Ricardo P. Schiavon
A major uncertainty in the spectroscopic dating of extragalactic globular clusters concerns the degenerate effect that age and horizontal branch morphology have on the strength of Balmer lines. In this {Letter} we show that the ratio between the equivalent widths of \hdf\ and \hb\ is far more sensitive to horizontal branch morphology than to age, thus making it possible to break the degeneracy. We show that it is possible to distinguish intermediate-age globular clusters from those whose Balmer lines are strengthened by the presence of blue horizontal branch stars, {purely on the basis of the clusters' integrated spectra}. The degeneracy between age and horizontal branch morphology can be lifted with \hb\ and \hdf\ line strengths from spectra with S/N $\simgreat$ 30/${\AA}$, which is typical of current studies of integrated spectroscopy of extragalactic globular clusters. %***[NOTE: I THINK THIS LAST SENTENCE IS UNNECESSARY, I.E., %IMPLIED IN PREVIOUS ONE... This indicator might be very useful for the %determination of spectroscopic ages of distant GCs for which color %magnitude diagrams cannot be obtained.] ***
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\newcommand{\myemail}{ripisc@virginia.edu}
\newcommand \hdf {$H\delta_F$}
\newcommand \hb  {$H\beta$}
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\shorttitle{Blue Horizontal Branch Stars in Globular Clusters}
\shortauthors{Schiavon et al.}


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%% \title{A Method to Identify the Signature of Blue Horizontal Branch
%% Stars in the Integrated Spectra of Globular Clusters
\title{The Identification of Blue Horizontal
Branch Stars in the Integrated Spectra of Globular
Clusters\footnotemark[1]\footnotetext[1]{Based on observations
collected with the CTIO 4m Blanco telescope}} %SC


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\author{Ricardo P. Schiavon}
\affil{Astronomy Department, University of Virginia, P.O. Box 3818, 
Charlottesville, VA 22903-0818}
\email{ripisc@virginia.edu}

\author{James A. Rose}
\affil{Department of Physics and Astronomy, CB 3255, University of North
Carolina, Chapel Hill, NC 27599}
\email{jim@physics.unc.edu}

\author{St\'ephane Courteau}
\affil{Department of Physics, Queen's University, Kingston, ON Canada K7L 3N6}
\email{courteau@physics.queensu.ca}

\author{Lauren A. MacArthur}
\affil{Department of Physics \& Astronomy, University of British Columbia, 
6224 Agricultural Road, Vancouver, BC Canada V6T 1Z1}
\email{lauren@physics.ubc.ca}



%% Notice that each of these authors has alternate affiliations, which
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%\altaffiltext{1}{Visiting Astronomer, Cerro Tololo Inter-American Observatory.
%CTIO is operated by AURA, Inc.\ under contract to the National Science
%Foundation.}
%\altaffiltext{2}{Society of Fellows, Harvard University.}
%\altaffiltext{3}{present address: Center for Astrophysics,
%    60 Garden Street, Cambridge, MA 02138}
%\altaffiltext{4}{Visiting Programmer, Space Telescope Science Institute}
%\altaffiltext{5}{Patron, Alonso's Bar and Grill}


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\begin{abstract}
A major uncertainty in the spectroscopic dating of extragalactic globular
clusters concerns the degenerate effect that age and horizontal
branch morphology have on the strength of Balmer lines. In this {\it
Letter} we show that the ratio between the equivalent widths of \hdf\
and \hb\ is far more sensitive to horizontal branch morphology than to
age, thus making it possible to break the degeneracy.  We show that it is
possible to distinguish intermediate-age globular clusters from those whose
Balmer lines are strengthened by the presence of blue horizontal branch
stars, {\it purely on the basis of the clusters' integrated spectra}.
The degeneracy between age and horizontal branch morphology can be lifted
with \hb\ and \hdf\ line strengths from spectra with S/N $\simgreat$
30/${\rm\AA}$, which is typical of current studies of integrated
spectroscopy of extragalactic globular clusters.

%***[NOTE: I THINK THIS LAST SENTENCE IS UNNECESSARY, I.E.,
%IMPLIED IN PREVIOUS ONE... This indicator might be very useful for the 
%determination of spectroscopic ages of distant GCs for which color 
%magnitude diagrams cannot be obtained.] ***

\end{abstract}

%% Keywords should appear after the \end{abstract} command. The uncommented
%% example has been keyed in ApJ style. See the instructions to authors
%% for the journal to which you are submitting your paper to determine
%% what keyword punctuation is appropriate.

\keywords{stars: evolution; stars: horizontal-branch; globular clusters:
general; galaxies: evolution; galaxies: stellar content}


%% From the front matter, we move on to the body of the paper.
%% In the first two sections, notice the use of the natbib \citep
%% and \citet commands to identify citations.  The citations are
%% tied to the reference list via symbolic KEYs. The KEY corresponds
%% to the KEY in the \bibitem in the reference list below. We have
%% chosen the first three characters of the first author's name plus
%% the last two numeral of the year of publication as our KEY for
%% each reference.

