%\documentclass{aastex}
%\input psfig
\documentclass{paastex2}
\usepackage{myemulateapj5}
\usepackage{graphicx}
\usepackage{lscape}
\usepackage{dblfnote}
%\documentclass[preprint2]{aastex}
\newcommand{\Ha}{H$\alpha$}
\newcommand{\Hb}{H$\beta$}
\newcommand{\La}{Lyman-$\alpha$}
\newcommand{\Lb}{Lyman-$\beta$}
\newcommand{\Bg}{B$\gamma$}
\newcommand{\Ls}{L$_{\odot}$}
\newcommand{\Ms}{M$_{\odot}$}
\newcommand{\mlow}{$M_{\rm{low}}$}
\newcommand{\mup}{$M_{\rm{up}}$}
\newcommand{\Zs}{Z$_{\odot}$}
\newcommand{\Myr}{M$_{\odot}$~yr$^{-1}$}
\newcommand{\kms}{km~s$^{-1}$}
\newcommand{\vinf}{$v_{\infty}$}
\newcommand{\Md}{$\dot{M}$}
\newcommand{\Teff}{$t_{\rm{eff}}$}
\newcommand{\civ}{C~IV $\lambda$1550}
\newcommand{\siiv}{Si~IV $\lambda$1400}
\newcommand{\nv}{N~V $\lambda$1240}
\newcommand{\heii}{He~II $\lambda$1640}
\newcommand{\ovi}{O~VI $\lambda$1035}
\newcommand{\cii}{C~II $\lambda$1335}
\newcommand{\oi}{O~I $\lambda$1300}
\newcommand{\siii}{Si~II $\lambda$1526}
\newcommand{\siiia}{Si~II $\lambda$1260}
\newcommand{\alii}{Al~II $\lambda$1670}

%\slugcomment{Submitted to The Astrophysical Journal}

%\shorttitle{Time Dependent Dust Obscuration}
%\shortauthors{Leitherer et al.}

\begin{document}
\vspace*{\fill}
\title{Ultraviolet Spectra of Star Forming Galaxies}
\title{with Time Dependent Dust Obscuration}

\author{Claus Leitherer}%\altaffilmark{1,2,3}}
\affil{Space Telescope Science Institute\/\altaffilmark{1}\\
       3700 San Martin Drive\\
       Baltimore, MD 21218\\
       leitherer@stsci.edu}


\author{Daniela Calzetti}
\affil{Space Telescope Science Institute\/\altaffilmark{1}\\
       3700 San Martin Drive\\
       Baltimore, MD 21218\\
       calzetti@stsci.edu}

% \and

\author{Lucimara P.\ Martins}
\affil{Instituto Astron\^omico e Geof\'{\i}sico, USP,\\ 
        Av.\ Miguel Stefano, 4200\\
        04301-904 S\~ao Paulo, Brazil\\
        lucimara@astro.iag.usp.br}

\vspace{1in}
\centerline{\titlelarge To be published in the 20 July 2002 issue of \emph{The Astrophysical Journal}}
\received{8 February 2002}
\accepted{25 March 2002}
\sluginfo
\vfill

\altaffiltext{1}{Operated by AURA, Inc., under NASA contract NAS5-26555.}


\begin{abstract}
We study the effect of time dependent dust obscuration on synthetic 
ultraviolet line profiles of a young stellar population. If the
youngest and most massive stars are more obscured than the older,
less massive stars, the \civ\ and other stellar-wind lines are
significantly diluted with respect to a simple foreground screen
model for the dust. We propose to use stellar-wind lines as a
probe of the dust-obscuration model instead of the previously
employed nebular emission lines. Since purely stellar diagnostics
are utilized, uncertain assumptions on the nebular properties are
unnecessary. Photoionization models demonstrate that the \civ\
emission is typically dominated by stellar winds, and nebular 
contamination is negligible. A first comparison with a galaxy
sample observed with HUT favors a dust geometry affecting ionizing
and non-ionizing stars equally. We point out the need for 
higher-quality data for a more rigorous comparison. HST
is capable of obtaining such data in the future.
\end{abstract}

\keywords{galaxies: ISM---galaxies: starburst---galaxies: 
stellar content---ultraviolet: galaxies}
   

\section{Motivation}

Multi-wavelength observations of spectral energy distributions
allow us to infer the stellar content in galaxies.
The wavelength resolution of the observations may range from crude
broadband photometry to high-resolution spectroscopy, providing 
information on the overall emitted power (e.g., Schmitt et~al.\
1997) as well as on specific phenomena seen in narrow energy
ranges. Examples of the latter are stellar absorption-line indices (e.g.,
Worthey 1994) or nebular emission lines powered by stars (e.g.,
Contini \& Viegas 2001).

Interpreting observations often requires spectral synthesis models. Such
models became feasible with the availability of sufficient computing
power (e.g., Tinsley 1968) and have since reached a remarkable 
degree of sophistication (see Bruzual 2001 for a recent review). It would
be desirable to produce spectral synthesis from first principles, 
with as few ad hoc assumptions as possible. {\em Evolutionary} synthesis
aims to achieve this goal but still suffers from major shortcomings.

First and foremost, stellar evolution is far from being fully 
understood, even on the main sequence, but more so for post-main-sequence
evolution. Mixing and diffusion processes are poorly understood and
cannot be treated self-consistently (e.g., Chiosi et~al.\ 1995). 
Stellar rotation has recently been recognized to be a fundamentally
important driver of stellar evolution and therefore of the emergent
stellar radiation (Maeder \& Meynet 2000). In addition,
stellar evolution is known to be non-conservative. Stellar 
winds during all stages of evolution lead to the evaporation of
stars, recycling almost 50\% of the stellar mass back to the
galaxy within $\sim$10$^8$~yr (Leitherer et~al.\ 1999). Unfortunately,
the mass-loss rates are difficult to determine in those phase
where they matter the most, and the rates required by stellar
evolution are either unknown or disagree with the observed
 rates (e.g., Salasnich, Bressan, \& Chiosi 1999).

Another category of uncertainties comes from model atmospheres, in
particular at high energies where direct observations are infeasible. 
Indirect indicators (such as nebular diagnostics) do indeed suggest
serious shortcomings (Kewley et~al.\ 2001).

Most of the above deficiencies can to some degree be calibrated and
re-adjusted via comparison with resolved, nearby stellar populations.
While such empirically corrected models are no longer derived from
first principles and their astrophysical basis is ad hoc, they
nevertheless provide a useful description for the spectra of more distant,
unresolved populations. Most major evolutionary synthesis codes
(e.g., Mas-Hesse \& Kunth 1991; Bruzual \& Charlot 1993; 
Rieke et~al.\ 1993; Fioc \& Rocca-Volmerange 1997;
Sternberg 1998;  Leitherer et~al.\ 1999) fall into this category.

The observed spectral properties of galaxies are not only determined
by stars and gas, but also by dust. Yet, almost all spectral 
synthesis models either treat dust in an ad hoc way or completely
neglect it. Notable exceptions are the modeling efforts of Dwek (1998),
Silva et~al.\ (1998), Poggianti et~al.\ (1999),
or Gordon et~al.\ (2001). The justification for 
not including the detailed radiative transfer for dust in star-forming
galaxies is the
observational evidence for a simple ``effective'' obscuration law
(Calzetti 2001). Although its physical basis is not fully
understood, it acts similarly to an inhomogeneous foreground screen,
thereby affecting all stars in the same way. 

It is well known that the obscuration in very dusty star-forming
galaxies can be so high that some (e.g., Beck, Turner, \&
Gorjian 2001 for He~2$-$10) or even most (e.g., Genzel et~al.\ 1998
for ultraluminous infrared [IR] galaxies) 
of the stars are completely hidden from view. Even
in ultraviolet(UV)-bright galaxies, dust may play a crucial
role in the determination of the stellar content. There is empirical
evidence for different reddening of the stellar continuum in
the UV and the gaseous emission in the visual. The latter has
a reddening almost twice as high as that of the former (Fanelli,
O'Connell, \& Thuan 1988; Calzetti, Kinney, \& Storchi-Bergmann
1994). Two alternative explanations have been suggested: Calzetti
(1997) favors a different spatial distribution of the UV-bright
stars and the ionized gas. In this picture, the newly formed
stars {\em in any mass range} have lower dust obscuration than the
gas since stellar winds have effectively swept away gas and
dust from the birth place early in the star-formation
episode. Alternatively, Charlot \& Fall
(2000) proposed that stars and gas are spatially coincident
over the size of the starburst, but have a cloud-like 
substructure. Here, the higher reddening seen in emission
lines  results from the larger dust attenuation
of the massive, ionizing stars still embedded in their birth
clouds. Subsequently, the birth clouds disperse, and {\em less
massive} stars exhibit a {\em less reddened} UV continuum. An
equivalent way of stating the two alternative interpretations
is via the involved time scales: in Calzetti's picture the
dust dispersal occurs rapidly, on a scale shorter than the
O-star lifetime, whereas the Charlot \& Fall mechanism operates
over a scale longer than the age of O stars (about 10--20~Myr).

Both interpretations are equally successful in accounting
for the different reddening seen in optical nebular emission
lines and in the stellar UV continuum. We will refer to this
difference as the stars vs.\ dust {\em dichotomy}.
In this paper we propose
an alternative observational test to distinguish between the
two interpretations for the dichotomy.
 The key is not to combine both stellar
and nebular diagnostics, which confuses the issue. Instead
of using \Ha\ or \Hb\ as a proxy for the most massive stars,
we propose to use stellar-wind features in the UV, which 
trace the same mass regime as \Ha, but are not affected by
assumptions on the gas. The paper is organized
as follows: First we review stellar emission and absorption
lines and their sensitivity to stars in different mass
regimes (Section~2). Then we discuss a suite of synthetic
UV line profiles which allow us to address the effects
of dust on the emitted UV light (Section~3). In Section~4
we evaluate the relative strength of stellar versus nebular
\civ\ emission. A first comparison to observations is
done in Section~5. Finally, we provide conclusions and directions
for further work in Section~6.



