\documentstyle[psfig,paasms4,pptwocol,pub]{preprint}
\begin{document}
\title{WFPC2 IMAGING OF THE GLOBULAR CLUSTER SYSTEM\\
~\\
OF THE S0 GALAXY 
NGC~3115\thanks{Based on observations with the NASA/ESA {\it Hubble Space 
Telescope\/}, obtained at the Space Telescope Science Institute, which is 
operated by the Association of Universities for Research in Astronomy, Inc., 
under NASA contract NAS5-26555.}   }

\author{Arunav Kundu\thanks{Astronomy Department, University of Maryland,  
College Park, MD 20742.} \& Bradley C.\ Whitmore\\
\\
Space Telescope Science Institute\\
\\
3700 San Martin Drive\\
\\
Baltimore, MD 21218\\
\\
Electronic Mail: akundu@stsci.edu; whitmore@stsci.edu}

\tobe{December 1998}{The Astronomical Journal}
\recacc{4 June 1998}{19 August 1998}
\url{1287}
\maketitle

\abstract{
Using the WFPC2 on-board the HST we have identified 144  globular cluster
candidates in the inner region of the S0~galaxy NGC~3115. The color distribution
of the globular clusters is bimodal, with a blue peak at V--I~= 0.96~mag and a 
red peak at V--I~= 1.17~mag. There are roughly equal number of blue and red clusters. 
The mean \mbox{[Fe/H]} of the blue population is --1.36~dex while the red 
population has a mean \mbox{[Fe/H]} of --0.37~dex.  Following  Ashman \& Zepf 
(1992), we interpret the bimodality as evidence of multiple epochs of cluster 
formation in the metal enhancement history of the galaxy. Assuming that the blue 
peak corresponds to a primordial, 15~Gyr old,  metal poor  population of globular 
clusters, we derive an age  of $11\pm3$~Gyr for the younger, red clusters. We 
also find that the age and metallicities of the globular clusters are consistent 
with the red giant stars in the  field, suggesting that the two systems formed  
coevally. The spatial distribution of both the red and the blue clusters are 
flattened with major and minor axes roughly colinear with the galaxy light profile. 
While the red clusters follow the thick disk component of the galaxy, the blue 
cluster distribution is less flattened and mimics the bulge/halo light 
distribution. We find a local specific frequency of $1.3\pm0.1$ in the  central 
region of the galaxy,  which is lower than the global specific of $2.3\pm0.5$ 
reported in the literature, implying that the specific frequency increases with 
galactocentric distance. We have also observed the globular cluster luminosity 
function 2~magnitudes deeper than the turnover luminosity and derived a distance 
modulus of m--M~$=29.8\pm0.3$~mag to NGC~3115. The average half light radii of 
the clusters is $\approx2$~pc and the red clusters are $\approx20$\% smaller 
than the blue ones. }
\base

\section{INTRODUCTION}
The advent of charge coupled devices and the Hubble Space Telescope has led to 
a recent explosion  in the number of studies of Globular Cluster Systems (GCS)
of external galaxies. The concept of what constitutes a typical GCS has 
undergone a revision due to the discovery of bi-modal color distribution of 
clusters  in many elliptical galaxies.  Most studies of GCSs have concentrated 
on ellipticals due to their rich cluster systems and the ease of identifying 
globular clusters in the field. In this paper, we study the GCS of the nearby  
S0 field galaxy NGC~3115, and attempt to trace it's formation history in the 
context of the various evolutionary scenarios that have been suggested. NGC~3115 
is an isolated, edge-on galaxy,  with a 10$^9$~M$_\odot$ black hole believed to 
be at it's center (Kormendy \& Richstone 1992; Kormendy et~al.\ 1996). The
only previous detailed analysis of the cluster system of this galaxy is by 
Hanes \& Harris (1986). The current study considerable improves our knowledge 
of the GCS of NGC~3115 since it goes $\sim2$~mag deeper than the turnover 
magnitude of the Globular Cluster Luminosity Function (GCLF) and includes 
observations in two colors.

