New abundance measurements in \uks, a very metal-poor galaxy in the Local Group \thanks{Based on data collected at the European Southern Observatory, La Silla, Chile, prop. No. 63.N-0726B}
I. Saviane\inst{1}
L. Rizzi\inst{2,3}
E. V. Held\inst{2}
F. Bresolin\inst{4}
Y. Momany\inst{3}
%% is: added 'brightest' We present new results from optical spectroscopy of the brightest H{\sc ii} region in the dwarf irregular galaxy \uks\ in Sagittarius (SagDIG). From high signal-to-noise spectra, reddening-corrected line flux ratios have been measured with typical uncertainties of a few percent, %-- rev-evh: from which the oxygen abundance is rediscussed, and new abundance estimates are derived for N and Ne. The O abundance in SagDIG, estimated with the empirical abundance indicator $\remp$ and other methods, is in the range \( 12+\log ({O/H})=\dodoh \) to \( \dodmax \). %--\( 12+\log ({O/H})=\dodoh \pm \dodoherrtot \). %-- , which makes this dwarf %-- galaxy the second most metal-poor galaxy known after I\,Zw\,18. The fact that SagDIG is \( \sim 10 \) times closer than I\,Zw\,18 makes it an ideal target to test the hypothesis of the existence of young galaxies in the present-day universe. Indeed, stellar photometry suggests that this galaxy may harbor a stellar population older than a few Gyr, and possibly an old stellar component as well. The case of SagDIG therefore supports the view that very low chemical abundances can be maintained throughout the life of a dwarf stellar system, even in the presence of multiple star formation episodes. \keywords{galaxies: individual (SagDIG) -- galaxies: abundances -- galaxies: dwarf -- galaxies: ISM -- galaxies: irregular -- Local Group}
uks1927-177.tex

%-- cerca is3 per l'ultimo cambio
%-- cerca is2 per le mie aggiunte/commenti/tagli!
%-- CERCA ``rev-evh'' per i miei pezzi post-referee.
%-- CERCA le XXX per punti da completare.

%-- seguono considerazioni di carattere generale (v. referee). Ti mando
%   il mio lavoro fin qui, ci risentiamo il 7/1/2002. In effetti c'e' 
%   un minimo da discutere che cosa vogliamo concludere ... Ho messo un 
%   cenno a  McGaugh dal tuo MEMO, aggiungerei accenno a Pilyugin.
%
%   is2: vedi mia discussione
%
%-- Nota 0. Il valore di 7.23 e' sempre li', ma e' da ridiscutere.
%
%   is2: lo tengo, vista la mia discussione
%
%-- NOTA 1. Sarebbe bello in effetti fare la Fig.2 usando soltanto 
%           dati omogenei (low-metal Hii regions e dIrr); penso sarebbe
%           abbastanza veloce, sia tu che Luca avete gia' i dati. 
%           
% is2: beh, ho tagliato la parte metal rich, e aggiunto le regioni di
%      McGaugh (che sono un soprainsieme di quelle di Skillman)
%
%-- Nota 2. La questione del basso U. Nella nostra fig.2 si vede 
%           [come dice il testo] che la relaz. di Skillman corrisp.
%           a logU=-2. In base a McGaugh SagDIg ha logU circa -2.8. 
%           Per 12+log(/H) ritrovo in effetti 7.5 usando a spanne
%           le figure di McGaugh. In particolare il rapporto oII/oIII
%           vincola logU, e quindi la Te.
%
% is2: vedi mia discussione, concludo che non e' affatto chiaro se
%      il logU teorico, e i suoi equivalenti, abbia senso!
%
%           POCO MALE ANCHE SE RITROVIAMO per O/H un valore piu' alto
%           di 7.23, e simile al precedente. La metodologia 
%           e' piu' affidabile. Non vedo motivi religiosi per una  
%           temperatura di 20.000 gradi, che e' supportata solo dai
%           risultati one-zone di Nebular. La tabella calcolata dal
%           referee, che usa un modello a due o piu' zone, 
%           da' dei valori un po' diversi per \otwo. Skillman K. H. 1989 
%           fanno rif. alla relaz. di Campbell Terlevich Melnick 1986.
%           Inoltre loro usano Ne=100 anziche' 1000. 
%
% is2: la Ne influisce ben poco!
%
%-- Nota 3. La ``contraddizione'' nel tuo memorandum (fig. 2 vs. 
%           fig.3 del memo). La fig.3 (Skillman) coincide con la 
%           nostra Fig.2. Si tratta quindi di vedere come i tuoi run
%           di Cloudy riproducono la Fig. di McGaugh. Secondo me si 
%           tratta di diversita' nei modelli. In effetti una figura
%           tipo la Fig.10 di Mc Gaug, rifatta con i conti di Cloudy,
%           e con sopra in nuovi dati, potrebbe non star male nel paper. 
%
% is2: la contraddizione e' anche INTERNA a McGaugh, e vedi anche mia
%      discussione di Pilyugin. Una figura tipo 10 di McGaugh sarebbe
%      buona, in quanto farebbe vedere le difficolta' teoriche di
%      interpretazione. Nevertheless, non ho qui i file originali 
%      prodotti con cloudy, e comunque penso di aver scritto abbastanza
%      per risultare convincente.
%
%-- Nota 4. E' possibile ricavare log (N/O) totale e fare un minimo di 
%           confronto con le altre galassie ? Vedi ad es. Larsen et al. 
%           2001, MNRAS 323, 555. 
%
% is2: trovata seguendo Peimbert and Costero, ma anche quel metodo ha
%      diverse incertezze, quindi quelle quattro parole che ci sono
%      nell'articolo mi sembrano sufficienti.
%
%-- Nota 5. Luca ha misurato la vel. radiale con r.m.s.; e' da scrivere
%           da qualche parte.
%
% is2: pensateci voi, se proprio volete! ma non si era detto che non
%      si riusciva a misurare, oppure con un errore enorme?
%
%-- Nota 6. Idem dicasi per la misura della riga dell'He; visto che e' 
%           stata misurata perche' non metterla in tabella ?
%
% is2: pero' mi pare che Luca aveva dei motivi per non crederci...
%
%--QUESTE SONO VECCHIE (MA SEMPRE VALIDE) NOTE:
%- v. abbondanze in LeoA
%- non ci sono commenti su abbondanze di N e Ne ?
%- perche' non calcoliamo log (N/O) ?
%- una perplessita' di fondo riguarda il calcolo della mass-loss:
%    non sarebbe corretto utilizzare un vero modello con outflow ?

%------------------------------------------------- define values ---------
\def\dodoh{7.26}
\def\dodm91{7.35}
\def\dodmax{7.50}
%evh2:
\def\dodoherrtot{0.20}
\def\dodoherr{0.05}
\def\dodohSTM{7.41}
\def\dodohIZw{7.17}

\def\TeHummer{20000}
%% usando Te=20000 si ottene la seguente C 
%%
         \def\cred{0.335}           \def\ecred{0.061}
     \def\credGamma{0.67}       \def\ecredGamma{0.82}
\def\credRegression{0.36}  \def\ecredRegression{0.05}

%% e con tale C, i flussi corretti per reddening sono i seguenti
%%
      \def\oii{1.36}        \def\eoii{0.21}
    \def\hnine{0.10}      \def\ehnine{0.16}
     \def\neii{0.17}       \def\eneii{0.08}
   \def\height{0.16}     \def\eheight{0.16}
 \def\hepsilon{0.22}   \def\ehepsilon{0.16}
   \def\hdelta{0.23}     \def\ehdelta{0.08}
   \def\hgamma{0.41}     \def\ehgamma{0.11}
    \def\hbeta{1.00}      \def\ehbeta{0.01}
\def\oiiishort{0.46}  \def\eoiiishort{0.01}
 \def\oiiilong{1.30}   \def\eoiiilong{0.01}
   \def\halpha{2.74}     \def\ehalpha{0.05}
      \def\nii{0.11}        \def\enii{0.07}

%% con questi flussi, il R23 e corrispondente O/H sono i seguenti
\def\rduetre{3.12}
\def\erduetre{0.26}
\def\rduetremin{0.46}
\def\rduetremax{0.53}

\def\dodOHrMin{7.19} 
\def\dodOHrMax{7.28}

%%% ed infine, con Te=20000, Ne=1000, ecco le altre abbondanze
%%%
%  \def\oiiHii{1.10}      \def\eoiiHii{0.17}
% \def\oiiiHii{0.72}    \def\eoiiiHii{ 0.01} 
% \def\oiioiii{1.82}     \def\eoiioiii{0.18}
%%
% \def\niihii{0.045}     \def\eniihii{0.026}
%\def\heiiihii{0.22}    \def\eheiiihii{0.10}

%%% ( 12 + log(O/H) = 7.26111412 )

%% ed infine, con Te=18000, Ne=1000, ecco le altre abbondanze
%%
  \def\oiiHii{1.09}      \def\eoiiHii{0.17}
 \def\oiiiHii{0.73}    \def\eoiiiHii{ 0.01} 
 \def\oiioiii{1.82}     \def\eoiioiii{0.18}
%
 \def\niihii{0.062}     \def\eniihii{0.039}
\def\heiiihii{0.28}    \def\eheiiihii{0.16}

%% ( 12 + log(O/H) = 7.357186317 )

%% is2: con Teff = 18000
%% seguendo Peimbert and Costero, possiamo trovare le abbondanze totali.
%% vedi anche la discussione in Kunth and Sargent 1983 sulle obiezioni
%% a tale metodo
%%
%%  N/H =   N+/H+ * O/O+
%% Ne/H = Ne++/H+ * O/O++
%%
%% ora in teoria O/H = (O+/H+ + O++/H+) * (He+ + He++)/H+
%%
%% ma dalle loro tab. 5 e 7 si vede che O/H identico a (O+ + O++)/H+
%%
%% per cui noi possiamo scrivere 
%%
%% O/H = (O+ + O++)/H+ = 2.28 +/- 0.23   (vedi qui sopra)
%%
%% allora O/ O+ = [2.28 +/- 0.23] / [1.37+/-0.22] = 1.66 +/- 0.44
%%        O/O++ = [2.28 +/- 0.23] / [0.91+/-0.01] = 2.51 +/- 0.28
%%
%% dunque
%%
%%  N/H = 0.067+/-0.039 * 1.66+/-0.44 = 0.11 +/- 0.09
%% Ne/H = 0.27 +/-0.13  * 2.51+/-0.28 = 0.68 +/- 0.40
%%
%% e 
%%
%% log( N/H) = -6.23 +/- 0.54
%% log(Ne/H) = -5.26 +/- 0.30
%%
%% infine
%%  N/O = 0.048 +/- 0.04
%% Ne/O = 0.30 +/- 0.21
%% 
%% e 
%%
%% log( N/O) = -1.72 +/- 0.69  
%% log(Ne/O) =  -0.67 +/- 0.37


\def\dodNH{6.01}
\def\errdodNH{0.}
\def\dodNeH{6.85}
\def\errdodNeH{0.30}
\def\logNO{-1.36}
\def\errlogNO{1.17}
\def\logNeO{-0.52}
\def\errlogNeO{0.36}

%% is2: con Teff = 20000
%% seguendo Peimbert and Costero, possiamo trovare le abbondanze totali.
%% vedi anche la discussione in Kunth and Sargent 1983 sulle obiezioni
%% a tale metodo
%%
%%  N/H =   N+/H+ * O/O+
%% Ne/H = Ne++/H+ * O/O++
%%
%% ora in teoria O/H = (O+/H+ + O++/H+) * (He+ + He++)/H+
%%
%% ma dalle loro tab. 5 e 7 si vede che O/H identico a (O+ + O++)/H+
%%
%% per cui noi possiamo scrivere 
%%
%% O/H = (O+ + O++)/H+ = 1.82 +/- 0.18  (vedi qui sopra)
%%
%% allora O/ O+ = [1.81+/-0.18] / [1.10+/-0.17] = 1.65 +/- 0.42
%%        O/O++ = [1.81+/-0.18] / [0.72+/-0.01] = 2.51 +/- 0.28
%%
%% dunque
%%
%%  N/H = 0.045+/-0.026 * 1.65+/-0.42 = 0.074 +/- 0.062 x 1e-5
%% Ne/H = 0.22 +/-0.10  * 2.51+/-0.28 = 0.55  +/- 0.31  x 1e-5
%%
%% e 
%%
%% log( N/H) =  -6.39 +/- 0.52  ( 12 + ... = 5.61)
%% log(Ne/H) =  -5.34 +/- 0.27  ( 12 + ... = 6.66)
%%
%% infine
%%  N/O = 0.074+/-0.062 / 1.82+/-0.18 = 0.041 +/- 0.038
%% Ne/O = 0.55 +/-0.31  / 1.81+/-0.18 = 0.30  +/- 0.20
%% 
%% e 
%%
%% log( N/O) =  -1.81 +/- 0.71
%% log(Ne/O) =  -0.65 +/- 0.35