\section{Introduction}

In the last decade, the use of equivalent widths (EWs) of Balmer
absorption lines in the integrated spectra of old stellar populations
(SPs) has become a standard procedure for estimating their light-weighted
mean ages (Worthey 1994). The rationale of the method is to utilize the
sensitivity of Balmer lines to the temperature and luminosity of turn-off
stars, which depend strongly on the age of a SP.  However, a well known
caveat is the existence of old hot stars that may be sufficiently bright
and numerous to boost the EWs of Balmer lines, mimicking the signature
of young stars in the integrated spectra of old SPs. That is the case
of blue horizontal branch (BHB) stars, which are present in such old
systems as Galactic globular clusters (GCs) with a range of metallicities
(see Moehler 2001 for a review), the halo and bulge fields (e.g. Kinman
et al.\ 2000), the cores of Local Group galaxies (Brown et al.\ 1998;
Brown et al.\ 2000; Brown 2003), and old Galactic clusters like NGC~6791
(e.g. Peterson \& Green 1998). So far the only way of unequivocally
distinguishing between an old SP that hosts BHB stars from a truly young
SP is by construction of color-magnitude diagrams (CMDs), which of course
can only be achieved for nearby systems, within the Local Group.

Theoretical studies (Freitas Pacheco \& Barbuy 1995; Lee et al.\ 2000;
Maraston \& Thomas 2000) have shown that the impact of BHB stars on the
integrated spectra of old SPs can be significant, and in particular it may
affect the ages inferred from the EWs of Balmer lines. However, in view
of the lingering theoretical uncertainties regarding the morphology of
the horizontal branch (HB) and its connection with mass loss on the red
giant branch (Catelan 2000), an accurate assessment of the contribution
of BHB stars to the integrated optical light of old SPs still depends
on the availability of CMDs.

In this {\it Letter} we present the first results of an ongoing project
to calibrate our SP synthesis models with high quality integrated spectra
of star clusters. While this has already been attempted by a number of
groups (Beasley et al.\ 2002; Schiavon et al. 2002a,b; Puzia et al.\
2002; Maraston et al.\ 2003; Leonardi \& Rose 2003), our work benefits
from the combination of a homogeneous set of high S/N integrated GC
spectra newly collected by our group, with the HST CMDs of several tens
of Galactic GCs, most of them reaching the GC turn-off, by Piotto et al.\
(2002).  We demonstrate the ability of well-calibrated models, combined
with high S/N spectra, to single out the contribution of BHB stars to the
integrated light of GCs. We use a combination of an Fe-sensitive index
and an index comprised of the ratio of two Balmer line EWs to uniquely
constrain the BHB contribution to the integrated light. This enables us
to distinguish, {\it purely on the basis of their integrated spectra},
truly intermediate-age or young clusters from those which are old,
but whose Balmer lines are strengthened by the contaminating light of
BHB stars. We envisage a direct application of our method to studies
of extragalactic GC systems, where the determination of GC ages and
metal abundances can yield insight into the star formation and merger
histories of the host galaxies (e.g. Cohen, Blakeslee \& C\^ot\'e 2003;
Larsen et al. 2003; Hempel et al.\ 2003).

%This {\it Letter} is presented as follows: In \S\ref {obs} we briefly
%describe our observations, and we consider in \S\ref {res} the effect
%of BHBS on the EWs of Balmer lines.  Our main result, that
%the signature of BHB stars in the integrated spectra of GCs can be well
%constrained from a combination of an iron index with an index composed
%of the ratio of two Balmer line EWs, is presented and discussed
%in \S\ref {detect}.  An epilogue and look towards future applications
%are offered in \S\ref{final}.

\section{Observations} \label{obs}

We collected integrated spectra for 40 Galactic GCs with the
R-C spectrograph on the 4m Blanco telescope at CTIO in April 2003.
We scanned a 5\arcmpt5-long slit across the core diameter of each target
GC. For the GCs located towards the bulge, on-target exposures were
interspersed with raster scans of adjacent sky regions $\sim$5\arcmin\
away from the GC centers. The instrumental setup consisted of grating
KPGL1, with 632 l/mm, and a Loral CCD with 3k $\times$ 1k 15$\mu$m-sized
pixels. The resulting spectra cover the range 3365--6435${\rm\AA}$
with 2.8${\rm\AA}$ FWHM resolution. The exposure times were scaled
to yield spectra with S/N $\simgreat$ 100 per resolution element. Data
reduction used standard IRAF routines for longslit spectra. Final
1D integrated spectra were obtained by extracting an aperture covering
the core diameter along the slit direction. Since the exposures were
trailed over a core diameter perpendicular to the slit, the resulting
1D spectra are representative of the stellar content in the cores of the
GCs. The spectra were flux calibrated using observations of flux
standards taken throughout the observing run. The 1D integrated spectra
had their resolution degraded to match that of the Lick/IDS system,
and EWs of absorption lines were measured following the definitions
of Worthey et al.\ (1994) and Worthey \& Ottaviani (1997). Consistency
with the Lick/IDS index system was achieved by comparing EWs measured in
spectra of standard Lick/IDS stars, taken throughout the observing run,
with standard values from Worthey et al.\ (1994).