\section{Dust sensitive spectral lines}

Our goal is to identify stellar emission and absorption lines
which originate from stars of different mass than those
seen in the adjacent
continuum. This will allow us to probe the relative importance
of stars as a function of mass from the emission- and 
absorption-line strength. We restrict our discussion to young
stellar populations with ages of order $10^8$~yr or less. The
corresponding stellar spectral types that contribute to the
integrated spectrum are OB stars in the UV to optical, and 
late-type supergiants in the red/IR part of the spectrum.

Populations with ages of more than 5--10~Myr have a strong
red supergiant component which dominates the near-IR. Their
imprints are, e.g., the CO 2.3~$\mu$m band head and the Ca~II
triplet at 8498--8662~\AA. These features are, however, not useful
for our purpose because both the features and the adjacent
continuum come from the same stars and therefore are not
sensitive to population variations. An identical effect is
that these features are entirely independent of changes in
the stellar initial mass function (IMF; Leitherer et~al.\ 1999).
This behavior is a simple consequence of stellar structure and
evolution: the red and near-IR part of the spectrum is dominated
by red supergiants only, with little contribution from main-sequence stars.
Therefore the mass-luminosity relation is quite degenerate,
and the line equivalent widths of a population
are almost independent of mass. Ultimately this is due to
the structure of the Hertzsprung-Russell diagram, which shows
an enormous luminosity spread between supergiants and main-sequence
stars in the red part. Conversely, evolved and unevolved stars
have comparable luminosities in the blue part of the diagram.
Therefore the integrated spectrum of a population with OB
stars being dominant should display the effect we are looking
for. This is indeed the case, as we will demonstrate.

OB stars have strong H and He lines in the visual part of the
spectrum, but few metallic lines. The former are in principle
strongly dependent on the population mix, both via age or  
IMF changes. In practice they become only useful in the absence
of massive ionizing stars due to blending with strong nebular emission
(e.g., Gonz\'alez Delgado, Leitherer, \& Heckman 1999). The
metallic lines at optical wavelengths are photospheric in origin
and tend to be weak in a stellar population with any reasonable
parameters. However, the situation dramatically changes in the UV.

The strongest lines in OB stars originate in powerful stellar winds.
The wind densities and ionization conditions favor transitions
from ground levels in ionized metals, like 
C$^{2+}$,
C$^{3+}$, 
N$^{4+}$,
O$^{5+}$, 
Al$^{2+}$, or
Si$^{3+}$. The corresponding wavelengths are below 3000~\AA, and the
most prominent spectral features are, e.g., \nv, \siiv, or \civ. Other
strong lines with wind features are C~III $\lambda$1175 or Al~III
$\lambda\lambda$1854,62 but their wavelengths are outside the
range that is commonly used for observations of O stars. Therefore
their behavior is less well studied.
Comprehensive discussions of the associated phenomena were given
by Robert, Leitherer, \& Heckman (1993), Leitherer, Robert, \&
Heckman (1995), Gonz\'alez Delgado, Leitherer, \& Heckman (1997),
de~Mello, Leitherer, \& Heckman (2000), and 
Leitherer et~al.\ (2001). Two key properties are relevant in the
present context: (i)~Winds are stronger in more massive and more
luminous stars. Therefore, stellar-wind {\em lines} show a dependence
on mass and luminosity in stars of any given spectral type, with
the most pronounced wind effect occurring in the earliest O types and
in stars with the highest luminosities.
(ii) The UV {\em continuum} is emitted by late-O and early-B stars.
As we observe the Rayleigh-Jeans part of the spectrum in these stars,
the UV continuum is very insensitive to changes in the population
mix over an age range of tens of Myr. Taken together, (i) and (ii)
suggest that the strength of stellar-wind lines relative
to the continuum should be very 
sensitive, e.g., to age and IMF. The previously quoted references
show that this is indeed the case. Obviously, this implies that
these lines will be affected by dust obscuration effects as well.

Wolf-Rayet (WR) stars are another class of objects whose spectral
features are strongly population dependent. WR stars are highly
evolved descendants of the most massive stars with the strongest
(non-eruptive) hot-star winds known (Willis 1999). Some of the
most prominent wind features are C~III $\lambda$5696, He~II $\lambda$4686, 
C~III/N~III $\lambda$4650,  and He~II $\lambda$1640.
While these lines are extremely strong in individual WR stars, they
are not very conspicuous in populations due the short life times
and correspondingly small numbers of WR stars (e.g., Schaerer et al. 2000).
Another complication affects the  WR lines in the optical. It is
very likely that the optical continuum is not only due to the
current star-formation episode but has a significant older component
as well. Evidence comes from the generally diminished \Ha\ and \Hb\
equivalent widths in starbursts, which can be understood in terms
of dilution effects (Raimann et al. 2000; 
Stasi\'nska, Schaerer, \& Leitherer 2001).
(Note, however, that this is not expected to be significant in the UV.)
Above all, WR evolution is still poorly understood and undergoing
revision (e.g., Maeder 1999). Therefore it would not be prudent to abandon
\Ha\ and \Hb\ as a tracer of the most massive stars in favor of
WR-star features. Given this situation, we find the O-star wind lines in
the UV as the best --- and the only --- stellar lines which allow
us to probe obscuration effects with purely stellar diagnostics. 






\section{Calculation of synthetic UV spectra}

Dust obscuration effects in UV line profiles were previously 
considered by Gasda \& Leitherer (1996), who modeled \siiv\
and \civ\ profiles for a population with {\em mass} dependent
obscuration. Their exploratory study suggested that even
small reddening values of $E(B-V) \approx 0.2$ would significantly
decrease both the absorption and emission parts of the P~Cygni
profiles. Since the resulting P~Cygni profiles were weaker
than typically observed for a reddening even lower than appropriate
for representative starburst regions, Gasda \& Leitherer 
concluded that mass dependent dust obscuration is unlikely to
be important.

Our new models are intended to investigate {\em time} dependent
obscuration. Under this hypothesis, no explicit dependence of
the dust obscuration on stellar mass is imposed. However, there
will be an indirect dependence on mass in populations with ongoing
star formation: the older the population, the more it will be
weighted towards lower mass, as less and less massive stars reach
an equilibrium between birth and death. This has an important
observational consequence: while the mass dependent dust obscuration
models of Gasda \& Leitherer (1996) predict diluted P~Cygni profiles
for any star-formation history, including single stellar populations,
time dependent models produce this effect only for mixed 
populations with continuing star formation. Star clusters formed
in an instantaneous burst would always display the same UV wind lines if
the obscuration has a time- but no explicit mass-dependence.

We parameterized the dust obscuration and its dependence on time
as follows. Stars are assumed to be embedded in fully opaque dust
clouds at age $t=0$, when their light is 100\% attenuated. For 
numerical reasons, the attenuation was actually taken to be 99.9\%
in the models, which leads to the same profiles. The attenuation 
gradually decreases with time due to the dispersal of the
circumstellar dust material. We can identify three relevant epochs
which are suggested by stellar evolutionary time scales. 
(i)~$t=5$~Myr: termination of the WN phase in which stars develop
their strongest winds; (ii)~$t=15$~Myr: evolutionary time scale
of ionizing O stars and transition to B-star phase; 
(iii)~$t=50$~Myr: lifetime of the least massive stars capable of 
producing a core-collapse Type~II supernova. We adopt these
time scales to define $t_{\rm{dust}} = 5$, 15, and 50~Myr, when
the attenuation reaches 0\%. The detailed functional behavior
between $t=0$ and $t=t_{\rm{dust}}$ is not constrained by observations.
We chose the simplest functional form, i.e., a linear decrease
of the attenuation from 100\% at $t=0$ to 0\% at $t_{\rm{dust}}$,
and 0\% thereafter. This parameterization allows a 
qualitative comparison with the effective
absorption curve of Charlot \& Fall (2000), which predicts
a decline of the attenuation by a factor of three at 10~Myr. This
is close to our model with $t_{\rm{dust}} = 15$~Myr.

It is worth pointing out that our models do neither depend on the
specifics of the amount of dust obscuration, nor on its functional
dependence on wavelength. We will exclusively discuss line profiles
which are {\em normalized to a continuum level of unity}, thereby 
circumventing uncertainties related to dereddened monochromatic
{\em luminosities}. The shape of the obscuration law is irrelevant as 
well, as the width of the P~Cygni profiles is small in
comparison with the wavelength interval over which the 
obscuration law varies.

The synthetic spectra were computed with the publicly available 
version 3.0 (December 2000) of Starburst99 (Leitherer et~al.\ 1999).
Computational details can be found in the Appendix. Our standard
model is for a power-law IMF parameterized as
$N(M) \propto M^{-\alpha}$ with
$\alpha = 2.35$. This approximates a Salpeter IMF. The upper and
lower mass limits are \mup~=~100~\Ms\ and \mlow~=~1~\Ms, 
respectively. The metallicity of the models shown in this paper
is solar, but calculations for quarter solar models give the
same quantitative results. Star formation ran continuously for
up to 50~Myr, which is about the age limit inferred for UV selected
starburst galaxies. The escaping UV light was attenuated 
according to the three attenuation laws described above. The
three sets of models with $t_{\rm{dust}} = 5$, 15, and 50~Myr
are the baseline for our parameter study.

The strongest UV lines affected by time dependent dust obscuration
are \ovi, \nv, \siiv, and \civ. \ovi\ and \nv, although included
in our modeling, will not be considered further. \ovi\ is
shortward of 1200~\AA, which places this line in a quite
inconvenient spectral region
for most spectrographs. Moreover, blending with Lyman-$\beta$ and
C~II $\lambda$1037 introduces systematic effects that would bias
the relatively subtle variations in \ovi. \nv\ is located in the
red wing of the \La\ absorption profile. This makes it rather
difficult to define the continuum level, both in our models and
in the observations. Since this line behaves very similarly to
\civ\ and provides no additional insight, we do not  discuss 
\nv\ any further.

