars implies that the solution cannot depend
on a star's individual characteristics.
Departures of the atmospheric structure from that predicted
by a standard (MARCS) model must be  similar across the
sample of stars. If departures are attributable to deposition of
mechanical energy, the flux of deposited mechanical energy cannot vary
widely across the sample. 

Uncovering the carbon problem was a surprise. Confidence in model
atmospheres was shaken. Extension of  the problem to the [C\,{\sc i}]
lines is also discomforting.
In spite of the fact that the RCB stars in several ways are very
complex, they have the one property that should make the analysis
of C\,{\sc i} lines in principle simpler than for almost all other
stars. This is the fact that the lines are due to the same element is the dominant
opacity source and lines and opacity originate from levels of 
similar character and excitation. This provides a test, that is not offered for
other stars. This test, however, fails.



 Perhaps, the problems are restricted
to very peculiar stars such as the rare RCB stars. But this
supposition deserves to be tested. Luminous normal (i.e., H-rich) supergiants 
of the temperature of the RCB stars are the stars from
which chemical compositions of external galaxies are now being
derived. Are their atmospheric structures  reliably simulated
by our codes? Or is there a problem analagous to the carbon problem
awaiting discovery?



This research has been supported in part by The Robert A. Welch
Foundation through a grant to DLL.




\begin{thebibliography}{99}

\bibitem{} Allende Prieto, C., Lambert, D. L., Asplund, M., 2002, ApJ, 573, L137
\bibitem{} Asplund, M., Gustafsson, B., Kiselman, D., Eriksson, K., 1997, A\&A, 318, 521
\bibitem{} Asplund, M., Gustafsson, B., Lambert, D. L., Rao, N. K., 2000, A\&A, 353, 287
\bibitem{} Asplund, M., Ryde, N., 1996, in:
     Hydrogen deficient stars, Jeffery C.S., Heber U.
    (eds.). ASP conf. series vol. 96, p. 57
\bibitem{} Bauschlicher, C. W., Jr., Langhoff, S. R., Taylor, P. R., 1988, ApJ, 332, 531
\bibitem{} Bakker, E. J., Lambert, D. L., 1998, ApJ, 502, 417
\bibitem{} Carlsson, M., 1986, Uppsala Astronomical Report, No. 33
\bibitem{} Cayrel, R., 1963,  C. R. Acad. Sci., 257, 3309
\bibitem{} Cottrell, P. L., Lambert, D. L., 1982, ApJ, 261, 595
\bibitem{} Davis Sumner, P., Phillips John, G., 1963, The red system (A$^2$$\Pi$$-$X$^2$$\Sigma$) of the CN molecule,
     Berkeley Analyses of Molecular Spectra, Berkeley: University of California Press.
\bibitem{} Galavis, M. E., Mendoza, C., Zeippen, C. J., 1997, A\&AS, 123, 159
\bibitem{} Gustafsson, B., Asplund, M., 1996, in:
     Hydrogen deficient stars, Jeffery C.S., Heber U.
    (eds.). ASP conf. series vol. 96, p. 27
\bibitem{} Gustafsson, B., Bell, R. A., Eriksson, K., Nordlund, A., 1975, A\&A, 42, 407
\bibitem{} J{\o}rgensen, U. G., Larsson, M, 1990, A\&A, 238, 424
\bibitem{} Lambert, D. L., Rao, N. K., Giridhar, S., 1990, JAA, 11, 475
\bibitem{} Lambert, D. L., Rao, N. K., 1994, JAA, 15, 47
\bibitem{} Lawson, W. A., Cottrell, P. L., Clark, M., 1991, MNRAS, 251, 687
\bibitem{} Liu, X. W., Barlow, M. J., Danziger, I. J., Clegg, R. E. S., 1995, MNRAS, 273. 47
\bibitem{} Luck, R. E., Lambert, D. L., 1981, 245, 1018
\bibitem{} Mihalas, D. 1978, Stellar Atmospheres, Freeman \& Co.:San Francisco
\bibitem{} Moore, Ch. E.: 1972, A Multiplet Table of Astrophysical Interest,
    NSRDS - NBS 40, Washington
\bibitem{} Moore, Ch. E.: 1993, Tables of Spectra of Hydrogen, Carbon, Nitrogen, and
    Oxygen Atoms and Ions, Editor: Jean W. Gallagher, CRC Series in Evaluated Data
    in Atomic Physics, CRC press
\bibitem{} Nave, G., Johansson, S., Learner, R. C. M., Thorne, A. P., Brault, J. W., 1994, ApJS, 94, 221
\bibitem{} Rao, N.K., Lambert, D.L., 1990, AJ, 105, 1915
\bibitem{} Rao N. K., Lambert D. L.,1996, in:
    Hydrogen deficient stars, Jeffery C.S., Heber U.
    (eds.). ASP conf. series vol. 96, p. 43
\bibitem{} Rao N. K., Lambert D. L., 1997, MNRAS, 284, 489
\bibitem{} Rao N. K., Lambert D. L., Adams, M. T., Doss, D. R., Gonzalez, G., Hatzes, A. P.,
    James, R., Johns-Krull, C. M., Luck, R. E., Pandey, G., Reinsch, K.,
    Tomkin, J., 1999, MNRAS, 310, 717
\bibitem{} Sch\"{o}nberner, D., 1975, A\&A, 44, 383
\bibitem{} Searle, L., 1961, ApJ, 133, 531
\bibitem{} Swensson, J. W., Benedict, W. S., Delbouille, L., Roland, G., 1970, The
    Solar Spectrum from $\lambda$7498 to $\lambda$12016. A Table of Measures and
    Identifications. Mem. Soc. R. Sci. Li\`{e}ge, Spec. Vol. No. 5
\bibitem{} Tull, R. G., MacQueen, P. J., Sneden, C. and Lambert, D. L., 1995, PASP,
    107, 251
\bibitem{} Wiese, W. L., Fuhr, J. R., Deters, T. M., 1996, Journal of Physical and Chemical Reference Data, Monograph No. 7