%Sky subtraction was achieved by extracting
%the same region on the separate sky frame as was extracted on the cluster
%frame, hence a direct subtraction was made.  


\section{Effect of Blue Horizontal Branch Stars on Balmer Lines}
\label{res}

%% In this analysis, we only employ a supsample of our GC spectra. Because we
%% want to highlight the effects on model predictions due to HB morphology
%% differences, we exclude from the analysis GCs with [Fe/H] $\simgreat$
%% --0.5, because all the Galactic GCs above that metallicity have similar,
%% stubby red HBs, so that they would not add much to our discussion of the
%% effects of BHBS on absorption line indices. We also limit the analysis
%% to GCs with more extreme HB morphologies, for clarity, and in order to
%% enhance the impact of HB morphology on the absorption line indices.
%%SC What is a "stubby" red HBs?  Here's a rewording of this section:

In order to highlight the impact of HB morphology on the integrated light,
we restrict our analysis to GCs with [Fe/H] $\simless -0.5$, because at
higher metallicities all GCs have similar red HBs, therefore not adding
much to our discussion. A plot of \hdf\ vs. [Fe/H] for this sub-sample is
shown in Figure \ref{fig1}. For most of the GCs, [Fe/H] values were taken
from Kraft \& Ivans (2003).  Those values were supplemented, whenever
needed, by data from Carretta \& Gratton (1997), Idiart, Th\'evenin \&
Freitas Pacheco (1997), and Harris (1996). Throughout this {\it Letter}
GCs are subdivided in terms of HB morphology, with $x_{HB}={B-R\over
B+V+R}$, as follows: GCs with {\it mostly blue} HBs ($x_{HB} > +0.4$),
GCs with {\it mostly red} HBs (--$0.9 < x_{HB} < -0.4$), GCs with {\it
strictly red} HBs ($x_{HB}$ $<$ --0.9), and GCs with intermediate HB
morphologies (--$0.4 < x_{HB} <$ +0.4). Values of $x_{HB}$ were taken
from Borkova \& Marsakov (2000). In all the figures, a few landmark
clusters are highlighted with special symbols (see caption of Figure
\ref{fig1}). Those are 47 Tuc, a metal-rich GC with [Fe/H] $\sim$ --0.8
and a red HB; M67, a $\sim3.5$ Gyr old, solar metallicity, open cluster
(data from Schiavon, Caldwell \& Rose 2004); NGC~6388 and NGC~6441, the
highest metallicity ([Fe/H] $\sim$ --0.6) Galactic GCs known to host BHB
stars (Rich et al.\ 1997), and NGC~5904, a metal poor GC with [Fe/H]
$\sim$ --1.27 (Kraft \& Ivans 2003) and whose HB has a sizeable blue
extension (Piotto et al.\ 2002).  Our sub-sample spans a wide range in
metallicity and HB morphology; most notably, GCs with --1.5 $\simless$
[Fe/H] $\simless$ --0.9 are strongly affected by HB morphology, in the
sense that \hdf\ in strictly blue-HB GCs is stronger by $\sim$1 ${\rm\AA}$
than in those with strictly red HBs at the same [Fe/H].

\begin{figure}
\plotone{f1.eps}
\caption{The sample of GCs analysed in this paper: GCs with mostly blue
HBs are depicted as asterisks, open circles represent those with {\it
strictly red} HBs, GCs with {\it mostly red} HBs are labelled as open
squares, and GCs with intermediate HB morphologies are represented by
small filled squares.  The two open triangles stand for NGC~6388 and
NGC~6441, the filled triangle represents NGC~5904, and the $\times$
stands for 47 Tuc. The star represents M67. The effect of HB morphology
on $H\delta$ can be seen as the $\sim$1 ${\rm\AA}$ vertical spread for
--1.5 $<$ [Fe/H] $<$ --0.9. Here and in all subsequent figures, typical
error bars are displayed in the upper right corner.
} \label{fig1} 
\end{figure}

\begin{figure*}
\plottwo{f2a.eps}{f2b.eps}
\caption{Comparison between data and models for single SPs in Balmer
vs. Fe4383 plots. Symbols are the same as in Figure \ref{fig1}. Dotted
lines connect same-age models, while isometallicity lines are solid. Model
lines are labelled with age and [Fe/H] values in panel (a). As a guide to
the eye, the nearly vertical line for [Fe/H]=0 (solid), and the nearly
horizontal line for 3.5 Gyr (dashed) are both thicker. In both panels,
GCs with blue HBs tend to appear younger than those with red HBs, even
though they are all old. GCs with blue HBs also look younger according to
$H\delta_F$ (panel b) than according to $H\beta$. The arrows indicate how
the indices for NGC~5904 change when BHB stars are artificially removed.
} \label{fig2} \end{figure*}