The basic trends can be seen in Figs.~\ref{10Myr}, \ref{20Myr},
and \ref{50Myr}, which show \civ\ at starburst ages of 10, 20,
and 50~Myr, respectively. We compare the models with dust attenuation
to a completely dust-free model in each figure. Formally, the
dust-free model corresponds to either $t_{\rm{dust}} = 0$ or $\infty$
because both cases lead to the same normalized spectrum. The
spectral features in this wavelength region in addition to
C~IV are S~V $\lambda$1502, \siii, Si~II $\lambda$1533, 
C~I $\lambda$1560, and numerous unresolved metallic (mostly iron)
lines, in particular longward of 1560~\AA. At each
of the three starburst ages, the continuum with its 
weak lines shows essentially
no variation with $t_{\rm{dust}}$. This is expected, as 
the variation of $t_{\rm{dust}}$ corresponds to a stellar
evolutionary time scale during which the UV continuum properties
are almost constant. 

The C~IV P~Cygni profile varies with
$t_{\rm{dust}}$, less strongly at 10~Myr, and more strongly
at 20 and 50~Myr. $t=10$~Myr is too young for a significant
population of B stars to accumulate. The spectrum is dominated
by stars which all have strong \civ, and a selective
attenuation with small or large $t_{\rm{dust}}$ has little
effect. This is quite different at 20 and 50~Myr. Now, a 
significant B-star population with weak C~IV contributes to
the spectrum, and selective dust obscuration leads to
conspicuous changes in both the absorption and the emission.
The fully dust-free model predicts C~IV emission at least
10\% above the continuum at any age between 0 and 50~Myr.
Around and after 20~Myr, the attenuated models with the 
three considered values of $t_{\rm{dust}}$ predict no 
C~IV emission above the continuum level. The absorption
strength is reduced in a similar way. Stellar-wind 
absorption should almost always be weaker than 60\% of
the continuum level with dust obscuration. These variations are highly 
significant and well above the noise level of the
library spectra.

In comparison with C~IV, \siiv\ is far less affected by 
time dependent dust obscuration. Fig.~\ref{SiIV} shows
the spectral region centered on the Si~IV line. Other 
features in this region are 
Fe~V $\lambda\lambda$1360--1380,
Si~III $\lambda$1417,
C~III $\lambda$1427,
and Fe~V $\lambda$1430. It would
be difficult to observationally verify the slight variations
predicted by the models. Si$^{3+}$ requires an ionization
energy of only 33~eV, as opposed to 49~eV for C$^{3+}$. Therefore
\siiv\ is strong in stars over a broader temperature range than \civ\ and
can be observed throughout the O to early-B sequence (Walborn, 
Nichols-Bohlin, \& Panek 1985; Walborn, Parker, \& Nichols 1995).
As a result, selective dust obscuration has 
little effect on Si~IV in the models discussed here. Note,
however, that we would see an effect for smaller
$t_{\rm{dust}}$ and younger ages. \siiv\ is strongly luminosity
dependent and develops a P~Cygni profile not until 3 to 4~Myr
after the onset of star formation when the first supergiants 
form. Attenuating the hot main-sequence
population at $t < 4$~Myr and leaving the most recently
formed  supergiant
population unobscured would actually {\em enhance} the profile.
This can be seen in Fig.~\ref{SiIV_5} where we have plotted 
\siiv\ for a 5~Myr old population and $t_{\rm{dust}} = 5$~Myr
and for a dust-free population of the same age. The profile 
with initial dust obscuration is about 10\%
stronger due to the diminished flux
from hot main-sequence stars. The trend in Fig.~\ref{SiIV_5} is
a result of the well-known luminosity effect in Si~IV (Walborn \&
Panek 1984). We do not expect this subtle effect to be
observationally significant, in particular since a careful
fine-tuning of parameters is required: the population must
be observed around 4--5~Myr, when the first O supergiants form,
and $t_{\rm{dust}}$ must be comparable to the starburst age.
Of course, the effect would be much stronger had we adopted  
an obscuration law with a stronger bias against the main-sequence
population (e.g., a step function with 100\% obscuration until
3~Myr and 0\% thereafter). Since the dust obscuration effect
is quite small for the time dependence adopted in the present
study, however, we will not discuss \siiv\ any further and
concentrate on \civ.
 
UV line profile variations are not unique evidence for or against
time dependent dust obscuration. One can always trade off 
dust obscuration for IMF variations or time dependent star
formation. A case in point is Fig.~\ref{IMF} where we compare
C~IV of two 20~Myr old populations: one with no dust obscuration,
and the other having obscuration with $t_{\rm{dust}} = 15$~Myr
and a flatter IMF with $\alpha = 1.7$. The two profiles are
virtually identical. The dust dilution has been compensated
by an increased fraction of the most massive stars.


An analogous case would be a trade off between dust obscuration
and star-formation history. As explained in the Appendix, a
decrease of the obscuration with time is identical to
a decrease of the star-formation rate with time. This may
be less intuitive than the result of an IMF variation but it
immediately follows from the decreasing ``effective'' mass
of the stellar contributors with time. This effective mass
decreases even more rapidly if the most recent
population (which is weighted towards higher mass) forms
at a lower rate. 

Care is required when comparing the theoretical profiles to
observations. It is essential to perform a profile, rather
than an equivalent width analysis. The equivalent width of
the full profile is not very meaningful since it is the
sum of the emission and the absorption which almost compensate
each other. Furthermore, the continuum level is quite
difficult to define in observations at low spectral resolution
due to line blanketing longward of 1560~\AA. When performing
a profile analysis, the observed P~Cygni absorption must
be corrected for interstellar absorption. Interstellar lines
tend to be much stronger in star-forming galaxies than in the
Milky Way (Heckman et~al.\ 1998). Since the library stars are
Galactic, the synthetic interstellar lines will be much
weaker than in actual galaxy spectra. In addition, most
observed galaxy spectra contain blueshifted Milky Way foreground
absorption lines, again adding a bias to the C~IV absorption
component. In contrast, the emission part of the P~Cygni 
profile is not affected by such biases and makes an excellent
tracer for dust obscuration effects.




\section{Stellar versus nebular \civ\ emission}

So far we have assumed that the \civ\ emission associated with a stellar
population is only produced in winds immediately surrounding the
stars.  We also expect some additional contribution to C~IV from 
H~II regions and the ISM. In this section  we utilize photoionization
models to estimate the relative strength of this nebular
contribution. The code AANGABA (Gruenwald \& Viegas 1992) is used to 
model a homogeneous, spherical nebula which is 
photoionized by a star cluster. The numerical simulation, under
the assumption of thermal equilibrium, 
predicts the physical conditions of the gas as a function of the central distance,
which then allows the calculation of emission-line strengths.

A grid of seven representative models was generated to explore how  typical
input parameters affect the emission-line intensities. The relevant
parameters are the ionizing spectra, the abundances and the electron density
of the gas. We adopted two choices for the ionizing spectra:
one using a Salpeter IMF and  the other with an IMF having
an exponent of $\alpha = 1.5$. The former corresponds to our
standard population model used to calculate the stellar-wind profiles, and
the latter is intended to simulate the effects of an excessively hard
input spectrum. Both spectra were calculated with Starburst99 and are
therefore consistent with the wind-profile models having no dust 
obscuration. The age is 20~Myr and star formation is continuous.
We adopted four different abundance values: (i)~solar abundances 
(Anders \& Grevesse 1989); (ii)~solar abundances with dust depletion as
discussed by Kewley et~al.\ (2001); (iii)~quarter solar abundances; 
(iv)~quarter solar with dust depletion. We note that our depleted models
are not fully self-consistent. We only remove the coolants from the
gas phase, thereby increasing the electron temperature. The change
in the thermal balance via photo-electrons from dust grains and the
modification of the ionizing input spectrum by dust grains are not
taken into account (see Shields \& Kennicutt 1995 for a discussion
of these effects). The nominal density of the nebula
is $n_{\rm e} = 100$~cm$^{-3}$, but higher and lower values
of 300 and 30~cm$^{-3}$, respectively, were considered as well.
 
A summary of the seven photoionization models is given in Table~1. In
addition to the already mentioned parameters, this table lists the \Hb\
flux normalized to a unit distance of $D = 1 \times 10^{24}$~cm. The
\Hb\ luminosities are obtained after multiplication by $4 \pi D^2$. As
\Hb\ measured through an identical aperture is often unavailable for 
a galaxy observed in the UV, we also give the relative strength of
C~III $\lambda$1907 (col.~6). The ionization parameter $U$ is in col.~7.
$U$ is defined as
\begin{equation}
    U = \frac{Q_{\rm H}}{4 \pi R_{\rm S}^2 n_{\rm e}c},
\end{equation}
where $Q_{\rm H}$ is the number of hydrogen ionizing photons (s$^{-1}$), $R_{\rm S}$
is the radius of the Stromgren sphere (cm), and $c$ is the light speed (cm~s$^{-1}$).

Table~2 is a compilation of the 20 strongest lines of each 
model between 1200 and 2000~\AA\ (excluding Lyman-$\alpha$). 
We use Model~1 as a baseline
for a discussion of the relevant issues. The Starburst99
spectrum at \Zs\ tends to overestimate the photon flux in the He$^+$
continuum above 54~eV (Bresolin \& Kennicutt 2002). Therefore
recombination lines like He~II $\lambda$1640 are probably too strong
in the models and should be taken with care.
The strongest line overall
in this wavelength range is \cii.
The equivalent width of C~II in Model~1 is 0.6~\AA\footnote{Equivalent
widths of all lines can easily be calculated by noting that
the Salpeter-type stellar spectrum has a continuum luminosity 
of $1.7 \times 10^{40}$~erg~s$^{-1}$~\AA$^{-1}$ at 1500~\AA\ and
that the continuum behaves like $F_\lambda \propto \lambda^{-2.5}$
in the UV.}. The C~II equivalent widths of the other models are
smaller but comparable to those of Model~1. The C~II line was
detected in MS~1512-cB58, a galaxy with the best
available UV spectrum of any star-forming galaxy
in existence to date (see Section~5.4 for
a discussion). Pettini et~al.\ (2000) found weak but significant
emission, made possible by the blueshift of the interstellar 
absorption component. In principle, \cii\ emission could also
be stellar, as very luminous B supergiants display the line 
in emission as well (Walborn et~al.\ 1995). Synthesis models predict that
the stellar emission contribution should be weak and undetectable. The
photoionization models indicate that nebular \cii\ can be 
present under standard conditions if it is not blended with
interstellar (or stellar) absorption. This appears to be
the case in MS~1512-cB58. 