%% Teff = 20000
%\def\dodNH{5.61}
%\def\errdodNH{0.52}
%\def\dodNeH{6.66}
%\def\errdodNeH{0.27}
%\def\logNO{-1.81}
%\def\errlogNO{0.71}
%\def\logNeO{-0.65}
%\def\errlogNeO{0.35}

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

%%\documentclass[referee]{aa} 
\documentclass{aa} 
%%%\usepackage{psfig} 
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%-- personal preprint packages:
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%++++++++++++++++++++++++++++++ start custom Tex macros +++
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              \begin{figure}
              \centerline{\includegraphics[width=8.8cm]{#1}}
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% ---------------------------------- 1) comandi temporanei, operativi etc .
\newcommand{\uks}{UKS\,1927--177}
\newcommand{\efosc}{{\sc efosc2}}
\newcommand{\remp}{R_{ 23}}
% ---------------------------------- 1) comandi temporanei, operativi etc .
%\newcommand{\changed}{\bf}                         %% draft
\newcommand{\changed}{}                         %% draft
\newcommand{\abbrev}[1]{{\bf #1}}     %%  boldface, draft
%\newcommand{\abbrev}[1]{{ #1}}      %%  plain, final
\newcommand{\bib}[1]{{\bf #1}\\}        %%  draft
\newcommand{\tag}[1]{{\bf #1}\\}         %%  draft
\newcommand{\red}[1]{{\color{red} #1}\\}         %%  draft

%++++++++++++++++++++++++++++++ this paper ++++++++
\def\ss{{SagDIG}\ }
\newcommand{\di}{dIrr} 
\newcommand{\psf}{PSF} 
\newcommand{\trgb}{{\sc trgb}} 
\newcommand{\agb}{{\sc AGB}}
\newcommand{\rgb}{{\sc RGB}} 
\newcommand{\allframe}{{\sc allframe}}
%++++++++++++++++++++++++++++++++++++++++++++++
 
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%
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%+++++++++++++++++++++++++++++++++++++++++++++++++++
\begin{document}

\title{New abundance measurements in \uks, 
a very metal-poor galaxy in the Local Group
\thanks{Based on data collected at the European Southern Observatory,
         La Silla, Chile, prop. No. 63.N-0726B}
}

\author{
I. Saviane\inst{1}
\and 
L. Rizzi\inst{2,3}
\and 
E. V. Held\inst{2}
\and 
F. Bresolin\inst{4}
\and 
Y. Momany\inst{3}
}
\offprints {I. Saviane}

%---------------
\institute{
European Southern Observatory, Casilla 19001,
Santiago 19, Chile \\
\email{isaviane@eso.org} 
\and
Osservatorio Astronomico di
Padova, vicolo dell'Osservatorio 5, I-35122 Padova, Italy \\
\email{(rizzi, held)@pd.astro.it}
\and
Dipartimento di Astronomia, Universit\`a di
Padova, vicolo dell'Osservatorio 2, I-35122 Padova, Italy \\
\email{momany@pd.astro.it}
\and
Institute for Astronomy, University of Hawaii, 
2680 Woodlawn Drive, Honolulu, HI 96822 \\
\email{bresolin@ifa.hawaii.edu}
}
\date {Accepted 17 May 2002}
\titlerunning{New abundance measurements in \uks}

%+++++++++++++++++++++++++++++++++++++++++++++++++++
\abstract{
%% is: added 'brightest'
We present new results from optical spectroscopy of the brightest
H{\sc ii} region in the dwarf irregular galaxy \uks\ in Sagittarius
(SagDIG).  From high signal-to-noise spectra, reddening-corrected line
flux ratios have been measured with typical uncertainties of a few
percent,
%-- rev-evh:
from which the oxygen abundance is rediscussed, and new abundance
estimates are derived for N and Ne.
The O abundance in SagDIG, estimated with the empirical abundance
indicator $\remp$ and other methods, is in the range 
\( 12+\log ({\rm O/H})=\dodoh \) to \( \dodmax \).
%--\( 12+\log ({\rm O/H})=\dodoh \pm \dodoherrtot \).
%-- , which makes this dwarf
%-- galaxy the second most metal-poor galaxy known after I\,Zw\,18.  
The fact that SagDIG is \( \sim 10 \) times closer than I\,Zw\,18 makes
it an ideal target to test the hypothesis of the existence of young
galaxies in the present-day universe. Indeed, stellar photometry
suggests that this galaxy may harbor a stellar population older than a
few Gyr, and possibly an old stellar component as well. The case of
SagDIG therefore supports the view that very low chemical abundances
can be maintained throughout the life of a dwarf stellar system, even
in the presence of multiple star formation episodes.
\keywords{galaxies: individual (SagDIG)
-- galaxies: abundances -- galaxies: dwarf -- galaxies: ISM --
galaxies: irregular -- Local Group}
}

\maketitle

%-----------------------------------------------------------------
\section{Introduction}

The quest for very metal-poor galaxies started soon after the Searle and
Sargent (\cite{ssargent72}) discovery that the interstellar medium (ISM)
of I\,Zw\,18, a blue compact dwarf galaxy, has an oxygen abundance that is
\( \sim 1/50 \) 
%-- (or $12 + \log (\rm O/H) = \dodohIZw$) 
of the solar one.
%
% REPLACES:
%The quest for very metal-poor galaxies started soon after the work
%of Searle and Sargent (\cite{ssargent72}; SS72). They discovered
%that the interstellar medium (ISM) of I\,Zw\,18, a blue compact dwarf
%(BCD) galaxy, has an oxygen abundance that is \( \sim 1/50 \) of
%the solar one. 
%
As Searle \& Sargent (\cite{ssargent72}) pointed out, this is compatible
with two scenarios: either the galaxy is now experiencing its first
star-formation (SF) episode, or, if it contains \( \sim 10\, \rm Gyr \)
old stars, then its SF proceeds in strong bursts separated by quiescent
phases. The first hypothesis (that of a \emph{primordial} galaxy) is
relevant in the context of cosmological models that assume hierarchical
formation of structures in a cold dark matter  dominated
Universe. In that scenario, dwarf dark matter halos (\( \sim 10^{8}\, \rm
M_{\odot } \)) would be the
\emph{first} objects 
to condense out of the Hubble flow, so it is important to understand
whether they already contained stars before merging into larger
galaxies. As a first step, we can start by answering a simple
question: is there any galaxy with a very metal-poor young population
and, at the same time, a \( >10\, \rm Gyr \) old generation of stars?
Such a finding would lead to the conclusion that very low metal
abundances do not necessarily imply a very young age.

Although a red population was eventually revealed in I\,Zw\,18 
by Aloisi et al. (\cite{aloisi_etal99}), 
{\changed their data only allow for a look-back time}
of \( \sim 1 \)~Gyr, due to the large distance of the galaxy.
Izotov et al. (\cite{izotov_etal00}) argue that \( 100\, \rm Myr\) 
is a more realistic limit. As a consequence, active searches are
carried out to find closer and comparably metal-poor galaxies, in
order to have better chances to detect a cosmologically old
population. In particular, \emph{low luminosity} dwarf irregulars are
promising candidates. It is still debated whether a definite
luminosity-metallicity relation exists for dIrrs (Hidalgo-Gamez \&
Olofsson \cite{h-go98}; Hunter \& Hoffman \cite{hh99}), but it is
nevertheless reasonable to expect that chemical abundances will be
lower in low-luminosity galaxies.

Indeed, some new metal-poor, low-luminosity objects turned up in the
past as a result of 
%evh2:
similar efforts. The situation has been conveniently
summarized by Kniazev et al. (\cite{kniazev00}), who list two galaxies
(besides I\,Zw\,18) at \( 12+\log (\rm O/H)\leq 7.35 \).  A few months
later van Zee (\cite{vanZee00}) added UGCA~292 to the list of
metal-poor galaxies, quoting \( 12+\log ({\rm O/H})=7.30\pm 0.05 \),
and pointed out that Leo~A should be added to the list as well.

Stimulated by all these promising results, we turned our attention to
\uks\ (SagDIG), since (a) its low luminosity is in the range of
other metal-poor galaxies, 
%evh: ?? (see Table~1 in K00), 
(b) there was some indication that the oxygen abundance obtained by
Skillman et al. 
%-- Skillman  Terlevich, \& Melnick
%evh2: referenza !!!
(\cite{stm89}, hereafter STM89), 
\( 12+\log ({\rm O/H})= \dodohSTM \), could be revised
downwards (see below), and (c) the other metal-poor galaxies (with the
exception of Leo~A) are more distant than \( 3\, \rm Mpc \). SagDIG
is more than two magnitudes closer in distance modulus (thus its HB is
within reach of HST observations in reasonable observing time).  

%% is: aggiungo un rapido sommario sulle regioni Hii
%%
The galaxy was discovered in 1977 by Cesarsky et
al. (\cite{cesarsky_etal77}), and a shallow CMD was presented in Cook
(\cite{cook87}). A CCD survey of its \hii\ regions was carried out by
Strobel et al. (\cite{strobelEtal91}; SHK91), using the
$2.1~\rm m$ telescope at Kitt Peak.  Three objects were detected in a
field of view of $\sim 2\farcm 5 \times 2\farcm 5$, with 
%-- apparent
H$\alpha$ fluxes of $28.1\times 10^{-15}$, $20.0\times 10^{-15}$, and
$4.09\times 10^{-15}~\rm erg\,cm^{-2}\,s^{-1}$.  Optical
spectrophotometry of the brightest object (object \#3 in Table~8 of
SHK91) was obtained by STM89, and an oxygen abundance
\( \sim 1/30 \)  of the 
%
solar value\footnote{{\changed In this paper we adopt the classical
value $12+\log(\rm O/H) = 8.87$ from Anders \& Grevesse 
(\cite{ande+grev89}). However, the solar oxygen abundance is
currently debated, and a lower value might be appropriate (Holweger
\cite{holweger01}; Allende
Prieto et al. \cite{alle+01}).  }}
%
could be determined, i.e. $12+\log(\rm O/H) \simeq \dodohSTM $. 
%
Radio maps of the galaxy exist as well. Young \& Lo (\cite{youngLo97})
presented deep VLA observations of the H{\sc i} in SagDIG. An amount
of \( 1.3\times 10^{7} \) solar masses of H+He were detected, forming
a circular ring which extends \( \sim 3 \) times farther than the 
optical extent of the star-forming regions.

The compactness of the galaxy, its low Galactic latitude,
%evh: NO (with consequent large extinction) 
and the high foreground contamination had discouraged further broadband
photometry until recently. In the last two years, imaging of this galaxy
was presented by 
Karachentsev et al. (\cite{kara_etal99}) and 
Lee \& Kim (\cite{lee_kim00}), who established the distance modulus at
$(m-M)_0=25.13\pm 0.25$, and the RGB stellar metallicity at extremely
low values ([Fe/H]$\la -2.45$).
From surface photometry, Karachentsev et
al. (\cite{kara_etal99}) found an integrated magnitude of
$M_V=-11.74$, and with a synthetic CMD analysis, they concluded that
the galaxy is currently experiencing a burst of star formation, the
star formation rate (SFR) being $10$ times larger now than the life
average.

A new imaging study of SagDIG is presented in a companion paper
(Momany et al. \cite{yaz_sagdig}). The most important results of our
photometry is represented by an upward revision of the stellar
metallicity of the galaxy and our analysis of the spatial distribution
of different stellar populations as a function of age.  Assuming a low
reddening for old RGB stars, we derived [Fe/H]$=-2.08\pm 0.20$ for the
red giant branch stars. 