In Figure \ref{fig2} we contrast the data with our model predictions for
Fe4383 and Balmer lines in single SPs. The Balmer indices were chosen to
provide a large baseline and because model predictions were available for
them at the time of this analysis. The models used are those by Schiavon
(2004, in preparation, but see Schiavon et al.\ 2004). 
For our present purposes, derivation of absolute ages
and [Fe/H]s is not essential, and we focus our discussion on relative
spectroscopic ages.

The main issue we address in this {\it Letter} is evident in Figure
\ref{fig2}: GCs with mostly blue HBs (asterisks) have stronger Balmer
lines and thus appear younger than GCs with red HBs (open circles) even
though they are all equally old (e.g. Rosenberg et al.\ 2002). This
effect has been pointed out by many (e.g. Freitas Pacheco \& Barbuy
1995; Lee et al.\ 2000; Beasley et al. 2002; Maraston et al.\ 2003)
and results from the fact that our SPS models do not account for BHB
stars, which have strong Balmer lines and mimick a younger age. This
is a source of confusion for cluster age estimation based on the EWs of
Balmer lines.  Furthermore, while the ages of GCs with red HBs do not vary
according to different Balmer lines, the same is not true for GCs with
blue HBs: their mean ages are $\sim$3.5 and $\sim$6 Gyr as given by \hdf\
and \hb, respectively. The same effect is also observed in the data for
NGC~6388 and NGC~6441 (open triangles), whose inferred ages are $\sim$8
or $\sim$11 Gyr according to \hdf\ or \hb, respectively. In contrast,
for M67, which is a cluster of truly intermediate age ($\sim$3.5 Gyr;
Schiavon et al.\ 2004) and whose spectrum is free from the influence
of BHB stars, the inferred age is independent of the age indicator. We
believe that these effects are caused by the increasing contribution
of BHB stars to the integrated GC light at bluer wavelengths, such that
\hdf\ is more affected than \hb. As a result, the GC spectroscopic ages
look younger according to \hdf\ than \hb.

%We also note that, because the spectrum of M67 was obtained
%from coadding individual spectra of stellar members, it is also free
%from the effects of blue stragglers.

To verify this hypothesis we computed integrated indices for NGC~5904
directly from its CMD in two different ways: first by taking all the stars
into account, and then removing all the BHB stars in order to assess
their effect on absorption line EWs. This is the most straightforward
and model independent way of assessing the contribution of BHB stars to
the integrated GC light. We chose NGC~5904 because its HB is mostly blue
and its CMD is indicative of an old age. The colors and magnitudes of
NGC~5904 stars were transformed into $T_{\rm eff}$ and $\log g$ using the
calibrations described in Schiavon et al.\ (2002a) and assuming [Fe/H]
= --1.27 (Kraft \& Ivans 2003). The resulting stellar parameters were
used as input to our fitting functions in order to predict line indices
for each star in the cluster CMD. The latter are integrated throughout
the CMD, with weights based on the photometry and stellar counts. The BV
photometry, its completeness as a function of V magnitude, the distance
modulus [$(m-M)_V = 14.46$], and the reddening towards NGC~5904 [$E(B-V)
= 0.03$] were taken from Piotto et al.\ (2002). The result is displayed
by arrows in Figure \ref{fig2} that indicate how index values change
when BHB stars are artificially removed from the CMD. The position of
the cluster in both plots is shifted towards an old age ($\sim$11--14
Gyr), which corroborates our working hypothesis. Furthermore the above
correction does not alter the agreement of our model predictions with
[Fe/H] determinations from spectroscopic analysis of individual stars
from NGC~5904 ([Fe/H] $\sim$ --1.3; Kraft \& Ivans 2003).

%compared with the indices
%measured in the integrated spectrum of NGC 5904 in Table 1. When all stars
%are included in the computation, the agreement with the measured values
%is in general good, and most of the residual differences can be ascribed
%to a combination of the uncertainties in the cluster's distance modulus,
%reddening, and metallicity, the cluster's nonsolar abundance ratios,
%uncertainties in the fitting functions for hot, metal-poor stars,
%and to the fact that the main sequence for stars cooler than $\sim$
%5200K is missing in the cluster's observed CMD. Again, we stress that
%for our present purposes, perfect agreement with the absolute observed
%values is not so relevant. Most importantly, 
%we call the attention to
%the great impact that removal of BHBS has on the Balmer lines. This can
%be seen in Figure 1, where we indicate by arrows how the indices change
%when BHBS are removed. 