We are mainly interested in \civ. Model~1 predicts a C~IV strength
of only 4\% the strength of C~III] $\lambda$1907. Since C~III] itself
is quite faint and often not even detected in star-forming galaxies,
nebular C~IV is predicted to have observationally insignificant
flux levels. This becomes even more apparent when considering
the equivalent width, which is far less than 0.1~\AA. Other models
with solar metallicity predict stronger C~IV, in particular when
depletion of refractory elements on dust grains is taken into
account. Even in these cases, however, \civ\ is still quite
weak and will remain undetectable. Model~6, which is for 
quarter solar metallicity with depleted element abundance ratios,
predicts the largest C~IV intensity. In this model, nebular
C~IV reaches the strength of C~III] and is no longer negligible.
If C~III] is detected in a galaxy spectrum, \civ\ could be
present as well.

We conclude that \civ\ is not expected to have a significant
nebular contribution at solar metallicity. Depending on the
amount and treatment of the depletion of heavy elements, nebular
C~IV may become comparable in strength to C~III] $\lambda$1907.
In that case, an analysis of C~III and other unblended nebular lines,
such as O~III $\lambda$1663, provides constraints on nebular C~IV.
Of course, spectral observations at sufficiently high resolution
could easily distinguish between broad stellar-wind 
and narrow nebular emission.



\section{Comparison to observations}

In this section  we will review
relevant data from local and distant star-forming 
regions. The goal is to highlight possible trends and
to point out future observations that would allow one
to investigate the relation between dust obscuration,
starburst evolution, and spatial morphology with the
method proposed in this paper. We will demonstrate
that the reddening dichotomy is not found universally
 but only in star forming regions observed through
apertures encompassing areas with scales of hundreds
of parsecs.

The question we pose can be stated as follows: 
Is the different reddening derived
from the stellar continuum and the Balmer decrement 
primarily due to spatial or due to age effects? In the
former case, purely stellar reddening indicators would
be age independent, as gas and dust become
spatially separated from both O and B stars
very early in the starburst evolution.
 In the latter case, these
reddening indicators would be time variable, as
dust obscuration varies on a time scale equal to the
evolutionary time scale of OB stars and therefore
affects different masses selectively.
The two cases would be indistinguishable using the
stellar UV continuum and nebular lines as diagnostics.
This is of course trivial, as the two explanations were
put forward to account for this dichotomy in the first
place. We propose to identify
the dominant mechanism via UV stellar-wind lines.


\subsection{30 Doradus and NGC~604}

Giant H~II regions may be scaled-down versions of
powerful starburst or they may lack one or more ingredient
that defines a starburst. For instance, their release
of mechanical luminosity may not be sufficient to
maintain ``superwinds'', massive outflows
ubiquitously seen in starburst galaxies which remove the
ambient gas and dust, regulate the star formation, and
drive the chemical evolution
in these objects (Heckman 2001). Nevertheless, the
proximity of regions like 30~Doradus in the Large Magellanic
Cloud and NGC~604 in M33 permits detailed studies of
individual stars and the surrounding gas and dust at
the highest spatial resolution. This may provide insight
into processes in genuine, more distant starbursts.

30~Doradus is the most luminous giant H~II region in the
Local Group (Kennicutt 1984) and has been called the Starburst Rosetta Stone
(Walborn 1991).  A plethora of data exist in the literature.
The powering source is the central $\sim$10$^5$~\Ms\ star cluster
R136 which hosts some of the most massive stars known (Massey
\& Hunter 1998).
Star formation is ongoing in the periphery of the cluster
(Brandner et~al.\ 2001), but the second-generation stars
are faint at optical and UV wavelengths. The optical/UV
morphology is dominated by an (almost) instantaneous
star-formation event over a few pc some 2--4~Myr ago
(e.g., de~Koter, Heap, \& Hubeny 1998).

Mathis, Chu, \& Peterson (1985) compared the stellar
extinction, as measured by Fitzpatrick \& Savage (1984),
to the values derived from the Balmer decrement. The
two extinction values agree. The ratio of \Hb/radio
flux suggests very little dust scattering close to
the star. Therefore most of the dust is not well mixed
with the stars and located preferentially in front (and the
back) of the cluster. Vacca et~al.\ (1995) verified this
result from the integrated UV light distribution over a
$3' \times 3'$ (50~pc~$\times$~50~pc) region. The IUE spectrum
indicates the same reddening as measured from the
Balmer decrement, and the stellar-wind lines are as
strong as expected for a standard stellar population.
These results are consistent with both of the above
dust scenarios but essentially rule out the mass dependent
obscuration discussed by Gasda \& Leitherer (1996).
Their model would predict weaker P~Cygni profiles. 

Similar arguments apply to NGC~604. At a distance of 850~kpc,
observations comparable (but inferior) to those of 30~Dor are
feasible. Gonz\'alez Delgado \& P\'erez (2000) performed
the analogous experiment of Vacca et~al.\ (1995). The
integrated UV spectrum over a region of tens of pc is 
quite similar to that of 30~Dor. The strong P~Cygni profiles
can be modeled with a standard, or even flatter, IMF and
a foreground dust screen. The dust obscuration derived
from the UV continuum and from the Balmer decrement is
the same within the measurement errors. 



\subsection{Resolved star clusters in galaxies outside the Local Group} 

The closest ``genuine'' starburst galaxies are several times
more distant than M33. Experiments similar to those done for
30~Dor or NGC~604 require HST, both for reasons of aperture size and
flux limits. The requirement is to observe both the UV continuum
and the nebular emission through the same aperture, which must
be small enough to encompass only a single star cluster. No
HST programs have been undertaken with this specific goal in
mind, but several galaxies have archival data which are suitable
for this purpose.
 
NGC~4214 is dwarf starburst
galaxy at a distance of 2.9~Mpc (Ma\'{\i}z-Apell\'aniz, Cieza, \&
MacKenty 2002), which hosts some of the UV brightest star clusters known.
UV spectra of NGC~4214-Ia, a luminous $\sim$4~Myr old
cluster containing WR stars, were
obtained with HST's FOS (Leitherer et~al.\ 1996). The cluster has
a size of less than 5~pc and is supposed to be the ionizing
source of a ring-like H~II region extending about 200~pc from
the center. Its UV spectrum displays characteristic, strong 
P~Cygni profiles, both in \siiv\ and \civ.
The spatial morphology was studied by MacKenty et~al.\
(2000), whose WFPC2 imagery indicates a significant displacement
between the UV and the \Ha\ peak. The strong winds associated
with NGC~4214-Ia have evacuated the surroundings within tens
of pc, creating the large cavity. Obviously this cluster has
cleared its dusty environment and all its stars will experience the same
obscuration at all times. The emission-line strengths outside the
cavity, however, are best modeled with dust and gas uniformly mixed 
(Ma\'{\i}z-Apell\'aniz et~al.\ 1999). The overall morphology
and dust geometry of NGC~4214-Ia agrees with Calzetti's (2001) 
interpretation of the UV-Balmer dichotomy: dust forms a 
foreground screen to the {\em stars} early in the evolution of
the cluster and affects all stellar types the same way. On
the other hand, NGC~4214-Ia does not rule out the time dependent
dust model of Charlot \& Fall (2000), as their model does not
apply to individual clusters.

NGC~5253 at $D=3.2$~Mpc is another nearby starburst galaxy
whose massive star clusters were studied with HST. Tremonti
et~al.\ (2001) obtained HST STIS longslit UV spectra of several
clusters with ages of a few Myr. The spectral morphology 
resembles that of NGC~4214-Ia, suggesting a strong O-star
population. The optical color excesses $E(B-V)$ as derived
from the UV continuum slopes are 0.16, 0.14, and 0.42 for
NGC~5253$-$2, NGC~5253$-$3, and NGC~5253$-$5 in their notation,
respectively. Available WFPC2 narrowband photometry for these
three clusters (Calzetti et~al.\ 1997) allows a determination of
$E(B-V)$ from the \Ha/\Hb\ ratio. While this method is less
secure than a Balmer decrement measurement from a spectrum
obtained from the ground,
it has the advantage of providing a value for $E(B-V)$ 
which is spatially coincident with that derived from the UV
spectrum. A re-remeasurement of the data presented by
Calzetti et~al.\ results in $E(B-V)<0.05$, $<0.05$, and 0.35 for
NGC~5253$-$2, NGC~5253$-$3, and NGC~5253$-$5, respectively. Taken
at face value, the reddening from the Balmer decrement of the 
first two clusters would
be lower than that from the UV slope. However, the
measurement uncertainties are 0.1. Therefore we
conclude that the two sets of values may in fact
 agree; at the very least, it seems
unlikely that the UV values are {\em smaller} than the Balmer
values. As we emphasized for NGC~4214, this test does not
support or exclude either of the two proposed obscuration
models. Rather, it demonstrates that there is no
evidence for the UV-Balmer dichotomy for compact
single-star populations and that UV and emission-line
observations give consistent results with standard
model populations.

M101 is a giant Sc galaxy at a distance of 7.5~Mpc. Rosa
\& Benvenuti (1994) used the FOS of HST to acquire integrated spectra
covering 1200--8000~\AA\ of four giant H~II regions. Spectral
synthesis modeling allowed them to study both the stellar
content and the nebular emission lines. The 
ionizing clusters have ages of $\sim$3~Myr, and their energy
production rate can account for the optical recombination lines.
The UV continuum is predominantly sensitive to reddening, 
with a mean $E(B-V)$ of approximately 0.2. This value agrees
with that derived from the Balmer decrement, which could 
be measured through matching apertures. Again, these 
observations suggest that stars and gas in compact star
clusters with a single stellar population have comparable 
reddening.