%rev-evh:
In this paper we present our new determination of the oxygen,
nitrogen, and neon abundances in the interstellar medium of 
\uks, and discuss it in the context
of the existence of primordial galaxies.

%\resizebox{1\columnwidth}{!}{\includegraphics{righe.ps}} 
%\figcaption[FileName]{Line identification for grism \#7 
% and \#9\label{fig:idgr7+9}}


%---------------------------------------------------------------
\section{Observations and data reduction}

%-------------------------------------------------->
%%%% FOR EMULATE APJ
\realfigure{righe.ps}{
The combined spectra of the brightest \hii\ region of SagDIG (object
No.3 in Strobel et al. \cite{strobelEtal91}).  The upper panel shows
the line identification for the averaged flux-calibrated spectrum
obtained with grism \#7.  The lower panel shows the same except for
grism \#9.
}{f_idgr7+9}
%--------------------------------------------------|

Spectroscopic observations of SagDIG were carried out on 1999 August
10, at the ESO observatory in La Silla, Chile. The imaging
spectrograph \efosc\ attached to the ESO 3.6m telescope was
used to secure low-resolution spectroscopy of SagDIG. The camera
was equipped with a Loral/Lesser thinned, AR coated, UV flooded, \(
2048^{2} \) pixel CCD (ESO CCD \( \#40 \)).  The \( 15\mu \) pixels
cover \( 0\farcs 157 \) on the sky, with a resulting field of view of
\( 5\farcm 4\times 5\farcm 4 \).
%evh
Two grisms were employed to cover the wavelength ranges \( 3270\div
5240 \)~\AA\ (grism \#7) and \( 4700\div 6770 \)~\AA\ (grism \#9),
respectively. The CCD was used in \( 2\times 2 \) binning mode to
increase the \( S/N \) ratio, the resulting dispersion being \( \sim 2
\)~\AA/pixel for both grisms.
%evh:
%The ESO FIERA controller was used in `normal' readout mode, which yields a
%readout noise of \( 7.5 \)~\( e^{-} \) and a conversion factor of \(
%1.3 \)~\( e^{-}/ \)ADU.

%% is: aggiungo quale oggetto e' stato osservato
%%
We observed the same \hii\ region studied by STM89, whose coordinates
are $\alpha=19$:$30$:$02.5$, $\delta=-17$:$41$:$26.1$ (J2000.0).
%%
The night was stable and photometric, with a fairly constant seeing of
\( \sim 1\arcsec \) (FWHM) at H\( \alpha \). Spectroscopic data
comprise exposures of \( 6\times 1800 \) s with grism \#7 and \( 4\times
1800 \) s using grism \#9. In both cases, a slit width of \( 1\farcs 5 \)
was used. 
%\changed: no signal was detected for Object 2 ! Commented out.
%, and the slit was rotated so that spectra of object \( 2
%\) of SHK91 could also be taken.
%
Spectrophotometric standards were observed during the night, selected
from the catalogs of Oke (\cite{oke90}) and Hamuy et al. (\cite{hamuy92},
\cite{hamuy94}).

Data reduction was performed using standard techniques in {\sc
iraf}\footnote{ IRAF is distributed by the National Optical Astronomy
Observatories, which are operated by the Association of Universities
for Research in Astronomy, Inc., under cooperative agreement with the
National Science Foundation.}.
%
The raw two-dimensional images were bias subtracted and trimmed.
Pixel-to-pixel variations were accounted for through division by 
a quartz lamp spectrum, after fitting a low-order polynomial along
the dispersion.  The splitting of the total exposure time in
%evh:
\( 1800 \) s exposures allowed us to remove the
contamination from cosmic rays.
%
Wavelength calibration was obtained from He-Ar exposures taken
immediately before and after the object observations.  Absolute flux
calibration was achieved using standard stars, obtaining a final
accuracy better than \( 5\% \). 
%
One-dimensional spectra were extracted from the two-dimensional images
by summing up all the pixels enclosed in the large profiles of
H\({\alpha } \) and H\( {\beta } \); the background and sky lines
were subtracted by averaging a small region on either side of the
extraction window.


%-evh: uncomment here for referee format:
%\placetable{t_ratios}
%-evh: comment the table below for referee format:
%------------------------------------------------------------------->
%-- {\footnotesize
\tablecaption{Line flux ratios relative to $\mbox{H}{\beta}$ \label{t_ratios}}
% \tablerefs{1) very nice  reference 1) very nice  reference }
% \tablecomments{real comment this is a real comment this is a real comment}
%--- NOTE NON-STANDARD (NON-AAS) USAGE OF TABLENOTETEXT ---
\tablenotetext{
   \begin{list}{}{}
     \item[$^{\mathrm{a}}$] {possible blend with 3968 [\otre].}
     \item[$^{\mathrm{b}}$] {estimated maximum error (0.15) (see text).}
   \end{list}
}
\begin{deluxetable}{cllcc}
%\tablewidth{0pt}
\tablehead{
\colhead{\(\lambda\)(\AA)} &
\colhead{Ident.} &
\colhead{Grism \#7} &
\colhead{Grism \#9} &
\colhead{Redd. corrected}
}
\startdata
3727&
 [\otwo] &
1.02\( \pm  \)0.20&
\nodata&
1.31\( \pm  \)0.21\\
3835&
H$_9$ &
0.08\( \pm  \)0.15\tablenotemark{b} &
\nodata&
0.10\( \pm  \)0.16\\
3868&
[\netre] &
0.13\( \pm  \)0.05&
\nodata&
0.16\( \pm  \)0.08\\
3889&
H$_8$ &
0.12\( \pm  \)0.15\tablenotemark{b}&
\nodata&
0.15\( \pm  \)0.16\\
3970&
H${\epsilon }$\tablenotemark{a}&
0.17\( \pm  \)0.15\tablenotemark{b}&
\nodata&
0.21\( \pm  \)0.16\\
4101&
H${\delta}$ &
0.19\( \pm  \)0.07&
\nodata&
0.23\( \pm  \)0.08\\
4340&
H${\gamma}$ &
0.36\( \pm  \)0.11&
\nodata&
 0.41\( \pm  \)0.11\\
4861&
H${\beta}$ &
1.00\( \pm  \)0.04&
1.00\( \pm  \)0.03&
1.00\( \pm  \)0.01\\
4959&
[\otre] &
0.48\( \pm  \)0.06&
0.46\( \pm  \)0.04&
0.46\( \pm  \)0.01\\
5007&
[\otre] &
1.35\( \pm  \)0.10&
1.34\( \pm  \)0.02&
1.30\( \pm  \)0.01\\
6563&
H${\alpha }$ &
\phantom{000}\nodata&
3.58\( \pm  \)0.03&
2.85\( \pm  \)0.05\\
6584&
[\ntwo] &
\phantom{000}\nodata&
0.15\( \pm  \)0.06&
0.12\( \pm  \)0.07 \\
%
\end{deluxetable}
%\tablenotetext{a}{possible blend with 3968 [\otre]}
%\tablenotetext{b}{estimated maximum error (0.15) (see text).}
%-- }

%------------------------------------------------|

Figure~\ref{f_idgr7+9} shows the combined spectra used for measurements.
The final resolution,
measured either on sky lines or comparison lamp lines, is \( 8.3 \)~\AA\
for grism \( \#7 \) and \( 9 \)~\AA\ for grism \( \#9 \).
%
Table \ref{t_ratios} presents the final results. From left to right,
the columns report the line wavelength and identification, the 
H\({\beta } \)-normalized fluxes for grisms \#7 and \#9, and the
reddening corrected fluxes (see below). 
%% is: cambio 
Three emission lines ([\otre] $\lambda\lambda 4959, 5007$ and H${\beta}$)
are in common between the two grisms, showing no significant flux
offset.
%evh2: 
In the worst case, that of the {[}\otre{]} \( \lambda 5007 \) line, 
the agreement is better than \( 5\% \).
%
A deblending procedure was applied to the case of the {[}\ntwo{]} \(
\lambda 6584 \) line to disentangle the contribution of the strong 
H\({\alpha } \).
%he result is presented in Figure \ref{fig:deblend}. 
%
%The error budget was estimated extracting the single exposures in
%the same way as the averaged resulting spectrum. Line fluxes were
%then estimated for each of the extraction and their rms dispersion
%estimated. The error of a particular line ratio was then obtained
%summing in quadrature the error of the line and the H\( {\beta } \)
%error. 
The error budget was estimated from the r.m.s. of line ratios measured
on the different individual exposures.  A maximum error of 15\% on the
flux ratio was attached to our estimate when the line is not detected
in the single-exposure spectra.


%%%%%%%%%%%%%%%%%%%%%%%%%%%% REDDENING CORRECTION

The measured line ratios were corrected for the effects of 
%--evh2: differential absorption, 
reddening, measured by comparing the observed and the expected
hydrogen line ratios.
%
%The 
%ratios H\( {\alpha } \)/H\( {\beta } \), H\( {\beta } \)/H\( {\gamma } \)
%etc. are commonly used, since they are relatively independent from
%\( T \), and since the large fluxes involved allow accurate measurements.
%
The theoretical H line ratios of Hummer \& Storey
(\cite{hummer}) were used, assuming an electron temperature 
\( T_{\rm e}\simeq \TeHummer \) K.
%% is: questo non serve piu' 
%
%, which applies to most HII cases (see Osterbrock \cite{agn2}; his
%Table~4.2).
%
We set 
%
\(R=I_{\lambda }/I_{\rm H\beta } \), 
\( R_{0}=I_{\lambda, 0}/I_{\rm H\beta, 0} \) and 
\( \varphi (\lambda )=[f(\lambda )-f(H\beta )] \) so that 
\( C=[\log R_{0}/R]/\varphi (\lambda ) \).  
%
The reddening constant was then computed using the 
Cardelli et al. (\cite{cardelli}) extinction law,
% is2: aggiunto definizione di f(lambda)
%
where $f(\lambda)=<A(\lambda)/A(V)>$, and we assumed $R_V=3.1$.
 We obtained 
\( C=\cred\pm \ecred \) from the H\( {\alpha } \)/H\( {\beta } \) ratio, and
\( C = \credGamma \pm \ecredGamma \) from the 
H\( {\beta } \)/H\( {\gamma } \) ratio. 
%
As an alternate route, one can plot \( \log (R_{0}/R) \) vs. \(
\varphi (\lambda ) \) and then find \( C \) with a linear regression
on the data. A weighted least-square fit on the first three H lines
yields 
%
\( C= \credRegression  \pm \ecredRegression \).  
%
Thus the value of the reddening constant is
essentially fixed by the more accurate 
H\( {\alpha } \)/H\({\beta } \) ratio, 
and \( C=\cred\pm \ecred \) will be adopted in the rest of this paper.
%
The reddening corrected flux ratios are listed in
Table~\ref{t_ratios}. When the same line is present in both the red
and blue spectra, a weighted mean is computed before the correction is
applied.  In order to compute the final errors, the error contribution
from the reddening correction
%-- evh2: exponential part 
was added in quadrature to the error on the flux.

%---------------------------------------------------------
\section{Results}

%evh:
Chemical abundances were derived from these measurements using the {\sc
iraf} \texttt{nebular} package (Shaw \& Dufour \cite{sd95}).
%--Since limited spectral information is available, 
%-- evh2: 
The abundances were calculated assuming a constant electron temperature
and density throughout the nebula.
%
For optical spectra covering the observed wavelength range,
%% is: cambiato qui 
the electron temperature should be obtained through $T_{\rm e}$-sensitive
line ratios, e.g. from
%evh2:
the ratio of {[}\otre{]} lines \( R_{4363}=[I(4959)+I(5007)]/I(4363) \)
(see Eq.~2.3 in Seaton \cite{seaton75}). However, the \( \lambda 4363
\) line was not detected even in our co-added spectra, so we 
%-- rev-evh:
%is2: includo un po' di discussione del metodo empirico, prima di usarlo
%per trovare la O/H
had to use the so-called empirical method, which gives the total
(O/H) abundance in terms of the \( \remp =
[I(4959)+I(5007)+I(3727)]/I(H\beta ) \) parameter (Pagel et
al. \cite{pagel_etal79}; Edmunds \& Pagel
\cite{edm_pagel84}). 