%%SC  This whole paragraph below was not clear and convincing to me...
%% It is still not clear to me that we don't err if we use solar-scaled 
%% isochrones for old GCs... I'll have to return to this discussion. 
%% Before proceeding with our discussion, we should mention the caveat that

%***[NOTE: Here's my opinion about this paragraph.  I think we should pretty 
%much emilinate it.  I think that this kind of discussion will be essential
%for a later paper in the series wehere we try to make very detailed
%comparison with the models.  However, for this paper we want to minimize the
%role of the  models.  We want to stay focussed as much as possible on the
%very nice discrimination that H-deltaF/H-beta provides for detecting BHB vs
%age, and this is mostly independent of zero-point issues of the models.  Which
%is what you've ended up with in the last sentence anyway.  My impression is
%that you're concerned about the referee and/or other readers sniping at the
%results by seeing potential problems in the models, and you're producing a
%pre-emptive strike to defend them.  But I think that fully explaining this
%situation about scaled-solar vs alpha-enhanced, and how that relates to what
%actual abundancesd people are measuring spectroscopically, and thus how one
%needs to account for that, requires more explanation than you give in this 
%paragraph, so it just ends up in causing confusion and making the models sound
%more important than they are.  I think that in a future larger paper, you 
%should give the full explanation that you gave to Lauren, so that people can
%see how one must "adjust" whatever people are measuring observationally into
%a scaled abundance in the models.  But I don't think this is necessary here.
%Thus I agree with Stephane that this whole paragraph can be basically
%eliminated.  However, I tend to take an oberlky sloppy approach to modelling,
%am overly free of zero-point worries, and thus my opinion requires an 
%accompanying pound of salt!]***
%Our models are based on solar-scaled theoretical isochrones and are,
%in principle, not adequate for comparison with GC data, whose abundance
%patterns are known to be enhanced in $\alpha$ elements (e.g. Carney
%1996). However, we argue that this does not impact our conclusions. 
%%% The metallicity scale of our models is expressed in terms of [Fe/H], because
%%% the fitting functions describing line index behaviour in terms of stellar
%%% parameters are defined in terms of this quantity. The same applies to
%%% the metallicity scale of the GCs under study, which are based on iron
%%% abundance analysis of stellar members (Kraft \& Ivans 2003). 
%First, the transformation from solar-scaled to $\alpha$-enhanced models 
%must be performed at fixed [Fe/H]. This affects model predictions in two 
%ways, via the separate effects of the abundance mixture on the isochrones
%and stellar spectra. Going from solar-scaled to $\alpha$-enhanced
%theoretical isochrones at fixed [Fe/H] brings relatively minor changes,
%mostly related to the temperature and brightness of turn-off stars, which
%become cooler and fainter, thus lowering the ages inferred from Balmer
%lines by $\sim$ 1 Gyr (Salaris, Chieffi \& Straniero 1993).  Regarding
%the spectroscopic effects, the Fe indices adopted in this analysis are
%not much affected by variations in the abundances of $\alpha$-elements
%(Trippico \& Bell 1995). This explains why our models, when applied
%to GCs with red HBs, predict ages and [Fe/H]s that are in good general
%agreement with values from CMD analysis and abundance analyses
%of stellar members. Note, for instance the case of 47 Tuc (large
%cross), which has [Fe/H] $\sim$ --0.8 (Kraft \& Ivans 2003) and a model
%age of $\sim$ 11--12.5 Gyr (depending on the theoretical isochrones adopted,
%see e.g. Schiavon et al. 2002b). In panels (a) and (b) of
%Figure \ref{fig2}, 47~Tuc  %% fall on the right [Fe/H] locus, while being 
%is consistent with being $\sim$ 14 Gyr. This slightly over-estimated 
%figure is due to a combination of the adoption of solar-scaled
%models and the mismatch between model and observed luminosity functions,
%as discussed at length in Schiavon et al. (2002b). Most importantly,
%we stress that the main focus of this {\it Letter} is on relative ages,
%and for that purpose, the models presently adopted are adequate.

%\begin{figure}
%\plotone{figure3.ps}
%\caption{
%} \label{fig3} 
%\end{figure}


%If this hypothesis is confirmed, a method can be devised that is able to
%distinguish, purely on the basis of their integrated spectra, truly young
%clusters from those which are older, but whose Balmer lines are boosted
%by a BHB population, purely on the basis of their integrated spectra: the
%spectroscopic age of a truly young cluster is independent of the Balmer
%line adopted, while clusters with a BHB have younger spectroscopic ages
%when bluer Balmer lines are employed. This is not a new idea. The use of
%higher order Balmer lines to flag the presence of stars much hotter than
%the turn-offs of intermediate-age SPs has been suggested before (see,
%for instance, Caldwell, Rose \& Concannon 2003, and references therein).