\subsection{Star formation in nearby galaxies observed with HUT}


Star-formation regions with sizes of hundreds of pc in 
external galaxies are the real test case for our models. 
The reddening discrepancy between the UV continuum and
the Balmer lines was first discovered in IUE observations
of star-forming galaxies through a $20'' \times 10''$ 
aperture, corresponding to linear areas of order 1~kpc$^2$
(Fanelli et~al.\ 1988). IUE spectra do not have sufficient
spectral resolution and signal-to-noise for the kind of
line-profile analysis required here. HST spectra on the
other hand would be perfectly suitable if they covered
a larger area. All existing HST UV spectral data were taken
through narrow slits with widths $2''$ or smaller so that
they generally sample only a few star clusters (Tremonti
et~al.\ 2002).

UV spectra obtained with the Hopkins Ultraviolet Telescope (HUT)
are the best data set currently in existence to test our
models. HUT provides higher signal-to-noise than IUE and a spectral
resolution of 3--4~\AA, which exceeds that of IUE by
about a factor of two. The HUT aperture is comparable to
that of IUE. Leitherer et~al.\ (2002) analyzed the complete
archival data set of star-forming galaxies observed with HUT.
We will use their sample for a model comparison. Of the 17
galaxies discussed by Leitherer et~al., five objects have
very poor signal-to-noise and/or host an active galactic
nucleus. The remaining 12 galaxies are:
NGC~2403,
IRAS~08339+6517,
NGC~3310,
NGC~3351,
Mrk~1267,
NGC~4214,
NGC~4449,
Mrk~66,
NGC~5194,
NGC~5236,
NGC~5253,
and NGC~6090. The objects span a range of galaxy types and 
properties, but they all have in common a UV spectral energy
distribution indicative of current star formation. Details
of the earlier analysis are given by Leitherer
et al. and are not repeated here. 
We combined the spectra of the 12 galaxies to create a spectral
template of high signal-to-noise. Each of the individual 
normalized spectra was added with the same weight.
We take this template as representative
of a starburst containing multiple star clusters and H~II regions
over an area over $\sim$1~kpc$^2$ and with a mean metallicity of
about solar. The mean age of the 12 regions, as derived by
Leitherer et~al., is about 20~Myr.

The synthetic profiles were smoothed over 5 pixels (corresponding
to 3.5~\AA) to match the HUT resolution and then compared to the
data. No correction for the wavelength zero point error of the data was
applied. The comparison for Si~IV is shown in Fig.~\ref{HUT_Si}.
The solid line in this figure is the average spectrum, and the
shaded area indicates the mean absolute deviation  
$\frac{1}{n}\sum_{i=1}^{n}|f_i(\lambda) - F(\lambda)|$,
where $F(\lambda)$ and $f_i(\lambda)$ denote the mean and the
individual fluxes of the 12 spectra, respectively. While
we do not expect to be able to constrain dust obscuration models
with this line, it serves as an important test
since there are virtually no free parameters, except the IMF.
Fig.~\ref{HUT_Si} suggests excellent agreement. As explained before,
the two narrow Si~IV absorption components are interstellar lines
whose strength and velocity has nothing to do with the {\em broad}
(extending to $\sim$1380~\AA) stellar-wind absorption. The Si~IV wind
emission and the photospheric Si~III and C~III lines at 1417 and
1427~\AA\ agree very well. Given the fact that this is by no means
a tailored model with any fine-tuning,  Fig.~\ref{HUT_Si} is
encouraging.

Next we compare the C~IV profiles in Fig.~\ref{HUT_C}. As in 
Fig.~\ref{HUT_Si}, the
solid line and the shaded area indicate the mean spectrum and the
mean absolute deviation, respectively. The model
parameters are the same as those used to reproduce Si~IV. Apart
from \civ, the other stellar lines in the data are
S~V $\lambda$1502 and a group of unresolved metal lines between
1560 and 1580~\AA. These features agree rather well with the models,
irrespective of the dust obscuration model. Interstellar \siii\ is extremely
broad and deep, a result of its intrinsic strength and the 
degraded HUT resolution due to aperture size effects. This
line is present in the models, but at much reduced strength,
for reasons explained above. The absorption part of C~IV is
heavily contaminated by interstellar absorption as well. In
particular, the absorption-line core is almost certainly
purely interstellar. If higher-resolution data were available,
deblending would be possible, but the HUT data themselves do not permit
a profile fit to the absorption. The emission component, however, provides
clues on the dust geometry. The standard model with no dust 
obscuration is in very good agreement with the mean spectrum, whereas
time dependent obscuration results in too weak emission. On the other
hand, the
model with  $t_{\rm{dust}} = 15$~Myr plotted in  Fig.~\ref{HUT_C}
underpredicts the emission-line flux in comparison with the mean spectrum.
However, the 12 HUT galaxies span a range of parameters,
as suggested by the variation around the mean. While some of the 
scatter is caused by noise, it represents real variations from galaxy
to galaxy as well. 

We caution that the trend in Fig.~\ref{HUT_C} should be considered
merely as suggestive, and definitely not conclusive. The data set
is not rigorously defined, and the quality is marginal for the 
purpose of this study. Taken at face value, the comparison between
models and observations favors {\em on average} similar dust 
obscuration for all stars with ages between 0 and 20~Myr.





\subsection{Lyman-break galaxies} 

Studies of star-forming galaxies at distances of tens of Mpc are
all affected by aperture-size effects. The possibility exists,
and is even likely, that spectroscopic data obtained even through
large apertures may miss a significant fraction of the emitted
luminosity. This is a particular concern for star-formation rates
derived from recombination-line
data which rely on photons which are not created in situ, where the
stars are formed. Galaxies which are small in comparison with the
aperture size are suitable objects to circumvent this problem. The
population of Lyman-break galaxies which has been intensively studied
in recent years (Pettini et~al.\ 2001; Shapley et~al.\ 2001) makes
an ideal test case.

By far the best studied example of a Lyman-break galaxy is MS~1512$-$cB58,
a lensed star-forming galaxy at a redshift of $z=2.7$, whose magnification
factor of $\sim$30 makes it a relatively easy target for a large
telescope (Yee et~al.\ 1996; Pettini et~al.\ 2000). The restframe UV
spectrum of MS~1512$-$cB58 was discussed and modeled by Pettini et al.
The UV lines of MS~1512$-$cB58 bear a striking resemblance to the
lines seen at low redshift and support the assumption that the UV
light is entirely dominated by a population of OB stars with 
characteristic ages of tens of Myr. If so, the UV continuum slope
can be used to estimate the stellar reddening, which turns out
to be $E(B-V)= 0.29$ or 0.24 for either a Calzetti-type or LMC-type
reddening law, respectively. The reddening from the Balmer decrement
is relatively difficult to measure due to the redshift of MS~1512$-$cB58.
Teplitz et~al.\ (2000) obtained near-IR spectroscopy and found 
$E(B-V)= 0.27$, in agreement with the color excess measured from
the UV slope. In this case we would not expect dilution of the
UV P~Cygni profiles for any of our dust obscuration models. 
The comparison with synthetic models performed by 
de~Mello et~al.\ (2000) and Pettini et~al.\ demonstrates
that this is indeed the case. The emission component of \civ\ in
MS~1512$-$cB58 reaches about 15\% above the continuum, which would
be in conflict with any of the time dependent dust obscuration models
discussed here. The other wind lines in MS~1512$-$cB58 are strong as well.

The data for other Lyman-break galaxies are not of high enough 
quality for a similar analysis. Pettini et~al.\ (2001) compared
the star-formation rates from the UV luminosity and from the \Hb\
recombination-line flux and found that there is no significant
offset between the two sets of star-formation rates prior to any
reddening correction. Earlier reports
of differences between the rates derived from the UV and from the Balmer
lines were the result of small number statistics. Unless there is
some fortuitous cancelation of counteracting effects  such as, e.g.,
leakage of radiation, an underlying older population, or a peculiar
star-formation history, the agreement suggests that there is little
differential reddening between the optical and the UV. Since almost
all Lyman-break galaxies have quite low reddening (lower than
MS~1512$-$cB58), any offset between $E(B-V)$ from the UV and from \Hb\
is easily hidden by uncertainties in the calibration of these
star-formation indicators.




\section{Conclusions and directions for further work}

We have presented a new approach to probe the properties of
the most massive, ionizing stars with respect to less massive, 
non-ionizing stars. The emission-line flux of stellar-wind 
P~Cygni profiles is well understood, both in terms of its
astrophysical origin and its systematic behavior. \civ\ in
particular traces the same mass and age range in a stellar
population as the nebular \Ha\ recombination line, with the
advantage of being independent of assumptions on the interstellar
medium. Stellar-wind lines in the UV are by far the best, and most
likely the only stellar lines whose behavior mirrors that of
nebular recombination lines.  No other spectral feature
from either hot or cool massive stars (except for WR stars)
is strong enough and/or
unblended to be easily observable in the spectrum of a
young stellar population. 
The new method can be used to probe the dust obscuration
properties of massive stars for different dust geometries.

The main results of our study are:

1) {\em The equivalent widths of wind lines are rather 
sensitive to selective attenuation of their stellar
contributors, for instance by dust obscuration.}
 The emission part of
the wind lines can be weakened or strengthened with respect
to the surrounding UV continuum for an increased or
decreased strength of the O-star population relative to older,
less massive B stars. The same holds for the absorption
part of the profile but this aspect is less useful in typical
galaxy spectra due to contamination by interstellar absorption.
\civ\ is best suited for this method due to its intrinsic
strength and the dominant dependence on spectral type, and 
therefore mass and age. \nv\ behaves similarly but is compromised
by its proximity to \La\ and the associated uncertainty of the
continuum level. \siiv\ is mostly luminosity dependent and
exhibits less dramatic changes when varying the ionizing relative to the
non-ionizing population.