%%% FOR EMULATE APJ
%-------------------------------------------------->
\realfigure{r23.ps}{
(O/H) determination for SagDIG using the empirical method based on
$\remp$. The $1\sigma$ range of variation of our value of $ \log(\remp)$
is shown by the vertical band. The empirical relation of Skillman
(\cite{skillman89}) is represented by the inclined heavy short-dashed
line.  The horizontal solid lines show the range of uncertainty of $\log
\rm (O/H)$.
%
For comparison, 
values of \( \log \rm (O/H) \) are plotted against the value
of \( \log \remp \) for the sample of \hii\ regions from 
D{\'{\i}}az \& 
P{\' e}rez-Montero (\cite{dpm00}) ({\it open circles}), Oey
\& Shields (\cite{oey00}) ({\it filled diamonds}),  Kennicutt et
al. (\cite{kennicuttEtal00}) ({\it open triangles}), and McGaugh
(\cite{mcga91}) ({\it filled circles}). Only objects more metal poor
than $12+\log(\rm O/H)=8.3$ are plotted.
%
%
A number of {\sc cloudy} simulations are also shown.
The thin solid curves represent the 
{\changed theoretical} relations obtained by
varying the ionization parameter from $\log U = -1 $ to $\log U = -4$
(see labels), and by varying [m/H] from $-0.4$ to $-2.4$ 
(top to bottom, {\it dotted lines}). An input ionizing spectrum from a 
\( T_{\rm eff}=40\,000\, \rm K \) star was assumed. The effect of
changing the input spectrum is shown by the long-dashed curves, 
which were calculated for \( T_{\rm eff}=50\,000\, \rm K \).
}{f_emp_meth}
%--------------------------------------------------|

%is2:
Oxygen abundances are plotted against $\remp$ in
Figure~\ref{f_emp_meth}, using literature samples of \hii\ regions
spanning a wide range of abundances (Oey \& Shields \cite{oey00}; 
D{\'{\i}}az \& P{\' e}rez-Montero \cite{dpm00};
Kennicutt et al. \cite{kennicuttEtal00};
McGaugh \cite{mcga91}). 
%
In the same diagram, we show a set of {\sc cloudy} (Ferland et
al. \cite{ferlandEtal98}) simulations, obtained assuming that the
ionizing source is a single star with a Mihalas (\cite{mihalas72})
continuum, for \( 4 \) choices of the ionization parameter (\( \log
U=-1,\, -2,\, -3 \) and \( -4 \)) and varying the metallicity from
$12+\log \rm (O/H)=6.47$ to $12 + \log\rm (O/H)=8.87$ (solar composition).
%
The effects of a change in the temperature of the
ionizing star were checked by running the simulations for both  
\( T_{\rm eff}=40\,000\, \rm K \) and 
\( T_{\rm eff}=50\,000\, \rm K \). 

This figure, an updated version of Fig.~2 of Skillman (\cite{skillman89};
hereafter S89) and Fig.~12 of McGaugh (\cite{mcga91}),
shows that, for abundances $12+\log(\rm O/H)\leq 8$,  
$\remp$ depends both on the metallicity and the ionization
parameter of the nebula.
% \[
% {U}=\frac{3\, Q}{4\, \pi \, c\, N\, R^{2}_{\rm S}}\] where \( Q \) is
% the ionizing luminosity, \( N \) is the gas density, and \( R_{\rm S} \)
% the Stromgrev radius.  Compact, high surface brightness regions should
% therefore have a larger ionization parameter. 
%
%--------------------------------------------
The S89 relation corresponds to
the lower branch of the \( \log U=-2 \) models
%--rev-evh: referee:
% is2: commento qui, visto che in effetti sono nebulae con varie
%      excitations
%(this is consistent with the fact that the S89 calibration is based
%primarily on high excitation nebulae)
%
and to the most metal-poor objects in the data samples,
irrespective of the effective temperature of the ionizing source.
%--rev-evh: referee:
% The high-abundance solution for $\remp$ (the so-called ``upper branch'')
% can be safely ruled out because the low luminosity and large gas content
% of the galaxy point to a chemically unevolved ISM.  Therefore our
%
% OLD PIECE.The double-valued estimates of log\,(O/H) which are 
% obtained from comparison of the
% $\remp$ with the models would make our line measurements compatible
% with both $12+ \log \rm (O/H) \sim 8.5$ and 
% $12 + \log \rm (O/H)= \dodoh $.
% However, 
% the former can be safely ruled out because the low luminosity and
% large gas content of the galaxy point to a chemically unevolved ISM.
% Our adopted value for the oxygen abundance in SagDIG is then
%
%\( ({\rm O/H})=(1.71\pm 0.16)\times 10^{-5} \), i.e. 
%
%The measured value for the oxygen abundance in SagDIG using the
%empirical method is 
%\( 12+\log (\rm {O/H})=\dodoh \pm \dodoherrtot \).
%The quoted error, larger than the formal error derived from
%application of the Skillman's (\cite{skillman89}) relation ($\pm
%\dodoherr$), takes into account the inherent uncertainty of the
%empirical method.
%
%evh2:
% Note that we would have obtained this value for (O/H) from the
% data in Table~\ref{t_ratios} using the \texttt{nebular}
% photoionization code calculations, had we assumed an electron 
% temperature $T_{\rm e}=20\,000$. 
%
%% is: aggiungere commento da Skillman (evh2: ???), 
% we are confident in our lower value
% of \( (\rm {O/H}) \), on which the rest of the discussion is based.
%evh:
%We first checked the correctness of
%our procedures, and recalculated the oxygen abundance using the line
%fluxes quoted in STM89 (their blue spectrum), and we found \( (\rm
%{O/H})=2.2\pm 0.8 \), which is almost coincident with their value of
%\( 2.3 \). 
%--------------------------------------------

Several authors have attempted to include both the effects of 
metallicity and ionization parameter 
in their empirical calibration methods. 
According to McGaugh (\cite{mcga91}), 
the parameter \( \log (\rm [\otre]4959,5007/[\otwo]3727) \)
is useful to constrain $U$. 
%
%In the case of SagDIG, we have $\log (\rm
%[\otre]4959,5007/[\otwo]3727)\simeq0.13$, which implies $\log U=-3$. 
%In turn, using this value of $\log U$ in our Fig.~\ref{f_emp_meth}, we
%would obtain $12+\log(\rm O/H)\simeq 7.5$.
%
%However, the predictions of the models are not totally
%self-consistent, as we now show. The five objects that have $12+\log
%(O/H) \leq 7.55$ in the M91 sample are all close to $\log U=-2$. On
%the other hand, their position in the $\log (\rm
%[\otre]4959,5007/[\otwo]3727)$ vs. $\remp $ plane implies $U$ varying by
%an order of magnitude, from $\log U = -3$ to $\log U = -2$. For
%example, the galaxy ESO\,380-27 (Tol~1223-359), with $\log (\remp) =
%0.75$, has an excitation parameter $\log (\rm
%[\otre]4959,5007/[\otwo]3727)=0.68$ (implying $\log U \sim -2$), and an
%oxygen abundance $12+\log(\rm O/H)=7.53$. With an excitation parameter
%ten times smaller ($\log (\rm [\otre]4959,5007/[\otwo]3727)=0.06$,
%implying $\log U \sim -3$), GR\,8 has a similar abundance $12+\log(\rm
%O/H)=7.43$ at $\log (\remp) = 0.70$.
%
% In summary, M91 turned the $\log(\rm O/H)=f(\remp,\log U)$ relation
% into a $\log(\rm O/H)=g(\remp,\log (\rm [\otre]4959,5007/[\otwo]3727)$
% relation, so the oxygen abundance of an
% \hii\ region is fixed by its position on the surface of equation $g$. 
% One has to use his Figg.~10 and 12 to find this position.  In a similar
% way, 
More recently, Pilyugin (\cite{pilyugin00}, \cite{pilyugin01a},
\cite{pilyugin01b}) has provided useful formulae to implement McGaugh
corrections into the empirical method. 
%
Pilyugin (\cite{pilyugin01b}; hereafter P01) transformed the
$\log(\rm O/H)=f(\remp,\log U)$ relation into a $\log(\rm
O/H)=h(\remp,R_3)$ relation, where 
$R_3=\log (\rm [\otre]4959,5007/H\beta$). In this case he gives the
equation of the surface, which is:
\[
12 + \log (\rm O/H) =  6.35 + 3.19 \times \log(\remp) - 1.74 \times \log(R_3)
\]
% 1.45 lgR3 - 3.19 (lgR3 - lgR23) = -1.74 lgR3 + 3.19 lgR23
These revised methods may have problems at low values of 
the abundance and the ionization parameter.
For example, the dwarf galaxies ESO\,380-27 and GR\,8 have similar $\log
(\remp)$, but $\log (R_3)=0.67$ for the former, and $\log
(R_3)=0.43$ for the latter. The low excitation parameter then implies
that GR\,8 should have a lower $\log U$, implying a higher oxygen
abundance. Indeed, we find $12+\log(\rm O/H)=7.58$ for the first, and
$12+\log(\rm O/H)=7.83$ for the second galaxy. The corrected empirical
method would then predict that GR\,8 should have an oxygen abundance
$0.25~\rm dex$ higher than ESO\,380-27, but in reality the 
measured abundance from
the temperature determinations is $0.10~\rm dex$ {\em lower} in GR\,8
than in ESO\,380-27 (see Kunth \& Sargent \cite{kunthSargent83}; 
Skillman et al. \cite{skillmanEtal88}).
%
% results of ~/papers/Spectroscopy/mg.mac -> tp
%#    lgOH    lg\otwo   lg\otre    lgR23 dodLogOH  trueDod        name
%#
%    -4.45     0.31     0.54     0.74     7.77     7.55 A1116+51
%    -4.47    -0.01     0.67     0.75     7.58     7.53 Tol_1223-359
%    -4.48     0.29     0.48     0.70     7.75     7.52 Sextans_A
%    -4.57     0.37     0.43     0.70     7.83     7.43 Gr_8
%    -4.76    -0.48     0.43     0.48     7.13     7.24 I_Zw_18
%
% Another puzzling behaviour of the P01 method is revealed as follows. If
% we use the line fluxes of the STM89 blue spectrum, SMT89 blue+red
% spectra, or our line fluxes, we obtain oxygen abundances $12+\log(\rm
% O/H)=7.48\pm 1.65$, $7.36 \pm+ 1.12$, and $7.47 \pm 0.13$,
% respectively. It therefore seems like the method gives almost the same
% oxygen abundance regardless of (very different) input fluxes. Indeed,
% any two \hii\ regions, A and B, having 
% \[
%      \frac{\log R_3^{\rm A}-\log R_3^{\rm B}}{\log P^{\rm A}-\log P^{\rm B}} = 2.2
% \]
% %
% will give the same oxygen abundance regardless of the input line
% fluxes. 
%
%-------------------------------------------------->
%%%% FOR EMULATE APJ
\realfigure{mcgaugh91.ps}{
A comparison of the observed values of \( \remp\ \) and the
[\otre]/[\otwo] ratio for the most metal-poor star-forming dwarfs with
a model grid obtained from {\sc cloudy} simulations, for $T_{\rm
eff}=50000~\rm K$.  The data are from Skillman (\cite{skillman89}), 
McGaugh (\cite{mcga91}), Van Zee (\cite{vzee+99}), and Kniazev et al. (\cite{kniazev00}).}{f_mcgaugh91}
%--------------------------------------------------|
%
Given all these uncertainties,  
we used both the original implementation of the empirical method
(as in STM89) and more recent methods to provide a range of
[O/H] for SagDIG. 
% is2: aggiunto qui un caveat
High $S/N$ spectra taken at 10-m class telescopes will be needed to 
place better constraints on the SagDIG abundance, by allowing the
measurement of the $\lambda 4363$ [\otre] line.
%

Therefore, we firstly used the relative O fluxes 
listed in Table~\ref{t_ratios} to obtain
\( \remp = \rduetre \pm \erduetre \) 
(i.e. \( \rduetremin \leq \log \remp \leq \rduetremax \)). 
The calibration of S89 
%% is: elimino citazione della formula, e metto 12 + \log (O/H)
%is \( 12+\log (\rm O/H)=1.21\, \log R_{23}+6.64 \), which
%yields \( -4.81\leq \log (\rm O/H)\leq -4.73 \), where the range in
%
gives \( \dodOHrMin \leq 12 + \log (\rm O/H)\leq \dodOHrMax \), where
the range in variation has been obtained by propagating the error 
on \( \remp \).
An application of the empirical calibration therefore indicates a
\emph{very low oxygen abundance} for the Sagittarius dwarf irregular
galaxy.  This value can be directly compared to the abundance
obtained by STM89 using the same calibration.  The new abundance
estimate is lower by $\sim0.2$ dex, and the difference is real and
traced to the new flux measurements.  Since our measurements from
the red and blue \efosc\ spectra are perfectly consistent, whereas the
measurements of STM89 presented a \( \sim 40\% \) internal scatter, we
are confident in our revised fluxes (indeed, using only the STM89 {\it
red} spectrum for the
\( \lambda \lambda 4959,5007 \) fluxes, we would obtain a value
\( 12 + \log (\rm {O/H})=7.22 \), in perfect agreement with ours). 