One of the most significant  discoveries in globular cluster research in recent 
times is that the broad-band color distribution of $\approx50$\% of the early 
type galaxies studied is bimodal. This effect, which is usually interpreted as 
a bimodality in the metal abundance of the clusters, suggests that globular 
clusters have formed during two distinct epochs in the metal enhancement histories 
of these galaxies. Schweizer (1987) and Ashman \& Zepf (1992) proposed that GCs 
may be formed during the interaction or merger of galaxies. While Zepf \& 
Ashman (1993) suggest that the bimodality in the color distribution is a natural 
consequence of the merger scenario, Forbes, Brodie \& Grillmair (1997) favor 
multiple phases of cluster formation during the collapse of a gas cloud into a 
galaxy. The search for bimodality in the the color distribution of globular 
clusters has  largely been limited to elliptical galaxies, with  NGC~1380 being
the only S0~galaxy confirmed to have a bimodal color distribution  
(Kissler-Patig et~al.\ 1997). 

Recently Elson (1997)  reported a bimodal metallicity distribution of red giant 
branch stars (RGB) in the halo of NGC~3115, suggesting that it underwent two 
distinct periods of star formation. The main goals of this paper are to study 
the properties of the globular clusters, compare the cluster population with 
the RGB stars, and attempt to differentiate between the various models of 
globular cluster formation. 

\section{OBSERVATIONS AND DATA REDUCTION}
This study is based on archival HST Wide Field and Planetary Camera~2 (WFPC2; 
hereafter WF~= Wide Field Camera and PC~= Planetary  Camera) images of the 
inner region of NGC~3115, obtained as part of HST program 5512 by Faber et~al.\ 
(cores of early type galaxies). We also re-analyze the red giant branch stars 
in the field studied by Elson (1997) using our detection methods and 
photometric offsets to check for consistency.  The images for this field 
were obtained as part of the Medium Deep Survey (Griffiths et~al.\ 1994).	

The relative positions of the two fields with respect to NGC~3115 are shown in 
Fig.~1. Henceforth, we shall refer to the field centered on the nucleus of 
NGC~3115 as the globular cluster field, and the field centered on the halo of 
the galaxy as the RGB field. 
\begin{figure*}%1
\centerline{\epsfxsize=\textwidth\epsfbox{pp1287f1.ps}}
\caption{The relative positions of the two fields being studied: The Globular 
Cluster Field is centered on the nucleus of NGC~3115. The Red Giant Branch Field
is located 8.5~arcmin (22.5~Kpc) East of the galaxy. Arrow points North, bar 
points East, and each is 20~arcmin long \label{fig1}}
\end{figure*}

\subsection{The Globular Cluster Field}

The nuclear region of NGC~3115 was imaged with the WFPC2 on Nov.~3, 1994. 
3~sets of 350~secs images were  obtained with the F555W and F814W filters. We
used the STSDAS utility GCOMBINE and the IRAF task COSMICRAYS  to reject
cosmic rays and hot pixels respectively. We identified  cluster candidates 
using the technique described in detail in Kundu et~al.\ (1998). Briefly, 
DAOPHOT is used to identify sources using a very low cutoff (1.5$\sigma$). Then 
photometry is performed on each of the candidates and the signal to noise is 
estimated by taking the ratio of the counts within the aperture to the rms noise 
in the sky background. For this set of images we used a S/N cutoff of~3 to 
detect candidate clusters. We also used the ratio of the flux within an aperture 
of 0.5 and 3~pixels to weed out chip defects and compact background galaxies. 
Using a concentration criteria $2<\frac{Counts_{3pix}}{Counts_{0.5pix}}<13$ for 
the PC and $2<\frac{Counts_{3pix}}{Counts_{0.5pix}}<10$ for the WF, we 
identified 157 cluster candidates that satisfied the selection criteria in both 
the F555W and the F814W image. After studying the radial profiles of the 
candidates we found that all the sources with $Counts_{0.5pix}>1000$~DN were 
foreground stars and we rejected these. Our final candidate list 
contains 144 objects. The cluster candidates are marked with circles on the 
F555W image of the globular cluster field in Fig.~2.
\begin{figure*}%2
\centerline{\epsfxsize=4in\epsfbox{pp1287f2.ps}}
\caption{Mosaic of $3\times350$~sec V~(F555W) band WFPC2 images of the inner 
region of NGC~3115.  144 globular cluster candidates are identified. The 
distribution of clusters appears to be flattened and roughly follows the light 
distribution of the galaxy.
\label{fig2}}
\end{figure*}