2) {\em Nebular \civ\ is usually weak, but C~II $\lambda$1335 can
be strong and detectable.}
We computed a set of  photoionization models with parameters 
appropriate for the objects under investigation. The models verified
that at solar metallicity \civ\ is almost certainly purely stellar.
The nebular emission never exceeds more than 10\% of the stellar
emission in \civ\ and is typically much weaker than the C~III $\lambda$1909
line. Nebular C~IV increases in strength
at lower metallicity and with depleted element abundances but remains
weaker than C~III.
Since C~III itself tends to be rather weak in starburst spectra, very 
little nebular C~IV is expected. We found that \cii\ is the strongest
nebular line in the 1200--2000~\AA\ wavelength range. In the 
absence of blending by stellar and interstellar absorption, the C~II
line could be detectable in a typical spectrum of a young population. 

3) {\em Time dependent dust obscuration affects \civ\ differently
for different dust geometries}.
Variations of the ionizing relative to the non-ionizing population
in a mixed (as opposed to a single) stellar population can
be thought of in several ways: IMF variations, a star-formation
history which varies in time over the evolutionary time scale
of O stars, or time dependent dust obscuration affecting very
massive and less massive stars differently. In our study we
concentrated on the third mechanism, but each of them produces
the same effect in the line profiles. We emphasize that we
did not introduce an explicit mass dependence of the obscuration.
The reason for more massive stars being more affected than less 
massive stars is the temporal decrease of the ``effective'' mass of the
contributor to the UV light for a constant star-formation rate.

4) {\em A comparison with large-aperture data favors similar
dust obscuration for ionizing and non-ionizing stars.} 
We demonstrated that time dependent dust obscuration which 
attenuates the UV light from ionizing stars more than that
from non-ionizing stars produces a significant dilution of the
\civ\ profile. This effect can be a powerful diagnostic to
distinguish between alternative models for the dust morphology
in star-forming galaxies. A first application to a sample of
galaxies observed with HUT gives mild evidence in favor of
dust {\em not} selectively obscuring  massive, ionizing stars.
Otherwise the strength of the C~IV emission could only be
understood if non-standard and less plausible assumptions on
the IMF or the star-formation history were made. 

5) {\em There is no observational support for a different 
reddening of stars vs. gas in single stellar populations.}
A review of the existing body of observations suggesting
higher reddening of optical recombination lines than of the
stellar UV continuum was performed. The higher reddening
inferred from recombination lines is the main motivation
behind different dust obscuration scenarios. There is no
evidence for such an effect in the best studied nearby giant
H~II regions and starburst clusters, as expected for single
stellar populations with or without dust obscuration varying
with time. However, these observations virtually rule out
models in which dust obscuration has an explicit dependence
on stellar mass. Similarly, observations of the integrated
stellar populations of MS~1512$-$cB58, the
best studied Lyman-break
galaxy, indicate similar reddening for the UV continuum
and the recombination lines. The only types of observations
suggesting different reddening for the stellar UV continuum
and the nebular lines in the optical are those of kpc-sized 
regions in nearby star-forming galaxies. 

6) {\em The reddening dichotomy is likely the result of
aperture size and spatial morphology.}
The limited body of observations supports dust
obscuration models with little attenuation variation over
the first tens of Myr since the onset of star formation. 
Better quality data are, however, needed for a fully
conclusive answer. If such a test is performed locally,
HST is the instrument of choice for UV sensitivity and spectral resolution.
Due to the small projected slit sizes of HST's spectrographs,
spectral mapping of multiple H~II regions and star clusters
over hundreds of pc is required. Suitable galaxies should not
have one or few bright clusters dominating the integrated light.
Such observations would allow verification  and better 
astrophysical understanding of the HUT results. It seems 
likely that the UV/optical reddening dichotomy is 
merely a result of aperture size and spatial offsets between
the stellar, gaseous, and dust components in star-forming 
regions. The fact that MS~1512$-$cB58, a well studied
Lyman-break galaxy, does not show
this dichotomy supports this hypothesis.






\acknowledgments
We acknowledge funding by NASA from ADP grant NAG5-11023 and
partial support by FAPESP, Brazil. This research 
has made use of the NASA/IPAC Extragalactic Database (NED) which is
operated by the Jet Propulsion Laboratory, California Institute of 
Technology, under contract with the National Aeronautics and Space 
Administration. We thank Christy Tremonti for IDL support and for
making some of her HST data available to us. Max Pettini and
Sueli Viegas provided insightful comments on an earlier version of 
the manuscript. Helpful suggestions from the referee Michael~Rosa 
are gratefully acknowledged.

%\appendix
\newpage

\section*{Appendix}
\section*{Details of the model computations}

Starburst99 has two default modes for the calculation of star-formation
histories: (i)~a single stellar population resulting from an instantaneous
starburst and evolving without star formation after $t = 0$; (ii)~a mixed
population whose stars form continuously at a constant rate independent
of time. From a computational view point, a time dependent dust 
obscuration is identical to a time dependent star-formation rate.
Therefore we varied the star-formation history of a mixed population
in order to simulate the obscuring effect of dust. 

Any mixed population can be approximated by a superposition of
a sufficiently large number of instantaneous bursts covering the
same age range as the mixed population. We first computed spectra 
with an instantaneous burst model at 0.5~Myr intervals. This
resolution proved sufficient for approximating  a smooth, continuous
star formation later on. The resolution of 0.5~Myr refers to the
generation of individual spectra. The internal interpolation scheme
applied to the evolutionary tracks used the default value of 0.1~Myr.
This value is required to cover the short evolutionary phases of
massive stars.   Subsequently we
applied weight factors to each individual burst. The factors were
derived from the dust attenuation law and the normalization 
requirements. The weighted bursts were co-added to compute a
continuous model.

The routine used for the superposition is written in IDL and works
as an add-on to Starburst99. The IDL routine is available from C.L.\
upon request.