%--2nd method:
Secondly, the revisions of the empirical method were taken into
account by plotting in Fig.~\ref{f_mcgaugh91} the [\otre]/[\otwo]
ratio against the \( \remp\ \) for a subsample of very metal-poor
dwarfs (see Fig.~10 in McGaugh \cite{mcga91}).  The data are
overplotted onto a model grid derived from the same {\sc cloudy}
models as in Fig.~\ref{f_emp_meth}.
%
In the case of SagDIG, we have $\log (\rm
[\otre]4959,5007/[\otwo]3727)\simeq0.13$, which implies $\log U \approx -3$. 
%In turn, using this value of $\log U$ in our Fig.~\ref{f_emp_meth}, we
%would obtain $12+\log(\rm O/H)\simeq 7.5$.
%
The position of SagDIG with respect to the lines of constant O
abundance indicates a metallicity intermediate between those of
I\,Zw\,18 and GR\,8, and close to that of HS0822+3542, for which
Kniazev et al. (\cite{kniazev00}) give $12+\log (\rm O/H) = 7.35$.  Thus, taking into account
{\it in a differential way} the information provided by the ionization
indicator [\otre]/[\otwo], we find $12+\log (\rm O/H) \simeq \dodm91$.

Finally, we evaluated the O abundance by using {\sc cloudy} to
optimize the fit to the observed line ratios. The best-fit is obtained
for a value $12+\log (\rm O/H) = 7.41$, i.e. 1/29 solar, with an
electron temperature $T_e \sim 16\,000$ K and $\log U \simeq -2.8$.
Not surprisingly, this result is consistent with that inferred from
Fig.~\ref{f_mcgaugh91}.  
%These different approaches led us to adopt for SagDIG an O abundance
%\( 12+\log ({\rm O/H})=\dodm91 \pm 0.20\), where 
%the error takes into account the inherent uncertainty of the 
%empirical methods. 


%Mihalas, D. 1972, Non-LTE Model Atmospheres for B \& O Stars, NCAR-TN/STR-76
%Mihalas, D. 1978, Stellar Atmospheres, 2nd Edition (W.H. Freeman: San Francisco)

%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%
%------------ TEMPERATURA E DENSITA` ELETTRONICA ------------
%
% is: cambiato secondo FB
% What \( T_{\rm eff} \) is compatible with this (O/H)? For \( T=20000\,
% \rm K \) and \( N_{e}=1000\, \rm cm^{-3} \) we obtain \( N(\rm
% {O^{+}})/N(\rm H^{+})=(1.05\pm 0.17)\times 10^{-5} \) and \( N(\rm
% {O^{++}})/N(\rm H^{+})=(0.72\pm 0.01)\times 10^{-5} \), so \( (\rm
% {O^{+}}+\rm {O^{++}})/\rm {H^{+}}=(1.77\pm 0.18)\times 10^{-5}<\rm
% (O/H) \).  The single ionized oxygen abundance was calculated from the

%-- rev-evh: QUESTO E' IL PUNTO CENTRALE DEL PROBLEMA ...
To proceed further to estimate the N and Ne abundances in SagDIG, we
need making some considerations (and assumptions) on the electron
temperature and density in the nebula. 
% is2: commento questo, meglio non sollevare dubbi!
%In fact, the doubt may arise that
%the physical assumption on which the $\remp$\ calibration rests are not
%applicable to the SagDIG \hii\ region.
%--rev-evh: esperimenti di Luca.
A loose upper limit on the electron temperature is set by non-detection
of the \( \lambda 4363 \) {[}\otre{]} line.  Experiments simulating faint
emission lines yield a $2\sigma$ upper limit of \( 0.09 \) on the flux
of the \(\lambda 4363 \) line relative to H$\beta$, which implies
$T_{\rm e} < 34\,000$ K according to the Seaton's (\cite{seaton75})
relation.
%-- tengo conto del referaggio:
% is2: low surface brightness lo dice il referee, noi possiamo dire
%      che e' low excitation, ma come detto nella discussione
%      precedente, sembra che questo non vincoli molto il logU
%
%      Dunque commento questo pezzo, e del resto altrimenti c'e' 
%      inconsistenza nelle abbondanze di O e di tutto il resto
%
%This upper limit is not very useful, however. In such a low surface
%brightness \hii\ region, the ionization parameter and electron
%temperature are unlikely to be higher than in I\,Zw\,18, where the
%temperature in the O$^{++}$ zone is between 18\,000 and 20\,000 K
%(QUI-REFS-XXX).
%--rev-evh: qui implemento le tue considerazioni sul metodo empirico. 
% is2: le considerazioni le ho fatte io prima, e quindi commento pure
%      questo
%Some information on the most likely value of the ionization parameter
%(and $T_{\rm e}$) is provided by the ratio of the {[}\otre{]} to
%{[}\otwo{]} lines, which is a function of $U$ for any given [O/H] (see,
%e.g., Fig.~10 in Mc~Gaugh \cite{mcga91}).
%%
%A comparison of the measured values of \( \log(\remp) \) and
%\( \log ([I(4959)+I(5007)]/I(3727))=0.128 \pm XXX \) 
%with a grid of photoionization models in the low-metallicity range,
%%-- log(O/H)$=-5.1$ to $-3.9$, 
%yields \( \log U=-2.8 \), \( 12+\log ({\rm O/H})=7.5 \). This value of
%\( \log U \) corresponds to $T_{\rm e} \approx 16\,000$ K.
%
A temperature $T_{\rm e}=20\,000 ~ \rm K$ is needed in the
\texttt{nebular} calculations to obtain an (O/H) value compatible with
the results from the empirical method (assuming \( N_{\rm e}=1000\ \rm
cm^{-3} \)).
%
In fact, using these values we obtain
%
\(N(\rm {O^{+}})/N(\rm H^{+}) = ( \oiiHii  \pm \eoiiHii ) \times 10^{-5} \) 
and
\(N(\rm {O^{++}})/N(\rm H^{+}) = ( \oiiiHii  \pm \eoiiiHii) \times 10^{-5} \), 
yielding  
\( (\rm {O^{+}}+\rm {O^{++}})/\rm {H^{+}} = 
( \oiioiii \pm \eoiioiii) \times 10^{-5} \), 
(i.e. $12 + \log \rm [(O^+ + O^{++})/H^+] = 7.26$).  
%
%evh2: devo tagliare qualcosa ...
% The single ionized oxygen abundance was calculated from the
% \( \lambda 3727 \) emissivity, while the double ionized oxygen
% abundance is a weighted mean of the abundances calculated from the \(
% \lambda 4959 \) and \( \lambda 5007 \) lines. 
%
% Although this implies a total \( \rm (O/H) \) slightly larger than
% observed, a marginally higher temperature would reconcile the two
% abundances. However, only \( T_{\rm e} \leq 20\,000\, \rm K \) is
% allowed by \texttt{nebular}, and the discrepancy is well within the
% errors, so we compute all abundances by assuming the quoted parameters.
%
This choice of physical parameters predicts
\( R_{4363} \simeq 37 \), implying that the \( \lambda 4363 \) line
flux ratio should be \( \simeq 0.05 \), definitely lower than the
minimum normalized flux that we have been able to measure. 
%evh2:
%- and it is clear that the
%- \( \lambda 4363 \) line must be just below the detection limit.
%

%evh-new:
Yet, the low surface brightness of the nebula seems to suggest 
a lower electron temperature and ionization parameter, in accord with 
the above {\sc cloudy} model calculations. 
%In fact, the above comparison with 
%photoionization models in the metal-poor range indicates 
%\( \log U \approx -3 \) and $T_{\rm e} \approx 16\,000$ K.
%
In view of these uncertainties, we adopt 
$T_e = 18\,000 \pm 2\,000$ for electron temperature, implying
$12+\log(\rm O/H) = 7.37^{+0.13}_{-0.11}$ for the O abundance in SagDIG, 
in accord with the different empirical approaches discussed above. 
%
For the nitrogen and neon ionic abundances in \uks,
%\( (\rm {N^{+}/H^{+}})  = ( \niihii \pm \eniihii)\times 10^{-5} \) 
\( (\rm {N^{+}/H^{+}})  = ( 0.062 \pm 0.039)\times 10^{-5} \)
and
%\( (\rm {Ne^{++}/H^{+}})= ( \heiiihii  \pm \eheiiihii )\times 10^{-5} \).
\( (\rm {Ne^{++}/H^{+}})= ( 0.284 \pm 0.134 )\times 10^{-5} \) are derived.
% is2: calcoliamo anche le abbondanze totali
The total abundances can be computed following Peimbert \& Costero
(\cite{peimbertCostero69}) (see also van Zee \cite{vzee98}), 
obtaining $12+\log(\rm N/H)= 6.00^{+0.28}_{-0.98}$
 and $12+\log(\rm Ne/H) = 6.86^{+0.20}_{-0.38}$.  This
also implies $\log(\rm N/O)=-1.36^{+0.30}_{-1.76}$ and $\log (\rm
Ne/O)=-0.50^{+0.22}_{-0.51}$.  
The quoted uncertainties for the N and Ne abundances reflect the
measurements errors and are larger than those related to the temperature range.
Notwithstanding the above uncertainties, this represents the first
measurement of N and Ne abundances in \uks.

%-------------------------------------------------------------
\section{Discussion}

%is3: discussione originale dopo \end{document}

%evh2:
Our results confirm that SagDIG is indeed one of the most metal poor
galaxies, having abundances just $50\%$ higher than I~Zw~18.
%
%having only slightly higher abundances than I\,Zw\,18
%(approximately XXX??? \( 18\% \) higher oxygen abundance).
%
% evh:
%Indeed, the \emph{turnoff} of a \( \sim 10\, \rm Gyr \) old population
%has been revealed by {\sc hst} out to the distance of Leo~I (\( \sim
%250\, \rm kpc \); Mighell \& Rich \cite{mighell_rich96}, Gallart et
%al. \cite{gallart_etal99}), so we expect that similar observations
%could reveal the \emph{HB} of a comparably old population out to the
%distance of SagDIG (\( \simeq 1\, \rm Mpc \)).  
%
In the Local Group, only 
Leo~A has a comparably low metallicity.
%-- rev-evh:
%: it is closer than SagDIG, but its oxygen abundance seems to be
%slightly higher (\( 37\% \) above that of I\,Zw\,18).
%
The nitrogen abundance of SagDIG is consistent with that of
I\,Zw\,18 (Izotov \& Thuan \cite{izotov98}).
%-- rev-evh: 
For Leo~A, van Zee et al. (\cite{vzee+99}) report a nitrogen abundance
\( \log {\rm N/O} = -1.5 \pm 0.1 \), similar to other low-metallicity 
galaxies. 
% The ISM of SagDIG and Leo~A appear to be the most metal poor
% inside the Local Group.
%

Answering the question about whether SagDIG and Leo\,A are young
galaxies rests on the results of resolved photometry.  
For SagDIG, ground based observations indicate the presence of an
intermediate-age stellar population in addition to the young stars
(Karachentsev et al. \cite{kara_etal99}; Momany et
al. \cite{yaz_sagdig}), but the presence of a really old component is
still unproven.
As is the case of Leo\,I, where wide-field observations have revealed
an old HB (Held et al. \cite{held_etal00}, \cite{held_etal01}), deep
{\sc hst} observations of SagDIG may eventually answer the original
question of how old this metal-poor galaxy is.