After correcting for the geometrical distortion (Holtzman et~al.\ 1995a),
aperture photometry was performed using a 3~pixel radius aperture for the PC, 
and a 2~pixel aperture for the WF. The median pixel value between 5 and 8~pixels 
was used as the sky background. Since the profiles of the cluster candidates are 
slightly broader than a stellar PSF, we  re-derived the aperture correction by 
calculating the ratio of counts within a given aperture with respect to an 
aperture of 0.5$''$ for an `average' cluster in our field. In order to 
extrapolate the aperture correction to infinite radius we added 0.1~mag to the 
correction (Holtzman et~al.\ 1995b).   The aperture corrections used for the 
F555W filter were --0.43~mag (3~pixel radius aperture) and --0.42~mag (2~pixel 
radius aperture) for the PC and WF respectively. The corresponding numbers
for the F814W filter were  --0.49~mag  and --0.45~mag. We adopted photometric
zeropoints of 22.573~mag and 21.709~mag (following Whitmore et~al.\ 1997) and 
small color correction terms from Holtzman et~al.\ (1995b) to convert the F555W 
and F814W magnitudes to Johnson~V and  Cousins~I respectively. The foreground 
reddening in the direction of NGC~3115 is estimated to be A$_B=0.10$~mag (Burstein 
\& Heiles 1984). Using the reddening curve from Mathis (1990), we applied small 
corrections of A$_V=0.078$ and  A$_I=0.047$. We did not correct for the Charge 
Transfer  Efficiency (CTE)  because the CTE problem is expected to be minimal 
in the presence of the strong galaxy background (Holtzman et~al.\ 1995b; 
Whitmore \& Heyer 1997). 

\subsection{The Red Giant Branch Field}

The RGB field is located 8.5~arcmin east of the center of NGC~3115 
($\approx22.5$~Kpc for an assumed distance of 9.1~Mpc),  and 5~arcmin from the 
major axis of the disk of the galaxy. It was observed with the  F606W and F814W 
filters at~4  dithered positions,  between 27~Dec 94 and 30~Dec 94. The total 
observing time is 16500s in the F606W filter and 23100s  in the F814W filter. 

Initially, cosmic rays and hot pixels were removed at each of the dithered 
positions. Subsequently  the DRIZZLE algorithm developed by Fruchter \& Hook 
(1997) was used to re-construct a sharper image from the undersampled images at 
the 4~dithered positions. Following Elson (1997), we searched for RGB stars only 
in the WF chips of  the deeper F814W reconstructed image. We used the object finding 
routine described above, with the constraints,
$1.5<\frac{Counts_{3pix}}{Counts_{0.5pix}}<6$ and S/N $>2.5$ to identify 3956 
unresolved point sources. Aperture photometry was then performed within a 2~pixel 
radius aperture and sky radius between 5 and 8~pixels. Since the average flux 
from each of the stars is less than 200~DN and the background counts are low, 
the CTE effect is expected to be significant (Holtzman 1995b; Whitmore \& Heyer 
1997), though the long exposure times lead to background counts of $\approx20$ 
DN which somewhat reduces the problem. Using the CTE measurements of Whitmore 
\& Heyer (1997) and the time dependence of the CTE effect measured by Whitmore 
(1998),  the estimated Y-CTE in the F606W chip is 3.5\% and the X-CTE is 1\%. 
The corresponding values in the F814W fields are 4.5\% \& 1.5\%. We applied the 
CTE correction as a linear ramp, independent in the X and Y~directions, with 
zero correction at X~=~0,Y~=~0. We did not apply any magnitude dependence to 
the CTE ramp since this effect was  very minor ($\approx1$\%) over the relevant 
range of magnitudes at the time of these observations. 
  