\end{thebibliography}


\end{document}










astro-ph0604023.extracted_bib
\begin{thebibliography}{99}

\bibitem{} Asplund, M., Gustafsson, B., Lambert, D. L., Rao, N. K., 2000, A\&A, 353, 287
\bibitem{} Bond, H. E., Luck, R. E., Newman, M. J., 1979, ApJ, 233, 205 
\bibitem{} Drilling, J. S., 1973, ApJ, 179, L31
\bibitem{} Drilling, J. S., Sch\"{o}nberner, D., Heber, U., Lynas-Gray, A. E., 1984, ApJ, 278, 224
\bibitem{} Jeffery, C. S., Starling, R. L. C., Hill, P. W., Pollacco, D. L., 2001, MNRAS, 321, 111
\bibitem{} Jeffery, C. S., Woolf, V. M., Pollacco, D. L., 2001, A\&A, 376, 497
\bibitem{} Pandey, G., Rao, N. K., Lambert, D. L., Jeffery, C. S., Asplund, M., 2001, MNRAS, 324, 937
\bibitem{} Pandey, G., Lambert, D. L., Rao, N. K., Jeffery, C. S., 2004, ApJ, 602, L113
\bibitem{} Pandey, G., Lambert, D. L., Jeffery, C. S., Rao, N. K., 2006, ApJ, 638, 454
\bibitem{} Rao, N. K., Lambert, D. L., 2003, PASP, 115, 1304 
\bibitem{} Stephenson, C. B., Sanduleak, N., 1971, Publ. Warner \& Swasey Obs. 1, 1
\bibitem{} Th\'{e}venin, F., 1989, A\&AS, 77, 137
\bibitem{} Th\'{e}venin, F., 1990, A\&AS, 82, 179
\bibitem{} Vanture, A. D., Zucker, D., Wallerstein, G., 1999, ApJ, 514, 932 
\bibitem{} Wiese, W. L., Fuhr, J. R., Deters, T. M., 1996, {\it J. Phys. Chem. Ref. Data}, Monograph No. 7

\end{thebibliography}