%We thus conclude that the systematic lowering of
%the spectroscopic ages measured from bluer Balmer lines is consistent
%with the contribution by BHBS to the integrated GC light, and therefore
%it can be used to infer their existence in a given GC on the basis of
%its integrated spectrum.

%***[NOTE: I am also somewhat ambivalent about the necessity of this paragraph.
%The effect of CN anomalies is only an issue if the BHB clusters are CN-strong 
%and the RHB clusters are CN-weak, right?  Otherwise, we're again taking an
%unnecessarily defensive position towards this potential issue, I think.]***
%***A potential*** caveat that deserves discussion is the influence
%of anomalous carbon and nitrogen abundance ratios on \hdf, which is
%known to be severely blended by CN lines. This is a concern because all
%the GCs in our sample appear to have non-solar [C/Fe] and/or [N/Fe] ratios.
%This issue has been discussed in detail by Schiavon et al.*** (2002a),***
%and the extent of the influence of C and N abundance variations on
%Balmer lines 
%%%SC For what type (metallicity range) of GCs? Are the calibration 
%%%   and science samples drawn from a similar population?
%was shown to be small, at the level of 1--2 Gyr. The case
%of 47 Tuc, whose CN-strong integrated spectrum (Rose \& Trippico 1986;
%Rose 1994; Vazdekis 1999; Schiavon et al. 2002a) %% and yet doesn't seem 
%%% to have its spectroscopic age significantly affected argues
%appears to have little effect on its spectroscopic age, argues 
%against a strong influence of CN lines on \hdf.  Thus, we tentatively
%conclude that the abnormal carbon and nitrogen abundance ratios
%characteristic of these GCs are not %% sufficient to explain the effect.
%of relevance for this study. 
%%SC The sentence is still not completely clear; when you say "these GCs",
%% you refer to 47 Tuc plus the ones studies in Schiavon etal (2002), or
%% our GCs??
%***[Note: To echo Stephane's comment, we do have to be a bit careful about
%generalizing from our 47 Tuc experience.  The BHB clusters that we're mainly
%interested in here are quite metal-poor, thus the effects of CN on Balmer
%indices could in principle be different than for 47 Tuc...]***

In summary, we suggest that the different spectroscopic ages inferred
from $H\beta$ and $H\delta_F$ are due to the stronger impact of BHB
stars on the latter. Below, we use this observation to disentangle the
degeneracy between the effects of HB morphology and age on Balmer lines
measured in the integrated spectra of GCs.

%The latter is illustrated
%in panel (b), where models and data are contrasted in a plot of CN1
%against $H\delta_F$.  It can be seen that the GCs are all CN-strong,
%with the data being significantly shifted towards higher metallicities
%and younger ages relative to the models, whose abundance pattern reflects
%that of the solar neighbourhood. 

%\begin{deluxetable}{ccccccccc}
%%\rotate
%\tablecaption{Absorption Line Indices Measured \label{tbl-ind}}
%\tablewidth{0pt}
%\tablehead{
%\colhead {}& $H\delta_F$ & CN1 & $H\gamma_F$ & Fe4383 &  $H\beta$ & 
%Mgb & Fe5270 &  Fe5335} 
%\startdata
%NGC 5904 & 2.5 & -0.068 & 1.65 & 0.9 & 2.4 & 1.1 & 1.1 & 0.9 \\
%All stars & 2.2 & -0.086 & 1.20 & 1.1 & 2.4 & 1.2 & 1.1 & 0.6 \\
%No BHBS & 1.2 & -0.059 & 0.01 & 1.4 & 1.8 & 1.3 & 1.3 & 0.7 \\
%Error & 0.15 & 0.008 & 0.08 & 0.3 & 0.15 & 0.1 & 0.1 & 0.15 \\
%\enddata
%\end{deluxetable}

\section{Detection of Blue Horizontal Branch Stars in Globular Clusters
} \label{detect}

We have identified a distinctive signature of BHB stars in the integrated
spectra of Galactic GCs. The obvious implication is that we can devise
an indicator that identifies their presence in a given GC, solely on the
basis of absorption line EWs. Such an indicator would be most useful for
work on distant GCs, for which CMDs cannot be obtained, by providing the
means to distinguish truly intermediate-age GCs from old GCs with BHB stars.
The best indicator we found is the ratio between \hdf\ and \hb.  Figure
\ref{fig4} displays the main result of this {\it Letter}, where we plot
the data in the $H\delta_F/H\beta$ vs. Fe4383 plane, overlaid with our
models. It can be seen that GCs
with mostly blue HBs are displaced far above the locus occupied by the models,
indicating ages substantially younger than $\sim$2 Gyr, in stark constrast
with the ages inferred in Figure \ref{fig2}. In fact $H\delta_F/H\beta$
is more sensitive to HB morphology than age, given that the separation
in $H\delta_F/H\beta$ between GCs with red and blue HBs is about
twice that between models for 2 and 14 Gyr single SPs.