Aperture corrections of --0.30~mag for F606W and --0.36~mag for F814W, derived 
from our tests, and corresponding zeropoints of 22.887~mag and 21.709~mag were 
used to convert the magnitudes of the stars to V and Cousins~I. After applying 
the small reddening corrections quoted in the last section, color corrections 
from Holztman et~al.\ (1995b) were applied to convert V$_{(F606W)}$ to Johnson~V.
The behavior of  V$_{Johnson}$--V$_{(F606W)}$ in the vicinity of V--I~= 2.0~mag 
is significant and it's possible effect shall be discussed in \S3.2.

\section{RESULTS AND DISCUSSION}

\subsection{The Color Distribution}

The color distribution of the cluster candidates with photometric errors less 
than 0.2~mag is plotted in Fig.~3. The colors of the clusters lie in a narrow 
range between V--I~= 0.8~mag and V--I~= 1.4~mag and the  distribution is clearly 
bimodal with  blue and  red peaks  near V--I~= 0.94~mag and V--I~= 1.15~mag 
respectively. The mean of the color distribution of the entire population is 
V--I~$=1.11\pm0.02$~mag while that of the 11 brightest objects is  
V--I~$=1.06\pm0.03$~mag. The mean color of the candidate  object set in the 
range  $0.8<$~V--I$<1.4$, which is largely free of contamination, is 
V--I~$=1.06\pm0.03$~mag. 
\begin{figure*}%3
\centerline{\psfig{figure=pp1287f3.eps,width=4in,angle=-90}}
\caption{ Histogram of V--I colors of the cluster candidates.  The distribution 
is clearly bimodal, suggesting that the globular clusters formed during two 
distinct epochs in the metal enhancement history of the galaxy.
\label{fig3}}
\end{figure*}

We also used the KMM mixture modelling algorithm (Ashman, Bird \& Zepf 1994) to 
test from bimodality and to find objective partitions between the sub-populations. 
KMM tests on the 133 cluster candidates with colors $0.8<$~V--I~$<1.4$, 
indicate that the distribution of cluster colors can be divided into two 
statistically significant sub-populations. For a homoscedastic partition in 
which both populations are forced to have equal variances, we obtain 64 blue 
clusters with a mean color of V--I~$=0.96$~mag and 69 red clusters with a mean 
V--I color of 1.17~mag. The common variance is 0.004~mag and the threshold color 
dividing the two populations is V--I~$=1.06$~mag. In the case of a 
heteroscedastic partition in which both populations are allowed different
variances, the algorithm estimates the mean color of 58 blue clusters to be
V--I~$=0.94$~mag and that of 75 red clusters V--I~$=1.16$~mag. The variances of 
the red and blue clusters are 0.003~mag and 0.006~mag respectively. Throughout 
the rest of the paper we shall adopt V--I~$=1.06$~mag as our best estimate of 
the threshold color. Clusters with colors between $1.06<$~V--I~$<1.4$ will be 
referred to as red clusters while those in the range $0.8<$~V--I~$<1.06$ will 
be referred to as the blue clusters. We adopt  V--I~$=0.96$~mag for the mean 
color of the blue clusters and V--I~$=1.17$~mag for the mean of the red ones.  