\begin{thebibliography}{}
\bibitem{an} Anders, E.\ \& Grevesse, N. 1989, Geochim.\ Cosmochim.\
 Acta, 53, 197  
\bibitem{be} Beck, S.~C., Turner, J.~L., \& Gorjian, V. 2001, AJ,
   122, 1365
\bibitem{bra} Brandner, W., Grebel, E.~K., Barb\'a, R.~H.,
     Walborn, N.~R., \& Moneti, A. 2001, AJ, 122, 858
\bibitem{bre} Bresolin, F.\ \& Kennicutt, R.~C. 2002, ApJ, in press
\bibitem{br} Bruzual, G. 2001, Ap\&SS, 277, 221
\bibitem{bc} Bruzual, G.\ \& Charlot, S. 1993, ApJ, 405, 538
\bibitem{ca0} Calzetti, D. 1997, AJ, 113, 162
\bibitem{ca} ------. 2001, PASP, 113, 1449
\bibitem{CA2} Calzetti, D., Kinney, A.~L. \& Storchi-Bergmann, T. 
    1994, ApJ, 429, 582
\bibitem{ca10} Calzetti, D., Meurer, G.~R., Bohlin, R.~C., 
   Garnett, D.~R., Kinney, A.~L., Leitherer, C., \& 
    Storchi-Bergmann, T. 1997, AJ, 114, 1834
\bibitem{cf} Charlot, S.\ \& Fall, S.~M. 2000, ApJ, 539, 718
\bibitem{ch} Chiosi, C., Vallenari, A., Bressan, A., Deng, L.,
    \& Ortolani, S. 1995, A\&A, 293, 710
\bibitem{co} Contini, M.\ \& Viegas, S.~M. 2001, ApJS, 137, 75
\bibitem{dek} de Koter, A., Heap, S.~R., \& Hubeny, I. 1998, ApJ, 509,
   879
\bibitem{dem} de Mello, D.~F., Leitherer, C., \& Heckman, T.
   M. 2000, ApJ, 530, 251
\bibitem{dw} Dwek, E. 1998, ApJ, 501, 643
\bibitem{fa} Fanelli, M.~N., O'Connell, R.~W., \& Thuan, T.~X.
   1988, ApJ, 334, 665
\bibitem{fi} Fioc, M.\ \& Rocca-Volmerange, B. 1997, A\&A, 326, 950
\bibitem{gas} Gasda, S.\ \& Leitherer, C. 1996, in From Stars to 
        Galaxies: The Impact of Stellar Physics on Galaxy Evolution,
       eds.\ C.~Leitherer, U.~Fritze-von-Alvensleben, \& J.~Huchra
       (San Francisco: ASP), 451 
\bibitem{fitz} Fitzpatrick, E.~L.\ \& Savage, B.~D. 1984, ApJ, 279, 578
\bibitem{gen} Genzel, R., Lutz, D., Sturm, E., Egami, E., Kunze, D.,
  Moorwood, A.~F.~M., Rigopoulou, D., Spoon, H.~W.~W., Sternberg, A.,
  Tacconi-Garman, L.~E., Tacconi, L., \& Thatte, N. 1998, ApJ, 498, 579
\bibitem{go97} Gonz\'alez Delgado, R.~M., Leitherer, C., \& Heckman,
    T.~M. 1997, ApJ, 489, 601
\bibitem{go99} ------. 1999, ApJS, 125, 489
\bibitem{gp} Gonz\'alez Delgado, R.~M.\ \& P\'erez, E. 2000, 
     MNRAS, 317, 64
\bibitem{gor} Gordon, K.~D., Misselt, K.~A., Witt, A.~N.,
    \& Clayton, G.~C. 2001, ApJ, 551, 269
\bibitem{grue} Gruenwald, R.~B.\ \& Viegas, S.~M. 1992, ApJS, 78, 153
\bibitem{hec} Heckman, T.~M. 2001, in Gas and Galaxy Evolution, eds.\ J.~E.\
    Hibbard, M.~Rupen, \& J.~H.\ van Gorkom (San Francisco: ASP), 345
\bibitem{hec1} Heckman, T.~M., Robert, C., Leitherer, C., Garnett, D.~R.,
   \&  van der Rydt, F. 1998, ApJ, 503, 646
\bibitem{kien} Kennicutt, R.~C. 1984, ApJ, 287, 116
\bibitem{kew} Kewley, L.~J., Dopita, M.~A., Sutherland, R.~S.,
    Heisler, C.~A., \& Trevena, J. 2001, ApJ, 556, 121
\bibitem{ll} Leitherer, C., Le\~ao, J.~R.~S., Heckman, T.~M.,
     Lennon, D.~J., Pettini, M., \& Robert, C. 2001, ApJ, 550, 724
\bibitem{lll} Leitherer, C., Li, I.-H., Calzetti, D., \& Heckman, T.~M.
       2002, ApJS, 140 (in press)
\bibitem{lr} Leitherer, C., Robert, C., \& Heckman, T.~M. 1995, ApJS, 99, 173
\bibitem{le} Leitherer, C., Schaerer, D., 
    Goldader, J.~D., Gonz\'alez Delgado, R.~M., Robert, C., Foo Kune,
    D., de Mello, D., Devost, D., \& Heckman, T.~M. 1999, \apjs, 123, 3
\bibitem{leit} Leitherer, C., Vacca, W.~D., Conti, P.~S., Filippenko, A.~V.,
    Robert, C., \& Sargent, W.~L.~W. 1996, ApJ, 465, 717
\bibitem{mac} MacKenty, J.~W., Ma\'{\i}z-Apell\'aniz, J., Pickens, C.~E.,
     Norman, C.~A., \& Walborn, N.~R. 2000, AJ, 120, 3007 
\bibitem{mae1} Maeder, A. 1999, in IAU Symp.~193,
     Wolf-Rayet Phenomena in Massive Stars and Starburst Galaxies,
     eds.\ K.~A.\ van der Hucht, G.~Koenigsberger, \& P.~R.~J.\ Eenens
     (San Francisco: ASP), 177
\bibitem{mae} Maeder, A.\ \& Meynet, G. 2000, ARAA, 38, 143
\bibitem{maiz}  Ma\'{\i}z-Apell\'aniz, J., Cieza, L., \& MacKenty, J.~W.\
     2002, AJ, in press
\bibitem{mai} Ma\'{\i}z-Apell\'aniz, J., Mu\~noz-Tu\~n\'on, C., 
    Tenorio-Tagle, G., \& Mas-Hesse, J.~M. 1999, A\&A, 343, 64
\bibitem{mas} Mas-Hesse, J.~M.\ \& Kunth, D. 1991, A\&AS, 88, 399
\bibitem{mass} Massey, P.\ \& Hunter, D.~A. 1998, ApJ, 493, 180
\bibitem{mat} Mathis, J.~S., Chu, Y.-H., \& Peterson, D.~E. 1985, ApJ, 292,
    155
\bibitem{pet1} Pettini, M., Shapley, A.~E., Steidel, C.~C., Cuby, J.-G.,
     Dickinson, M., Moorwood, A.~F.~M., Adelberger, K.~L., \&
     Giavalisco, M. 2001, ApJ, 554, 981
\bibitem{pet2} Pettini, M., Steidel, C.~C., Adelberger, K.~L.,
        Dickinson, M. \& Giavalisco, M. 2000, ApJ, 528, 96 
\bibitem{po} Poggianti, B.~M., Smail, I., Dressler, A., Couch, W.~J., 
        Barger, A.~J., Butcher, H., Ellis, R.~S., \& Oemler, A.~J.
        1999, ApJ, 518, 576
\bibitem{rai} Raimann, D., Storchi-Bergmann, T., Bica, E.,
        Melnick, J., \& Schmitt, H. 2000, MNRAS, 316, 559
\bibitem{rie} Rieke, G.~H., Loken, K., Rieke, M.~J., \& Tamblyn, P. 1993,
    ApJ, 412, 99
\bibitem{rob} Robert, C., Leitherer, C., \& Heckman, T.~M. 1993, ApJ, 418, 749
\bibitem{ros} Rosa, M.~R.\ \& Benvenuti, P. 1994, A\&A, 291, 1
\bibitem{sal} Salasnich, B., Bressan, A., \& Chiosi, C. 1999, A\&A, 342, 131
\bibitem{sch} Schaerer, D., Guseva, N.~G., Izotov, Yu.~I., \&
        Thuan, T.~X. 2000, A\&A, 362, 53 
\bibitem{sc} Schmitt, H.~R., Kinney, A.~L., Calzetti, D., \&
     Storchi Bergmann, T. 1997, AJ, 114, 592
\bibitem{sha} Shapley, A.~E., Steidel, C.~C., Adelberger, K.~L., Dickinson, M.,
         Giavalisco, M., \& Pettini, M. 2001, ApJ, 562, 95
\bibitem{sh} Shields, J.~C.\ \& Kennicutt, R.~C. 1995, ApJ, 454, 807
\bibitem{si} Silva, L., Granato, G.~L., Bressan, A., \& Danese, L.
     1998, ApJ, 509, 103
\bibitem{sta} Stasi\'nska, G., Schaerer, D., \& Leitherer, C. 
     2001, A\&A, 370, 1
\bibitem{st} Sternberg, A. 1998, ApJ, 506, 721
\bibitem{tep} Teplitz, H.~I., McLean, I.~S., Becklin, E.~E., Figer, D.~F.,
         Gilbert, A.~M., Graham, J.~R., Larkin, J.~E., Levenson, N.~A., \&
         Wilcox, M.~K. 2000, ApJ, 533, L65
\bibitem{ti} Tinsley, B.~M. 1968, ApJ, 151, 574
\bibitem{tre} Tremonti, C.~A., Calzetti, D., Heckman, T.~M., \& 
     Leitherer, C. 2002, ApJS, submitted
\bibitem{tr} Tremonti, C.~A., Calzetti, D., Leitherer, C., \& 
     Heckman, T.~M. 2001, ApJ, 555, 322
\bibitem{vac} Vacca, W.~D., Robert, C., Leitherer, C., \& Conti,
      P. S. 1995, ApJ, 444, 647
\bibitem{walb} Walborn, N.~R. 1991, in Massive Stars in Starbursts,
     eds.\ C.~Leitherer, N.~R.\ Walborn, T.~M.\ Heckman, \& C.~A.\ Norman
     (Cambridge: CUP), 145 
\bibitem{wal1} Walborn, N.~R., Nichols-Bohlin, J., \& Panek, R.~J.
       1985, International Ultraviolet Explorer Atlas of O-type
       Spectra from 1200 to 1900~\AA\ (Washington: NASA)
\bibitem{wp} Walborn, N.~R.\ \& Panek, R.~J. 1984, ApJ, 280, L27 
\bibitem{wal2} Walborn, N.~R., Parker, J.~W., \& Nichols, J.~S. 
       1995, International Ultraviolet Explorer Atlas of B-type 
       Spectra from 1200 to 1900~\AA\ (Washington: NASA) 
\bibitem{wi} Willis, A.~J. 1999, in IAU Symp.~193,
     Wolf-Rayet Phenomena in Massive Stars and Starburst Galaxies,
     eds.\ K.~A.\ van der Hucht, G.~Koenigsberger, \& P.~R.~J.\ Eenens
     (San Francisco: ASP), 1
\bibitem{wo} Worthey, G. 1994, ApJS, 95, 105
\bibitem{yee} Yee, H.~K.~C., Ellingson, E., Bechtold, J., Carlberg, R.~G.,
   \& Cuillandre, J.-C. 1996, AJ, 111, 1783


\end{thebibliography}

%\clearpage

\begin{figure}
\centering
\includegraphics[width=12.8cm,width=9.6cm,angle=90]{f1.eps}
%\psfig{figure=f1.eps,angle=90,width=16cm,height=12cm}
\caption{\label{10Myr} \civ\ for a 
model with constant star formation at $t=10$~Myr. Solid: 
$t_{\rm{dust}} = \infty$ (no time dependent dust obscuration);
dotted: $t_{\rm{dust}} = 5$~Myr;
dashed: $t_{\rm{dust}} = 15$~Myr;
dash-dotted: $t_{\rm{dust}} = 50$~Myr. The profiles for
$t_{\rm{dust}} = 15$ and 50~Myr are almost identical. Further
parameters: $Z$~=~\Zs; $\alpha = 2.35$; \mup~=~100~\Ms;
\mlow~=~1s.}
\end{figure}
%\clearpage
\begin{figure}
\centering
\includegraphics[width=12.8cm,width=9.6cm,angle=90]{f2.eps}
%\psfig{figure=f2.eps,angle=90,width=16cm,height=12cm}
\caption{\label{20Myr} Same as 
Fig.~\ref{10Myr}, but at age $t=20$~Myr.}
\end{figure}

%\clearpage
\begin{figure}
\centering
\includegraphics[width=12.8cm,width=9.6cm,angle=90]{f3.eps}
%\psfig{figure=f3.eps,angle=90,width=16cm,height=12cm}
\caption{\label{50Myr} Same as 
Fig.~\ref{10Myr}, but at age $t=50$~Myr.}
\end{figure}

%\clearpage

\begin{figure}
\centering
\includegraphics[width=12.8cm,width=9.6cm,angle=90]{f4.eps}
%\psfig{figure=f4.eps,angle=90,width=16cm,height=12cm}
\caption{\label{SiIV} Same as 
Fig.~\ref{20Myr}, but for \siiv.}
\end{figure}

%\clearpage
\begin{figure}
\centering
\includegraphics[width=12.8cm,width=9.6cm,angle=90]{f5.eps}
%\psfig{figure=f5.eps,angle=90,width=16cm,height=12cm}
\caption{\label{SiIV_5} \siiv\ at age 
$t=5$~Myr. Solid: no dust obscuration; dotted: dust 
obscuration with $t_{\rm{dust}} = 5$~Myr. Si~IV is
{\em enhanced} for this particular set of parameters.} 
\end{figure}