\begin{figure}
\plotone{f3.eps}
\caption{The data are confronted with the models in an Fe vs. Balmer
line ratio plane. Symbols, lines and arrow have the same meaning as
in the previous figures. GCs with blue HBs are greatly displaced from
the model locus, indicating ages which are much younger than those
found in Figure \ref{fig2}. In constrast the ages and [Fe/H]s for M67
and GCs with strictly red HBs are consistent with those obtained from
Figure \ref{fig2}. The ratio $H\delta_F/H\beta$ is more sensitive to HB
morphology than to age, and therefore helps in breaking the degeneracy
between the two parameters.
} \label{fig4}
\end{figure}

GCs with strictly red HBs, such as 47 Tuc, are consistent with the ages
and metallicities inferred from Figure \ref{fig2}. As expected, GCs
with mostly red or intermediate HBs lie somewhere in between the loci
occupied by GCs with extreme HB morphologies.  Note that the data for
M67, which has a truly intermediate age, are (as in Figure \ref{fig2})
consistent with its known age and [Fe/H]. We conclude that the ratio
$H\delta_F/H\beta$ discriminates between younger clusters and those
whose Balmer lines are strengthened by the presence of BHB stars.

The high S/N of our spectra make it relatively easy to detect the
influence of the BHB population on the $H\delta_F/H\beta$ index. For
instance, for NGC~5904 which has $H\delta_F/H\beta = 1.04 \pm 0.05$
(1 $\sigma$), the value expected to be consistent with the age estimated
from \hb\ from Figure \ref{fig2} ($\sim$ 6 Gyr) would be $H\delta_F/H\beta
\sim 0.84$, or $\sim$4 $\sigma$ lower than the measured value. A useful
quantity for observers is the minimum S/N needed to distinguish between
a GC with BHB stars and a young GC. This was estimated by computing index
errors as a function S/N, assuming pure Poisson noise. In the case above,
one needs $H\delta_F/H\beta$ measured to better than 0.2 which, for a
2-$\sigma$ detection requires a S/N of roughly 30--40 per ${\rm\AA}$ at
$\simeq$ 4100 ${\rm\AA}$, a value that can be achieved with reasonable
integrations of extragalactic GCs with 8-10m class telescopes (e.g. Cohen
et al.\ 2003; Larsen et al.\ 2003; Hempel et al.\ 2003).

\section{Final Remarks} \label{final}

In this {\it Letter} we propose a method to detect the signature of
BHB stars in the integrated spectra of GCs. If \hdf, \hb, and an Fe line
are measured accurately, one can infer the presence of BHB stars in a
given GC if the age indicated by $H\delta_F/H\beta$ is substantially
younger than that indicated by $H\beta$. For intermediate-age clusters,
the ages indicated by the different indices are the same
within the uncertainties. For this method to work, {\it 
SP synthesis models employed must predict correct
ages and metal abundances for systems with truly intermediate/young
ages}. We showed that our models predict the correct age and [Fe/H] for
M67. While the case for applying this method to disentangle age from HB
morphology seems strong for single-age systems, {\it it is not clear that
it should succeed when applied to galaxies},
where multiple turnoffs due to a complex star formation history cannot
in principle be distinguished from the combination of an old turnoff
with other kinds of warm stars like BHB stars.

%***[NOTE: I suggest we remove this last paragraph. In a Letters paper, where
%space is at a premium, I see no reason to burn up space in advertizing future 
%plans.]***
%Globular clusters are the ultimate benchmark against which SPS models must
%be confronted. Therefore, a homogeneous and accurate set of integrated
%spectra of these systems is a precious asset, whose exploration will
%certainly yield a deeper understanding of the complexities inherent to
%an accurate modelling of stellar population spectra. This {\it Letter}
%is a first attempt in this direction. An in depth analysis of our data,
%combined with publically available deep color-magnitude diagrams will
%be the topic of forthcoming papers. 


\acknowledgments

We thank Francesca de Angeli and Manuela Zoccali for providing us with
data on the completeness of the photometry for NGC~5904. Ruth Peterson,
Duncan Forbes, Robert Proctor, and an anonymous referee are thanked
for useful suggestions. R.P.S. acknowledges financial support from
NSF grant AST 00-71198 to the University of California, Santa Cruz,
and from HST Treasury Program Grant GO-09455.05-A to the University
of Virginia.  S.C. and L.A.M. acknowledge support from the National
Science and Engineering Research Council of Canada, and J.A.R. from NSF
grant AST-99000720 to the University of North Carolina.