Stellar population modelling shows that the color difference between cluster 
populations is dominated by age effects if the younger population is less than
1~Gyr old, and by metallicity effects otherwise (Worthey 1994; Bruzual \& Charlot
1993). Since there is no evidence of recent star formation in the halo of NGC~3115, 
the color difference   is primarily due to metallicity variations. The peak of the 
blue clusters at V--I~$=0.96$~mag coincides exactly with the peak color of the 
old, metal poor, clusters found in many other galaxies e.g.\ Milky Way, M87, 
NGC~4406 (Whitmore et~al.\ 1997). Consequently, the red population must be more 
metal rich and have formed later in the metal enrichment history of the galaxy. 
In \S3.3 we argue that the red clusters are younger than the blue clusters and 
have a mean age of $11\pm3$~Gyr as compared to a 15~Gyr old blue population.

The  conversion  between the broad band colors and metallicities of globular 
clusters is known to be roughly linear from population synthesis models. The 
most commonly used, empirical relationship,  derived by Couture, Harris \& 
Allwright (1990), using the Milky Way globular clusters, is Fe/H~$=-6.096+5.051$(V--I) 
(Note: Actually, they derive V--I~$=0.198$[Fe/H]~$+1.207$ using [Fe/H] as the 
independent variable but we have inverted it to make the comparison below). 
However, since most of the Milky Way clusters are old and blue, this expression 
may not be accurate when extrapolated to higher metallicity. Kissler-Patig et~al.\ 
(1998) avoided this problem  by combining their spectroscopic observations of the 
metal-rich clusters in NGC~1399 with data on the Milky Way clusters to derive  
[Fe/H]~$=-4.50(\pm0.30)+3.27(\pm0.32)$(V--I). They point out that the slope of 
their expression is smaller than the value derived by Couture et~al.\ (1990), 
and that mean metallicities have probably been overestimated in the past for 
this reason. We re-derived the metallicity-color relationship using the 
McMaster Catalog of Milky Way clusters (Harris 1996) and the reddening curve 
from Mathis (1990) E$_{(V-I)}=1.22$E$_{(B-V)}$. Restricting the data points to 
clusters with E$_{(B-V)}<0.4$ we derive the following expressions:
[Fe/H]~$=-7.31(\pm0.74)+6.25(\pm0.60)$~(V--I) (using Fe/H as the independent 
variable and inverting) and [Fe/H]~$=-5.02(\pm0.35)+3.78(\pm0.38)$~(V--I) 
(using V--I as the independent variable).  The former expression is closer to 
the one derived by Couture et~al.\ (1990), while the  latter is closer to the
relation derived by Kissler Patig et~al.\ (1998). Hence, most of the difference  
between the two expressions is probably due to the choice of independent variables. 
The underlying reason for the variation in the slope may be a small non-linearity 
in the color-metallicity relationship that is predicted by population synthesis 
models, different levels of uncertainties in the determination of the broadband 
colors and metallicities of globular clusters, or a combination of the two.  
Since both Fe/H and V--I have internal errors due to the age and foreground 
reddening effects, there is no obvious reason to choose one over the other as 
the independent variable. We use the bisector between the two expressions 
derived by us to convert from broadband V--I colors to metallicity i.e.\ 
\begin{equation}  
[Fe/H] = -5.89 + 4.72  (V-I)
\end{equation}

Using this expression the peaks of the globular cluster color distribution at 
V--I~$=0.96$~mag and 1.17~mag, correspond to metallicities of Fe/H~$=-1.36$~dex 
and  --0.37~dex respectively. 

\subsection{The Distance to NGC~3115 using the Tip of the RGB and a Comparison 
of the Globular Cluster System with the Red Giant Stars in NGC~3115}

Our next goal is to compare the field stars in NGC 3115 with the globular 
cluster population.  However, before we can do this we need to determine the 
distance modulus in order to normalize the theoretical isochrones used in the 
comparison.