%\clearpage


\begin{figure}
\centering
\includegraphics[width=12.8cm,width=9.6cm,angle=90]{f6.eps}
%\psfig{figure=f6.eps,angle=90,width=16cm,height=12cm}
\caption{\label{IMF} \civ\ profiles
for a population without time dependent dust obscuration (solid)
and for a population with $t_{\rm{dust}} = 15$~Myr and a flatter
IMF with $\alpha = 1.7$ (dashed).}
\end{figure}

%\clearpage
\begin{figure}
\centering
\includegraphics[width=12.8cm,width=9.6cm,angle=90]{f7.eps}
%\psfig{figure=f7.eps,angle=90,width=16cm,height=12cm}
\caption{\label{HUT_Si} Average Si~IV profile
of 12 star-forming galaxies observed with HUT (solid), with the
mean absolute deviation indicated by the shaded area. The
observations are compared to a synthetic profile of age 20~Myr
(dashed) with
standard population parameters and no dust obscuration.}
\end{figure}

%\clearpage
\begin{figure}
\centering
\includegraphics[width=12.8cm,width=9.6cm,angle=90]{f8.eps}
%\psfig{figure=f8.eps,angle=90,width=16cm,height=12cm}
\caption{\label{HUT_C} Average C~IV profile
of 12 star-forming galaxies observed with HUT (solid), with the
mean absolute deviation indicated by the shaded area. The
observations are compared to synthetic profiles of age 20~Myr with
standard population parameters and no dust obscuration (dashed)
and dust obscuration with $t_{\rm{dust}} = 15$~Myr (dotted).}
\end{figure}

\clearpage

\begin{landscape}

\begin{deluxetable}{ccccccc}
\tablewidth{0pt}
\tablecaption{Photoionization model parameters}
%\begin{center}
\tablehead{
\colhead{Model} &
\colhead{$Z$} &
\colhead{$n_{\rm e}$} &
\colhead{IMF} &
\colhead{$F(H\beta)$} &
\colhead{$I_{\rm{CIII}}/I_{\rm{H}\beta}$} &
\colhead{$\log U$} \\
&& \colhead{(cm$^{-3}$)} &&
\colhead{(erg s$^{-1}$ cm$^{-2}$)}
}
\startdata
 1& \Zs & 100 & Salpeter   &$8.19 \times 10^{-10}$& $3.80 \times 10^{-2}$& 3.8 \\
 2& \Zs & 100 &$\alpha=1.5$&$3.07 \times 10^{-9\phn}$ & $4.52 \times 10^{-2}$& 4.3 \\
 3& \Zs, depleted	& 100&Salpeter&$8.25\times 10^{-10}$&$3.70 \times 10^{-2}$&3.8 \\
 4& \Zs, depleted	& 300&Salpeter&$8.22\times 10^{-10}$&$4.70 \times 10^{-2}$&3.3 \\	
 5& \Zs, depleted	&\phn30 &Salpeter&$8.23\times 10^{-10}$&$3.29 \times 10^{-2}$&4.3 \\
 6& $\frac{1}{4}$\Zs, depleted& 100&Salpeter&$8.14 \times 10^{-10}$&$3.63 \times 10^{-1}$&3.8 \\
 7& \Zs, depleted	& 100& $\alpha=1.5$&$3.09 \times 10^{-9\phn}$&$1.46 \times 10^{-2}$&4.3 \\
\enddata
%\end{center}
\end{deluxetable}

%\clearpage
\begin{deluxetable}{cccccccccccccc}
\tablewidth{0pt}
\tablecaption{Photoionization model results}
%\begin{center}
\tablehead{
\multicolumn{2}{c}{Model 1} &
\multicolumn{2}{c}{Model 2} &
\multicolumn{2}{c}{Model 3} &
\multicolumn{2}{c}{Model 4} &
\multicolumn{2}{c}{Model 5} &
\multicolumn{2}{c}{Model 6} &
\multicolumn{2}{c}{Model 7} \\
\colhead{Line} &
\colhead{$I$} &
\colhead{Line} &
\colhead{$I$} &
\colhead{Line} &
\colhead{$I$} &
\colhead{Line} &
\colhead{$I$} &
\colhead{Line} &
\colhead{$I$} &
\colhead{Line} &
\colhead{$I$} &
\colhead{Line} &
\colhead{$I$} 
}
\startdata
OV 1215&3.35& OV 1215&4.23& OV 1215&7.43& OV 1215&8.23&OV 1215&6.53& OV 1215&8.35& OV 1215&8.51\\
OV 1218&2.57& OV 1218&3.12& OV 1218&2.31& OV 1218&2.26&OV 1218&2.02& NV	1240&3.42& OV 1218&2.45\\
NV 1240&1.22& NV 1240&1.40& NV 1240&3.24& NV 1240&3.49&NV 1240&3.04&CII 1335&4.10& NV 1240&3.44\\
CII 1335&\llap{38}0\phd\phn\phn&CII 1335&\llap{32}7\phd\phn\phn&CII 1335&\llap{19}1\phd\phn\phn&CII 1335&\llap{14}9\phd\phn\phn&CII 1335&\llap{22}0\phd\phn\phn&OIV 1394&2.31&CII 1335&\llap{13}3\phd\phn\phn\\
OV 1371&0.74\rlap{0}&OV 1371&0.90\rlap{7}&OV 1371&0.68\rlap{6}&OV 1371&0.67\rlap{7}&OV 1371&0.59\rlap{6}&SiIV 1397&1.18&OV 1371&0.73\rlap{8}\\
OIV 1394&0.30\rlap{6}&OIV 1394&0.30\rlap{5}&OIV 1394&1.92&OIV 1394&1.95&OIV 1394&2.00&SIV 1400&1.26&OIV 1394&1.77\\
SiIV 1397&0.09\rlap{56}&SiIV 1397&0.10\rlap{5}&OIV 1401&0.58\rlap{9}&OIV 1401&0.61
\rlap{9}&OIV 1401&0.60\rlap{8}&OIV 1405&0.85\rlap{1}&OIV 1401&0.54\rlap{3}\\
OIV 1401&0.09\rlap{38}&OIV 1401&0.09\rlap{35}&OIV 1405&0.70\rlap{5}&OIV 1405&0.71
\rlap{1}&OIV 1405&0.74\rlap{2}&NIV 1483&3.25&OIV 1405&0.65\rlap{2}\\
OIV 1405&0.11\rlap{3}&OIV 1405&0.11\rlap{2}&NIV 1483&0.82\rlap{2}&NIV 1483&0.82
\rlap{6}&NIV 1483&0.875&NIV 1483&5.40&NIV 1483&0.728\\
NIV 1483&0.13\rlap{5}&NIV 1483&0.07\rlap{54}&NIV 1483&1.36&NIV 1483&1.37&NIV 1483&1.45&NIV 1486&2.14&NIV 1483&1.21\\
NIV 1487&0.61\rlap{3}&NIV 1483&0.12\rlap{5}&NIV 1486&0.54\rlap{1}&NIV 1486&0.54\rlap{7}&NIV 1486&0.57\rlap{8}&CIV 1549&\llap{9}3.9\phn&NIV 1486&0.47\rlap{9}\\
CIV 1549&4.18&NIV 1487&0.54\rlap{2}&NIV 1487&0.46\rlap{8}&CIV 1549&\llap{2}3.4\phn&NIV 1487&0.54\rlap{4}&HeII 1640&\llap{1}4.4\phn&CIV 1549&\llap{2}1.1\phn\\
NeV 1575&0.11\rlap{3}&CIV 1549&4.32&CIV 1549&\llap{2}2.8\phn&HeII 1640&\llap{10}3\phd\phn\phn&CIV 1549&\llap{2}3.5\phn&OIII 1663&\llap{2}9.2\phn&NeV 1575&0.39\rlap{8}\\
HeII 1640&\llap{12}7\phd\phn\phn&HeII 1640&\llap{11}7\phd\phn\phn&NeV 1575&0.35\rlap{4}&OIII 1663&4.72&NeV 1575&0.31\rlap{0}&SIII 1729&0.89\rlap{5}&HeII 1640&\llap{10}9\phd\phn\phn\\
OIII 1663&0.24\rlap{2}&OIII 1663&0.37\rlap{3}&HeII 1640&\llap{13}3\phd\phn\phn&NIII 1744&1.74&HeII 1640&\llap{15}0\phd\phn\phn&NIII 1744&7.41&OIII 1663&4.10\\
NIV 1719&0.17\rlap{6}&NIV 1719&0.15\rlap{6}&OIII 1663&2.95&NIII 1750&0.68\rlap{1}&OIII 1663&2.52&NIII 1750&2.40&NIII 1744&1.49\\
NIII 1744&0.11\rlap{9}&NIII 1744&0.17\rlap{3}&NIII 1744&0.95\rlap{4}&NIII 1752&0.57\rlap{7}&NIII 1744&0.75\rlap{4}&NIII 1752&2.64&NIII 1750&0.49\rlap{5}\\
NeIII 1815&1.21&NeIII 1815&1.07&NeIII 1815&1.51&NeIII 1815&1.31&NeIII 1815&1.64&NIII 1754&1.02&NIII 1752&0.52\rlap{5}\\
SiIII 1892&1.64&SiIII 1892&1.43&SiIII 1892&1.24&SiIII 1892&1.72&SiIII 1892&1.16&SiIII 1892&4.96&NeIII 1815&1.22\\
CIII 1907	&\llap{10}0\phd\phn\phn&CIII 1907&\llap{10}0\phd\phn\phn&CIII 1907&\llap{10}0\phd\phn\phn&CIII 1907&\llap{10}0\phd\phn\phn&CIII 1907&\llap{10}0\phd\phn\phn&CIII 1907&\llap{10}0\phd\phn\phn&CIII 1907&\llap{10}0\phd\phn\phn\\
\enddata
%\end{center}
%\normalsize
\tablenotetext{~}{Note: Line intensities are relative to C~III $\lambda$1907, with 
$I_\mathrm{CIII} = 100$.}
\end{deluxetable}
\end{landscape}
\end{document}