%
In Leo\,A, Tolstoy et al. (\cite{tolstoy_etal}) could not rule out the
presence of such a population in their {\sc hst} photometry, although
it must be a minor component.  Recently, Dolphin et
al. (\cite{dolp+02}) have detected a few candidate RR Lyrae variable
stars, indicating the presence of an old (globular cluster like)
stellar population accounting for at least 0.1\% of the galaxy's
luminosity.

%is3: cerchiamo di rassicurare i lettori che SagDIG e' metal-poor per
%     conto proprio, non per effetti di ambiente
{\changed Finally, it is important to check that 
the low metal content of SagDIG is the result of normal chemical
evolution, and not, for example, of gas stripping by a massive
neighbor. Indeed, its position at the margins of the Local Group seems
to ensure self-regulated evolution, an idea that we reinforce by
the following simple considerations of chemical evolution.  }


%is3: commento qui sotto
%-------- IS THE CLOSED-BOX MODEL A GOOD APPROXIMATION ? ------ 
%Even before deep {\sc hst} observations, some information on the
%evolutionary status of SagDIG can be derived by using the presently
%available data. 

In the closed-box model, the expected ISM oxygen
abundance is given by 
%
\( (\rm O/H) = (O/H)_{y} \ln ({1}/\mu) \)
%
(Searle \& Sargent \cite{ssargent72}).
%
A {}``standard{}'' value for the yield of oxygen, \( \rm (O/H)_{y} \),
is \( 2/3 \) of the solar abundance, i.e. \( \rm (O/H)_{y}=4.9\times
10^{-4} \) (e.g. Pagel \cite{pagel97}); and \(\mu =G/(G+S) \) is the
ratio of the gas mass \( G \) over total mass (\( S \) is the mass in
stars).

The gas mass given by Young \& Lo (\cite{youngLo97}) is \( 1.3\times
10^{7}\, \rm \textrm{M}_{\odot } \), and the mass in stars and
remnants, computed by integrating the SFR given by Karachentsev et
al. (\cite{kara_etal99}) in three age intervals, is \(
2\times 10^{6}\, M_{\odot } \).
%
We thus obtain \( \mu =0.86 \), a value that in the closed-box model
implies \( 12+\log (\rm O/H)=7.87 \), higher than the measured one.
We would need \( \mu =0.97 \) to reproduce our observation.  The
observed lower O abundance could be explained if part of the enriched
gas is lost into the intergalactic medium.
%- evh2:
%If we call \( G \) the gas mass, and \( S \) the star mass (thus \(
%\mu =G/(G+S) \)),
With some algebra, we find that the mass lost should be
%evh2: deltaG
\( \Delta G=0.11\times G=1.5\times 10^{6}\, \rm \textrm{M}_{\odot }
\).  \
According to Karachentsev et al. (\cite{kara_etal99}), the SFR
during the last \( 200\, \rm Myr \) was \( \sim 10 \) times higher than
in the past, with a further enhancement during the last \( 50\, \rm Myr
\). If we assume that the mass loss occurred during the last \( 200 \)
million years, then we obtain a mass outflow of \( \dot{M}=0.01\, \rm
\textrm{M}_{\odot }\, yr^{-1} \), while \( \dot{M}=0.03\, \rm
\textrm{M}_{\odot }\, yr^{-1}
\) if the mass outflow occurred during the last \( 50 \) million
years. These are not unreasonable values, since for example outflows
of the order of \( 0.2\div 2\, \rm M_{\odot }\, yr^{-1} \) are
observed in the case of NGC~1705 (Meurer et al. \cite{meurer_etal98}),
%is3: aggiungo qui
{\changed another isolated dwarf galaxy.}


%evh: CHE CI SIA GAS ATTORNO MI CONVINCE POCO ...
%A source of concern might be the fact that \( 1.5\times 10^{6}\, \rm
%M_{\odot } \) in gas would have been easily detected 
%by Young \& Lo \cite{youngLo97}, 
%%evh:
%unless the ejected gas was 
%transported out of the field of view of the HI mapping, which
%corresponds to \( \sim 1.5\, \rm kpc \) at the distance of SagDIG. Now
%a modest outflow velocity of \( 50\, \rm km/s \) would carry the gas
%%evh:
%out to \( 10\, \rm kpc \) in \( 200\, \rm Myr \), so it is
%possible that the gas is indeed present 
%%evh
%in the vicinity of the galaxy.

Using the higher value of the gas fraction, 
we find an {average} metallicity for
the \emph{stellar} component which is in agreement with the values
derived from broadband photometry.  Using Eq.~(8.7) of Pagel (1997),
an average oxygen abundance of \( 0.01 \) solar is obtained for
the stars. For solar-scaled metal abundances, this would mean
a metallicity \( \rm [Fe/H]=-2 \), 
%evh2: le evidenze per ora sono contro l'alfa-enhancem., quindi ne
%approfitto per guadagnare spazio ...
%
% , while if we take into account a possible enhancement of the oxygen
% with respect to iron, then the implied average stellar abundance would
% be \( \rm [Fe/H]\simeq -2.3 \), if we apply the typical enhancement
% observed for globular cluster stars (e.g. Carney \cite{carney96}).
%
%evh
which agrees with the stellar metallicity measured by 
Momany et al. (\cite{yaz_sagdig}), [Fe/H]$\simeq -2.1$. 
%evh2: concetto ripetuto...
% There is therefore added evidence that part of the processed gas must
% have been lost, possibly in the recent past of the galaxy.

%evh:
Our conclusion that a very metal-poor ISM is present in SagDIG goes in
the direction of the accumulating evidence that very low oxygen
abundances can be maintained throughout the life of a
%evh:
dwarf galaxy, even if it is not experiencing its first star formation
episode,
%is3:
{\changed and without invoking special environmental conditions.}
%
%
It is fairly common that galaxies once believed to contain
only a young population actually reveal an old halo when modern data
are obtained (see Kunth \& {\"O}stlin
\cite{kunthOstlin00} for a review; also Saviane et
al. \cite{saviane_etal01}).
%evh:
%, so it is our expectation that the known
%metal-poor galaxies that are farther than \( 3\, \rm Mpc \) will show
%a cosmologically old populations when e.g. {\sc ngst} will be on line
%in the near future.

%%%%%%%%%




%-----------------------------------------------------------------
\acknowledgements 
We thank Michael Sterzik for his precious help with the observations,
and Stefano Ciroi for sharing with us his expertise in {\sc cloudy}.
We also thank the referee, Dr. C. Leitherer, for helpful comments. 

%-----------------------------------------------------------------
\begin{thebibliography}{}

\bibitem[2001]{alle+01}
Allende Prieto, C., Lambert, D.~L., \& Asplund, M.\ 2001, \apjl, 556, L63 

\bibitem[1999]{aloisi_etal99} 
Aloisi, A., Tosi, M., \& Greggio, L.\ 1999, \aj, 118, 302

\bibitem[1989]{ande+grev89} 
Anders, E.,~\& Grevesse, N.\ 1989, Geochim. Cosmochim. Acta, 53, 197

\bibitem[1989]{cardelli} 
Cardelli, J.~A., Clayton, G.~C., \& Mathis, J.~S.\ 1989, \apj, 345, 245

%evh2: \bibitem[1996]{carney96} Carney, B.~W.\ 1996, \pasp, 108, 900

\bibitem[1977]{cesarsky_etal77} 
Cesarsky, D.~A., Lequeux, J., Laustsen, S., 
Schuster, H.-E., \& West, R.~M.\ 1977, \aap, 61, L31

\bibitem[1987]{cook87} Cook K.H., 1987, Ph.D. Thesis, University of Arizona

%\bibitem[1991]{deVaucouleurs_etal91}de Vaucouleurs, G., da Vaucouleurs, A., Corwin, H.G., Buta, R.J, Paturel,
%G., Fouque, P., 1991, \emph{Third Reference Catalogue of Bright
%Galaxies}

\bibitem[2000]{dpm00} 
D{\'{\i}}az, A.~I., \& P{\' e}rez-Montero, E.\ 2000, \mnras, 312, 130

\bibitem[2002]{dolp+02} 
Dolphin, A. E., Saha, A., Claver, et al. 2002, AJ in press [astro-ph/0202381]

\bibitem[1984]{edm_pagel84} 
Edmunds, M.~G., \& Pagel, B.~E.~J.\ 1984, \mnras, 211, 507

%\bibitem[1999]{gallart_etal99}Gallart, C., Freedman, W.L., Aparicio, A., Bertelli, G., Chiosi, C.,
%1999, AJ, 118, 2245 

\bibitem[1998]{ferlandEtal98}  
Ferland, G. J., Korista, K. T., Verner, D. A., Ferguson, J. W., 
Kingdon, J. B., \& Verner, E. M. 1998,  PASP, 110, 761

% \bibitem[1989]{grev+ande89} Grevesse, N.~\& 
% Anders, E.\ 1989, AIP Conf.~Proc.~183: Cosmic Abundances of Matter, 1.

\bibitem[1992]{hamuy92} 
Hamuy, M., Walker, A.~R., Suntzeff, N.~B., Gigoux, P., 
Heathcote, S.~R., \& Phillips, M.~M.\ 1992, \pasp, 104, 533

\bibitem[1994]{hamuy94} 
Hamuy, M., Suntzeff, N.~B., Heathcote, S.~R., Walker, A.~R., 
Gigoux, P., \& Phillips, M.~M.\ 1994, \pasp, 106, 566

\bibitem[2000]{held_etal00} 
Held, E.~V., Saviane, I., Momany, Y., \& Carraro, G.\ 2000, \apjl, 530, L85

\bibitem[2001]{held_etal01} 
Held, E.~V., Clementini, G., Rizzi, L., Momany, Y., 
Saviane, I., \& Di Fabrizio, L. 2001, ApJ, 562, L39

\bibitem[1998]{h-go98} 
Hidalgo-Gamez, A.~M., \& Olofsson, K.\ 1998, \aap, 334, 45

\bibitem[2001]{holweger01}
Holweger,�~H. 2001, in Proceedings of the
Joint SOHO/ACE workshop {\it Solar and Galactic Composition}, ed. 
R. F. Wimmer-Schweingruber, AIP Conf. Proc. 598,  
p.23


\bibitem[1987]{hummer} 
Hummer, D.~G., \& Storey, P.~J.\ 1987, \mnras, 224, 801

\bibitem[1999]{hh99} 
Hunter, D.~A., \& Hoffman, L.\ 1999, \aj, 117, 2789

%\bibitem[1986]{hunter_etal86}Hunter, D.A., Gillett, F.C., Gallagher, J.S., Rice, W.L., Low, F.J.
%1986, ApJ, 303, 171

% \bibitem[2000]{izotov_etal00} 
% Izotov, Y.~I., Papaderos, P., Thuan, T.~X., Fricke, K., Foltz, C.,
% Guseva, N. 2000. In Proc. of the MPA/ESO Workshop, The First Stars, 
% eds. A. Weiss, G. Abel, \& V. Hill (Berlin:Springer), 303


\bibitem[1998]{izotov98} Izotov, Y.~I.~\& Thuan, T.~X.\ 1998, \apj, 497, 227

\bibitem[2001]{izotov_etal00}
Izotov, Y.~I., Chaffee, F.~H., Foltz, C.~B., et al. 2001, \apj, 560, 222

\bibitem[1999]{kara_etal99} 
Karachentsev, I., Aparicio, A., \& Makarova, L.\ 1999, \aap, 352, 363

\bibitem[2000]{kennicuttEtal00} 
Kennicutt, R. C., Bresolin, F., French, H., \& Martin, P. 2000, ApJ, 537, 589