Following Elson (1997), we calculated the derivative of the luminosity function 
of the stars in the RGB field using different bin widths (0.1, 0.2, 0.3 \& 0.4~mag) 
and  bin centers. Averaging over all cases we find a peak in the slope of the 
luminosity function at m$_I=26.19\pm0.07$~mag.  Adopting the value of  M$_I\approx-4.0$ 
for the tip of the RGB from the estimate of Lee, Freedman \& Madore (1993), we 
derive a distance modulus of m--M~$=30.19\pm0.3$~mag. Using the same data, 
Elson (1997) found a peak in the derivative of the luminosity function at   
m$_I=26.14\pm0.04$~mag and a distance modulus of  m--M~$=30.21\pm0.3$~mag.

Fig.~4 is a I vs.\ V--I plot  of the stars in the RGB field. The objects to the 
right of the dotted line are mostly red-giant branch stars, while the blue 
objects to the left of the line are dominated by background galaxies. The dashed 
line is a plot of the location of the RGB stars in M15 (Da~Costa \& Armandroff 
1990), which have a metallicity of Fe/H~$=-2.17$~dex, normalized to a distance 
modulus of 30.19~mag. It  roughly corresponds to the expected isochrone of the 
most metal poor clusters in the galaxy.  The solid line to the left traces the 
red giant branch of NGC~6752 (Da~Costa \& Armandroff 1990). It represents a 
metallicity of Fe/H~$=-1.54$~dex and roughly follows the isochrone of the 
average old, blue, globular clusters in NGC~3115. The solid line to the right 
traces the theoretical isochrone of an 11~Gyr old stellar population with a 
metallicity of Fe/H~$=-0.4$~dex (Bertelli et~al.\ 1994). This represents our 
best guess of where the red giants would fall if they had the same age 
(discussed in \S3.3) and metallicity as the mean red globular cluster population 
in NGC~3115. The presence of metal rich stars in what is purportedly a 
population of halo stars is rather surprising. It is possible that the metal 
rich stars in the RGB field are actually linked to the disk and/or bulge, hence 
we shall refer to the candidates in the RGB field as the field stars instead of 
the halo population from here on.
\begin{figure*}%4
\centerline{\psfig{figure=pp1287f4.eps,width=4in,angle=-90}}
\caption{ Color magnitude diagram of the 3956 objects detected in the Red Giant 
Branch Field. The objects to the left of the dotted line are mostly background 
compact galaxies, while those to the right are Tip of the Red Giant Branch stars.
The solid curves trace the loci of the Red Giant Branch of (from left to right)
M15 (Fe/H~$=-2.17$~dex), NGC~6752 (Fe/H~$=-1.54$~dex), and a 11~Gyr model 
cluster with Fe/H~$=-0.4$~dex. 
\label{fig4}}
\end{figure*}

The isochrones predicted by the globular clusters in NGC~3115 are within the 
spread of the RGB stars.  Hence, the globular cluster system of NGC~3115 appears 
to have formed  roughly coevally with the field stars.  It should be pointed 
out that the derived metallicity of the blue clusters is not affected significantly 
by the choice of the color-metallicity relation and the clusters are found to be 
coeval with the field stars irrespective of the equation that we use. The position 
of the isochrone of the red clusters, on the other hand, is quite sensitive to 
the choice of the color-metallicity relationship, and the assumed age of the clusters. 
The calculated metallicity of the blue clusters in NGC~3115 varies by at most 
$\sim0.1$~dex when any of the above color-metallicity expressions are used, 
whereas the derived metallicities of the red clusters can change by $\sim0.5$~dex 
using the same expressions. 