 
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%% Generally speaking, only the figure captions, and not the figures
%% themselves, are included in electronic manuscript submissions.
%% Use \figcaption to format your figure captions. They should begin on a
%% new page.

\clearpage

%% No more than seven \figcaption commands are allowed per page,
%% so if you have more than seven captions, insert a \clearpage
%% after every seventh one.

%% There must be a \figcaption command for each legend. Key the text of the
%% legend and the optional \label in curly braces. If you wish, you may
%% include the name of the corresponding figure file in square brackets.
%% The label is for identification purposes only. It will not insert the
%% figures themselves into the document.
%% If you want to include your art in the paper, use \plotone.
%% Refer to the on-line documentation for details.


%\figcaption[sgi9259.eps]{This is the first figure and it uses sgi9259.eps as
%its EPS figure file. \label{fig1}}


%% Tables should be submitted one per page, so put a \clearpage before
%% each one.

%% Two options are available to the author for producing tables:  the
%% deluxetable environment provided by the AASTeX package or the LaTeX
%% table environment.  Use of deluxetable is preferred.
%%

%% Three table samples follow, two marked up in the deluxetable environment,
%% one marked up as a LaTeX table.


%% In this first example, note that the \tabletypesize{}
%% command has been used to reduce the font size of the table.
%% Note also that the \label command needs to be placed 
%% inside the \tablecaption.

%% You can append references to a table using the \tablerefs command.

%% The following command ends your manuscript. LaTeX will ignore any text
%% that appears after it.

\end{document}

%%
%% End of file `sample.tex'.

astro-ph0404592.extracted_bib
\begin{thebibliography}{}
\bibitem[Beasley et al. (2002)]{bhs02} Beasley, M.A., Hoyle, F. \& Sharples,
R.M. 2002, \mnras, 336, 168
\bibitem[Borkova \& Marsakov (2000)]{bm} Borkova, T.V. \& Marsakov, V.A. 2000,
Astronomy Reports, 44, 665
\bibitem[Brown (2003)]{b03} Brown, T.M. 2003, (astro-ph/0308509)
\bibitem[Brown et al. (2000)]{betal00} Brown, T.M., Bowers, C.W., Kimble, R.A.,
Sweigart, A.V. \& Ferguson, H.C. 2000, \apj, 532, 308
\bibitem[Brown et al. (1998)]{betal00} Brown, T.M., Ferguson, H.C., Stanford,
S.A. \& Deharveng, J.-M. 1998, \apj, 504, 113
%\bibitem[Caldwell et al. (2003)]{crc03} Caldwell, N., Rose, J.A. \& Concannon,
%K.D. 2003, \aj, 125, 2891
\bibitem[Carney (1996)]{car} Carney, B.W. 1996, \pasp, 108, 900
\bibitem[Carretta \& Gratton (1997)]{cg97} Carretta, E. \& Gratton, R.G.
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\bibitem[Catelan (2000)]{c00} Catelan, M. 2000, \apj, 531, 826
\bibitem[Cohen et al. (2003)]{cetal03} Cohen, J.G., Blakeslee, J.P., \&
C\^ot\'e, P. 2003, \apj, 591, 850
\bibitem[Freitas Pacheco \& Barbuy (1995)]{fpb} Freitas Pacheco, J.A. \&
Barbuy, B. 1995, \aap, 302, 718
\bibitem[Girardi et al. (2000)]{getal00} Girardi, L., Bressan, A., Bertelli,
G. \& Chiosi, C. 2000, \aaps, 141, 371
\bibitem[Harris (1996)]{h96} Harris, W.E. 1996, \aj, 112, 1487
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\bibitem[Lee et al. (2000)]{letal00} Lee, H.-c., Yoon, S. \& Lee, Y. 2000, \aj,
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\bibitem[Rosenberg et al. (2002)]{retal02} Rosenberg, A., Aparicio, A., 
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%\bibitem[Salasnich et al. (2000)]{slas} Salasnich, B., Girardi, L., Weiss, A.
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\bibitem[Schiavon et al. (2002a)]{setal02a} Schiavon, R.P., Faber, S.M.,
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\bibitem[Schiavon et al. (2004)]{setal04} Schiavon, R.P., Caldwell, N. \&
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%\bibitem[Trager et al. (2000)]{tetal00} Trager, S.C., Faber, S.M., Worthey,
%G. \& Gonz\'alez, J.J. 2000, \aj, 120, 188
%\bibitem[Trippico \& Bell (1995)]{tb95} Trippico, M.J. \& Bell, R.A. 1995,
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%\bibitem[Vazdekis (1999)]{v99} Vazdekis, A. 1999, \apj, 513, 224
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