\bibitem[2000]{kniazev00} 
Kniazev, A.~Y., Pustilnik, S. A., Masegosa, J., et al.\ 2000, \aap, 357, 101

\bibitem[2000]{kunthOstlin00} 
Kunth, D., \& {\"O}stlin, G.\ 2000, \aapr, 10, 1

\bibitem[1983]{kunthSargent83} Kunth, D., \& Sargent, W. L. W. 1983, ApJ,
273, 81

\bibitem[2000]{lee_kim00} 
Lee, M.~G., \& Kim, S.~C.\ 2000, \aj, 119, 777

\bibitem[1991]{mcga91} Mc~Gaugh, S. 1991, \apj, 380, 140

\bibitem[1998]{meurer_etal98} 
Meurer, G.~R., Staveley-Smith, L., \& Killeen, N.~E.~B.\ 1998, \mnras, 300, 705

%\bibitem[1996]{mighell_rich96}Mighell, K.J., Rich, R.M., 1996, AJ, 111, 777
%evh:

\bibitem[1972]{mihalas72} 
Mihalas, D. 1972, Non-LTE Model Atmospheres for B \& O Stars, NCAR-TN/STR-76

\bibitem[2001]{yaz_sagdig}
 Momany, Y., Held, E.~V., Saviane, I., \& Rizzi, L. 2001, \aap, 384, 393

\bibitem[2000]{oey00} 
 Oey, M.~S., \& Shields, J.~C.\ 2000, \apj, 539, 687

\bibitem[1990]{oke90} Oke, J.~B.\ 1990, \aj, 99, 1621

%\bibitem[1989]{agn2}Osterbrock D.E., 1989. \emph{Astrophysics of Gaseous Nebulae and Active
%Galactic Nuclei}. University Science Books, Mill Valley, California
%(AGN\( ^{2} \))

\bibitem[1989]{agn2} 
Osterbrock, D. E. 1989, Astrophysics of Gaseous Nebulae and Active
Galactic Nuclei (Mill Valley: Univ. Sci.) 

 
\bibitem[1997]{pagel97}
Pagel, B.~E.~J.\ 1997, {Nucleosynthesis and Chemical Evolution
of Galaxies} (Cambridge: Cambridge University Press)

\bibitem[1979]{pagel_etal79} 
Pagel, B.~E.~J., Edmunds, M.~G., Blackwell, D.~E., Chun, M.~S., 
\& Smith, G.\ 1979, \mnras, 189, 95

\bibitem[1969]{peimbertCostero69} Peimbert, M., \& Costero, R. 1969,
Bol. Obs. Tonantzintla y Tacubaya, 5, 3

\bibitem[2000]{pilyugin00} Pilyugin, L. S. 2000, A\&A, 362, 325
\bibitem[2001a]{pilyugin01a} Pilyugin, L. S. 2001a, A\&A, 369, 594
\bibitem[2001b]{pilyugin01b} Pilyugin, L. S. 2001b, A\&A, 374, 412

\bibitem[2001]{saviane_etal01}
Saviane, I., Held, E.~V., Momany Y., \& Rizzi, L. \ 2001, 
in Proceedings of the 44th National Congress of the Italian
Astron. Society, eds. L. A. Antonelli, G. Bono, G. Giobbi, 
\& N. Menci, Mem.S.A.It. 72, n.4, 773

\bibitem[1972]{ssargent72} 
Searle, L., \& Sargent, W.~L.~W.\ 1972, \apj, 173, 25

\bibitem[1975]{seaton75} 
Seaton, M.~J.\ 1975, \mnras, 170, 475

\bibitem[1995]{sd95} 
Shaw, R.~A., \& Dufour, R.~J.\ 1995, \pasp, 107, 896

\bibitem[1989]{skillman89} 
Skillman, E. D. 1989, ApJ, 347, 883 (S89) 

\bibitem[1989a]{skh89} 
Skillman, E.~D., Kennicutt, R.~C., \& Hodge, P.~W.\ 1989, \apj, 347, 875 

%evh2:
\bibitem[1989b]{stm89} 
Skillman, E.~D., Terlevich, R., \& Melnick, J.\ 1989, \mnras, 240, 563 (STM89)

\bibitem[1988]{skillmanEtal88} 
Skillman, E. D., Melnick, J., Terlevich, R., Moles, M. 1988, MNRAS, 230, 563

\bibitem[1991]{strobelEtal91} 
Strobel, N.~V., Hodge, P., \& Kennicutt, R.~C.\ 1991, \apj, 383, 148 (SHK91)

\bibitem[1998]{tolstoy_etal} 
Tolstoy, E., Gallagher, J. S., Cole, A. A., et al.\ 1998, \aj, 116, 1244

\bibitem[2000]{vanZee00} 
 van Zee, L.\ 2000, \apjl, 543, L31

%--rev-evh:
\bibitem[1998]{vzee98} 
  van Zee, L., Salzer, J.~J., Haynes, M.~P., O'Donoghue, A.~A., 
  \& Balonek, T.~J.\ 1998, \aj, 116, 2805

\bibitem[1999]{vzee+99}
 van Zee, L., Skillman, E.~D., Haynes, M.~P. 1999, AAS, 194, \#05.04

\bibitem[1997]{youngLo97} 
Young, L.~M., \& Lo, K.~Y.\ 1997, \apj, 490, 710

\end{thebibliography}
%---------------------------------------------

\end{document}

%+++++++++++++++++++++++++++++++++++++++++++++++++++
%REFEREE FMT: TABELLE

%+++++++++++++++++++++++++++++++++++++++++++++++++++
%REFEREE FMT: FIGURE

%+++++++++++++++++++++++++++++++++++++++++++++++++++
%\end{document}


% is3: questa e' la discussione originale


%evh2:
Our results confirm that SagDIG is indeed one of the most metal poor
galaxies, having abundances just $50\%$ higher than I~Zw~18.
%
%having only slightly higher abundances than I\,Zw\,18
%(approximately XXX??? \( 18\% \) higher oxygen abundance).
%
% evh:
%Indeed, the \emph{turnoff} of a \( \sim 10\, \rm Gyr \) old population
%has been revealed by {\sc hst} out to the distance of Leo~I (\( \sim
%250\, \rm kpc \); Mighell \& Rich \cite{mighell_rich96}, Gallart et
%al. \cite{gallart_etal99}), so we expect that similar observations
%could reveal the \emph{HB} of a comparably old population out to the
%distance of SagDIG (\( \simeq 1\, \rm Mpc \)).  
%
In the Local Group, only 
Leo~A has a comparably low metallicity.
%-- rev-evh:
%: it is closer than SagDIG, but its oxygen abundance seems to be
%slightly higher (\( 37\% \) above that of I\,Zw\,18).
%
The nitrogen abundance of SagDIG is consistent with that of
I\,Zw\,18 (Izotov \& Thuan \cite{izotov98}).
%-- rev-evh: 
For Leo~A, van Zee et al. (\cite{vzee+99}) report a nitrogen abundance
\( \log {\rm N/O} = -1.5 \pm 0.1 \), similar to other low-metallicity 
galaxies. 
% The ISM of SagDIG and Leo~A appear to be the most metal poor
% inside the Local Group.

Answering the question about whether SagDIG and Leo\,A are young
galaxies rests on the results of resolved photometry.  
For SagDIG, ground based observations indicate the presence of an
intermediate-age stellar population in addition to the young stars
(Karachentsev et al. \cite{kara_etal99}; Momany et
al. \cite{yaz_sagdig}), but the presence of a really old component is
still unproven.
As is the case of Leo\,I, where wide-field observations have revealed
an old HB (Held et al. \cite{held_etal00}, \cite{held_etal01}), deep
{\sc hst} observations of SagDIG may eventually answer the original
question of how old this metal-poor galaxy is.

%
In Leo\,A, Tolstoy et al. (\cite{tolstoy_etal}) could not rule out the
presence of such a population in their {\sc hst} photometry, although
it must be a minor component.  Recently, Dolphin et
al. (\cite{dolp+02}) have detected a few candidate RR Lyrae variable
stars, indicating the presence of an old (globular cluster like)
stellar population accounting for at least 0.1\% of the galaxy's
luminosity.


%-------- IS THE CLOSED-BOX MODEL A GOOD APPROXIMATION ? ------ 
Even before deep {\sc hst} observations, some information on the
evolutionary status of SagDIG can be derived by using the presently
available data. In the closed-box model, the expected ISM oxygen
abundance is given by 
%
\( (\rm O/H) = (O/H)_{y} \ln ({1}/\mu) \)
%
(Searle \& Sargent \cite{ssargent72}).
%
A {}``standard{}'' value for the yield of oxygen, \( \rm (O/H)_{y} \),
is \( 2/3 \) of the solar abundance, i.e. \( \rm (O/H)_{y}=4.9\times
10^{-4} \) (e.g. Pagel \cite{pagel97}); and \(\mu =G/(G+S) \) is the
ratio of the gas mass \( G \) over total mass (\( S \) is the mass in
stars).

The gas mass given by Young \& Lo (\cite{youngLo97}) is \( 1.3\times
10^{7}\, \rm \textrm{M}_{\odot } \), and the mass in stars and
remnants, computed by integrating the SFR given by Karachentsev et
al. (\cite{kara_etal99}) in three age intervals, is \(
2\times 10^{6}\, M_{\odot } \).
%
We thus obtain \( \mu =0.86 \), a value that in the closed-box model
implies \( 12+\log (\rm O/H)=7.87 \), higher than the measured one.
We would need \( \mu =0.97 \) to reproduce our observation.  The
observed lower O abundance could be explained if part of the enriched
gas is lost into the intergalactic medium.
%- evh2:
%If we call \( G \) the gas mass, and \( S \) the star mass (thus \(
%\mu =G/(G+S) \)),
With some algebra, we find that the mass lost should be
%evh2: deltaG
\( \Delta G=0.11\times G=1.5\times 10^{6}\, \rm \textrm{M}_{\odot }
\).  \
According to Karachentsev et al. (\cite{kara_etal99}), the SFR
during the last \( 200\, \rm Myr \) was \( \sim 10 \) times higher than
in the past, with a further enhancement during the last \( 50\, \rm Myr
\). If we assume that the mass loss occurred during the last \( 200 \)
million years, then we obtain a mass outflow of \( \dot{M}=0.01\, \rm
\textrm{M}_{\odot }\, yr^{-1} \), while \( \dot{M}=0.03\, \rm
\textrm{M}_{\odot }\, yr^{-1}
\) if the mass outflow occurred during the last \( 50 \) million
years. These are not unreasonable values, since for example outflows
of the order of \( 0.2\div 2\, \rm M_{\odot }\, yr^{-1} \) are
observed in the case of NGC~1705 (Meurer et al. \cite{meurer_etal98}).

%evh: CHE CI SIA GAS ATTORNO MI CONVINCE POCO ...
%A source of concern might be the fact that \( 1.5\times 10^{6}\, \rm
%M_{\odot } \) in gas would have been easily detected 
%by Young \& Lo \cite{youngLo97}, 
%%evh:
%unless the ejected gas was 
%transported out of the field of view of the HI mapping, which
%corresponds to \( \sim 1.5\, \rm kpc \) at the distance of SagDIG. Now
%a modest outflow velocity of \( 50\, \rm km/s \) would carry the gas
%%evh:
%out to \( 10\, \rm kpc \) in \( 200\, \rm Myr \), so it is
%possible that the gas is indeed present 
%%evh
%in the vicinity of the galaxy.

Using the higher value of the gas fraction, 
we find an {average} metallicity for
the \emph{stellar} component which is in agreement with the values
derived from broadband photometry.  Using Eq.~(8.7) of Pagel (1997),
an average oxygen abundance of \( 0.01 \) solar is obtained for
the stars. For solar-scaled metal abundances, this would mean
a metallicity \( \rm [Fe/H]=-2 \), 
%evh2: le evidenze per ora sono contro l'alfa-enhancem., quindi ne
%approfitto per guadagnare spazio ...
%
% , while if we take into account a possible enhancement of the oxygen
% with respect to iron, then the implied average stellar abundance would
% be \( \rm [Fe/H]\simeq -2.3 \), if we apply the typical enhancement
% observed for globular cluster stars (e.g. Carney \cite{carney96}).
%
%evh
which agrees with the stellar metallicity measured by 
Momany et al. (\cite{yaz_sagdig}), [Fe/H]$\simeq -2.1$. 
%evh2: concetto ripetuto...
% There is therefore added evidence that part of the processed gas must
% have been lost, possibly in the recent past of the galaxy.

%evh:
Our conclusion that a very metal-poor ISM is present in SagDIG goes in
the direction of the accumulating evidence that very low oxygen
abundances can be maintained throughout the life of a
%evh:
dwarf galaxy, even if it is not experiencing its first star formation
episode. It is fairly common that galaxies once believed to contain
only a young population actually reveal an old halo when modern data
are obtained (see Kunth \& {\"O}stlin
\cite{kunthOstlin00} for a review; also Saviane et
al. \cite{saviane_etal01}).
%evh:
%, so it is our expectation that the known
%metal-poor galaxies that are farther than \( 3\, \rm Mpc \) will show
%a cosmologically old populations when e.g. {\sc ngst} will be on line
%in the near future.

%%%%%%%%%

\changed{

------

A proposito del punto finale, vedi la mia risposta. Comunque puo' forse
valere la pena di stressare nel testo che SagDIG e' isolata, quindi la
sua evoluzione puo' essere considerata "normale", nel senso che non e'
necessario stripping per giustificare la abbondanza di ossigeno, e
quindi le nostre conclusioni si possono in effetti applicare agli
oggetti piu' lontani/isolati.  [ actually, quanto isolati siano non lo
so, ci vorrebbe tutta una esplorazione della letteratura, ma lasciamolo
ad un altro articolo! ]

}


astro-ph0205355.extracted_bib
\begin{thebibliography}{}

\bibitem[2001]{alle+01}
Allende Prieto, C., Lambert, D.~L., \& Asplund, M.\ 2001, \apjl, 556, L63 

\bibitem[1999]{aloisi_etal99} 
Aloisi, A., Tosi, M., \& Greggio, L.\ 1999, \aj, 118, 302

\bibitem[1989]{ande+grev89} 
Anders, E.,~\& Grevesse, N.\ 1989, Geochim. Cosmochim. Acta, 53, 197

\bibitem[1989]{cardelli} 
Cardelli, J.~A., Clayton, G.~C., \& Mathis, J.~S.\ 1989, \apj, 345, 245

%evh2: \bibitem[1996]{carney96} Carney, B.~W.\ 1996, \pasp, 108, 900

\bibitem[1977]{cesarsky_etal77} 
Cesarsky, D.~A., Lequeux, J., Laustsen, S., 
Schuster, H.-E., \& West, R.~M.\ 1977, \aap, 61, L31

\bibitem[1987]{cook87} Cook K.H., 1987, Ph.D. Thesis, University of Arizona

%\bibitem[1991]{deVaucouleurs_etal91}de Vaucouleurs, G., da Vaucouleurs, A., Corwin, H.G., Buta, R.J, Paturel,
%G., Fouque, P., 1991, \emph{Third Reference Catalogue of Bright
%Galaxies}

\bibitem[2000]{dpm00} 
D{\'{\i}}az, A.~I., \& P{\' e}rez-Montero, E.\ 2000, \mnras, 312, 130

\bibitem[2002]{dolp+02} 
Dolphin, A. E., Saha, A., Claver, et al. 2002, AJ in press [astro-ph/0202381]

\bibitem[1984]{edm_pagel84} 
Edmunds, M.~G., \& Pagel, B.~E.~J.\ 1984, \mnras, 211, 507

%\bibitem[1999]{gallart_etal99}Gallart, C., Freedman, W.L., Aparicio, A., Bertelli, G., Chiosi, C.,
%1999, AJ, 118, 2245 

\bibitem[1998]{ferlandEtal98}  
Ferland, G. J., Korista, K. T., Verner, D. A., Ferguson, J. W., 
Kingdon, J. B., \& Verner, E. M. 1998,  PASP, 110, 761

% \bibitem[1989]{grev+ande89} Grevesse, N.~\& 
% Anders, E.\ 1989, AIP Conf.~Proc.~183: Cosmic Abundances of Matter, 1.

\bibitem[1992]{hamuy92} 
Hamuy, M., Walker, A.~R., Suntzeff, N.~B., Gigoux, P., 
Heathcote, S.~R., \& Phillips, M.~M.\ 1992, \pasp, 104, 533

\bibitem[1994]{hamuy94} 
Hamuy, M., Suntzeff, N.~B., Heathcote, S.~R., Walker, A.~R., 
Gigoux, P., \& Phillips, M.~M.\ 1994, \pasp, 106, 566

\bibitem[2000]{held_etal00} 
Held, E.~V., Saviane, I., Momany, Y., \& Carraro, G.\ 2000, \apjl, 530, L85

\bibitem[2001]{held_etal01} 
Held, E.~V., Clementini, G., Rizzi, L., Momany, Y., 
Saviane, I., \& Di Fabrizio, L. 2001, ApJ, 562, L39

\bibitem[1998]{h-go98} 
Hidalgo-Gamez, A.~M., \& Olofsson, K.\ 1998, \aap, 334, 45

\bibitem[2001]{holweger01}
Holweger,�~H. 2001, in Proceedings of the
Joint SOHO/ACE workshop {\it Solar and Galactic Composition}, ed. 
R. F. Wimmer-Schweingruber, AIP Conf. Proc. 598,  
p.23


\bibitem[1987]{hummer} 
Hummer, D.~G., \& Storey, P.~J.\ 1987, \mnras, 224, 801

\bibitem[1999]{hh99} 
Hunter, D.~A., \& Hoffman, L.\ 1999, \aj, 117, 2789

%\bibitem[1986]{hunter_etal86}Hunter, D.A., Gillett, F.C., Gallagher, J.S., Rice, W.L., Low, F.J.
%1986, ApJ, 303, 171

% \bibitem[2000]{izotov_etal00} 
% Izotov, Y.~I., Papaderos, P., Thuan, T.~X., Fricke, K., Foltz, C.,
% Guseva, N. 2000. In Proc. of the MPA/ESO Workshop, The First Stars, 
% eds. A. Weiss, G. Abel, \& V. Hill (Berlin:Springer), 303


\bibitem[1998]{izotov98} Izotov, Y.~I.~\& Thuan, T.~X.\ 1998, \apj, 497, 227

\bibitem[2001]{izotov_etal00}
Izotov, Y.~I., Chaffee, F.~H., Foltz, C.~B., et al. 2001, \apj, 560, 222

\bibitem[1999]{kara_etal99} 
Karachentsev, I., Aparicio, A., \& Makarova, L.\ 1999, \aap, 352, 363

\bibitem[2000]{kennicuttEtal00} 
Kennicutt, R. C., Bresolin, F., French, H., \& Martin, P. 2000, ApJ, 537, 589

\bibitem[2000]{kniazev00} 
Kniazev, A.~Y., Pustilnik, S. A., Masegosa, J., et al.\ 2000, \aap, 357, 101

\bibitem[2000]{kunthOstlin00} 
Kunth, D., \& {\"O}stlin, G.\ 2000, \aapr, 10, 1

\bibitem[1983]{kunthSargent83} Kunth, D., \& Sargent, W. L. W. 1983, ApJ,
273, 81

\bibitem[2000]{lee_kim00} 
Lee, M.~G., \& Kim, S.~C.\ 2000, \aj, 119, 777

\bibitem[1991]{mcga91} Mc~Gaugh, S. 1991, \apj, 380, 140

\bibitem[1998]{meurer_etal98} 
Meurer, G.~R., Staveley-Smith, L., \& Killeen, N.~E.~B.\ 1998, \mnras, 300, 705

%\bibitem[1996]{mighell_rich96}Mighell, K.J., Rich, R.M., 1996, AJ, 111, 777
%evh:

\bibitem[1972]{mihalas72} 
Mihalas, D. 1972, Non-LTE Model Atmospheres for B \& O Stars, NCAR-TN/STR-76

\bibitem[2001]{yaz_sagdig}
 Momany, Y., Held, E.~V., Saviane, I., \& Rizzi, L. 2001, \aap, 384, 393

\bibitem[2000]{oey00} 
 Oey, M.~S., \& Shields, J.~C.\ 2000, \apj, 539, 687

\bibitem[1990]{oke90} Oke, J.~B.\ 1990, \aj, 99, 1621

%\bibitem[1989]{agn2}Osterbrock D.E., 1989. \emph{Astrophysics of Gaseous Nebulae and Active
%Galactic Nuclei}. University Science Books, Mill Valley, California
%(AGN\( ^{2} \))

\bibitem[1989]{agn2} 
Osterbrock, D. E. 1989, Astrophysics of Gaseous Nebulae and Active
Galactic Nuclei (Mill Valley: Univ. Sci.) 

 
\bibitem[1997]{pagel97}
Pagel, B.~E.~J.\ 1997, {Nucleosynthesis and Chemical Evolution
of Galaxies} (Cambridge: Cambridge University Press)

\bibitem[1979]{pagel_etal79} 
Pagel, B.~E.~J., Edmunds, M.~G., Blackwell, D.~E., Chun, M.~S., 
\& Smith, G.\ 1979, \mnras, 189, 95

\bibitem[1969]{peimbertCostero69} Peimbert, M., \& Costero, R. 1969,
Bol. Obs. Tonantzintla y Tacubaya, 5, 3

\bibitem[2000]{pilyugin00} Pilyugin, L. S. 2000, A\&A, 362, 325
\bibitem[2001a]{pilyugin01a} Pilyugin, L. S. 2001a, A\&A, 369, 594
\bibitem[2001b]{pilyugin01b} Pilyugin, L. S. 2001b, A\&A, 374, 412

\bibitem[2001]{saviane_etal01}
Saviane, I., Held, E.~V., Momany Y., \& Rizzi, L. \ 2001, 
in Proceedings of the 44th National Congress of the Italian
Astron. Society, eds. L. A. Antonelli, G. Bono, G. Giobbi, 
\& N. Menci, Mem.S.A.It. 72, n.4, 773

\bibitem[1972]{ssargent72} 
Searle, L., \& Sargent, W.~L.~W.\ 1972, \apj, 173, 25

\bibitem[1975]{seaton75} 
Seaton, M.~J.\ 1975, \mnras, 170, 475

\bibitem[1995]{sd95} 
Shaw, R.~A., \& Dufour, R.~J.\ 1995, \pasp, 107, 896

\bibitem[1989]{skillman89} 
Skillman, E. D. 1989, ApJ, 347, 883 (S89) 

\bibitem[1989a]{skh89} 
Skillman, E.~D., Kennicutt, R.~C., \& Hodge, P.~W.\ 1989, \apj, 347, 875 

%evh2:
\bibitem[1989b]{stm89} 
Skillman, E.~D., Terlevich, R., \& Melnick, J.\ 1989, \mnras, 240, 563 (STM89)

\bibitem[1988]{skillmanEtal88} 
Skillman, E. D., Melnick, J., Terlevich, R., Moles, M. 1988, MNRAS, 230, 563

\bibitem[1991]{strobelEtal91} 
Strobel, N.~V., Hodge, P., \& Kennicutt, R.~C.\ 1991, \apj, 383, 148 (SHK91)

\bibitem[1998]{tolstoy_etal} 
Tolstoy, E., Gallagher, J. S., Cole, A. A., et al.\ 1998, \aj, 116, 1244

\bibitem[2000]{vanZee00} 
 van Zee, L.\ 2000, \apjl, 543, L31

%--rev-evh:
\bibitem[1998]{vzee98} 
  van Zee, L., Salzer, J.~J., Haynes, M.~P., O'Donoghue, A.~A., 
  \& Balonek, T.~J.\ 1998, \aj, 116, 2805

\bibitem[1999]{vzee+99}
 van Zee, L., Skillman, E.~D., Haynes, M.~P. 1999, AAS, 194, \#05.04

\bibitem[1997]{youngLo97} 
Young, L.~M., \& Lo, K.~Y.\ 1997, \apj, 490, 710

\end{thebibliography}