One of the conclusion of Elson's (1997) study of the RGB stars in NGC~3115 
is that the distribution is bimodal.  In an attempt to verify this claim, we 
binned the color distribution of our data in similar magnitude ranges. The 
results are plotted in Fig.~5. We see possible evidence of bimodality, consistent 
with Elson's results, in some of the bins. In order to objectively test the 
bimodality in the color distribution we performed KMM tests on the color 
distribution in each of the bins. For a homoscedastic partition of objects in 
the magnitude range $26.0<$~V~$<26.3$ the algorithm returned two significant 
populations, blue objects with a mean V--I of 1.58~mag, red objects with a mean 
V--I of 2.60~mag, and a threshold dividing color of V--I~$=2.05$~mag. The objects 
in the magnitude range $26.3<$~V~$<26.5$ could also be divided into two 
sub-populations,  blue stars with a mean V--I~$=1.66$~mag, red stars with a mean 
V--I~$=2.73$~mag, and a  threshold color of V--I~$=2.22$~mag. Heteroscedastic  
partitions for these magnitude bins yield the following  results: RGB candidates 
in the brightness range $26.0<$~V~$<26.3$ can be divided into two populations 
with mean colors of V--I~$=1.36$~mag and 2.38~mag respectively while objects in 
the range $26.3<$~V~$<26.5$ yield mean colors of V--I~$=1.68$~mag and
V--I~$=2.76$~mag for the two sub-populations. However, KMM analysis of objects 
in the magnitude range $26.5<$~V~$<26.7$ did not return any reasonable 
partitions and somewhat weakens the case of a bimodal metallicity distribution 
of RGB stars. Furthermore, if we were to assume that two sub-populations derived 
in the brighter magnitude bins trace isochrones of different metallicity as
claimed by Elson (1997), the mean color of the red and blue objects in the range
$26.0<$~V~$<26.3$ should be {\it redder\/} than the corresponding peaks of 
the objects with $26.3<$~V~$<26.5$ whereas we find from the KMM tests that they
are {\it bluer\/} in each case. Thus the case for a bimodal distribution
of RGB stars is somewhat unconvincing. We also noticed that the minima between
the two peaks in all of Elson's plots and our own is at, or near, V--I~$=2.0$~mag. 
This is cause for some concern since this is the breakpoint for using different 
expressions to correct the instrumental colors to V--I colors.  Holtzman et~al.\ 
(1995b) show in their Fig.~10 that there is a systematic offset of $\approx0.1$~mag
between the true V~magnitude and the derived V$_{606}$ magnitude at V--I~$=2.0$.
We found that adding offsets of this order to the V--I colors in a small range
about V--I~$=2.0$~mag either wiped out, or shifted, the peaks of the color
distribution. We also found that we could induce false peaks that were 
$\approx20$\% higher than the surrounding bins  by adding the  photometric 
errors to simulated distributions.  Thus, even though the position of the 
minimum of the RGB color distribution is tantalizingly close to the isochrone 
corresponding to the minimum of the globular cluster color distribution (at 
V--I~$=1.06$~mag) we cannot rule out the possibility that the perceived lack of 
field star formation in both Elson (1997) and our plots are entirely instrumental 
effects.
\begin{figure*}%5
\centerline{\epsfxsize=4.5in\epsfbox{pp1287f5.eps}}
\caption{ The color distribution of the red giant branch stars in 3~magnitude 
ranges.
\label{fig5}}
\end{figure*}

We also searched for cluster candidates in the RGB field using the same 
parameters used in the globular cluster field. We found 18 candidates with
V~$<24$~mag with no obvious peak in the color or magnitude distribution.
Close inspection of the profiles suggest that 5 of them may be cluster candidates. 
The colors of these 5~candidates lie within $0.8<$~V--I~$<1.3$, suggesting that 
these are probably globular clusters while the rest are contaminating field 
objects. This suggests that we have about 13 contaminating objects per WFPC2
field, most of which lie outside the color range of clusters. That is consistent
with our observation that 11 of the 144~GC objects in the GC field have colors 
redder than V--I~$=1.4$ or bluer than V--I~$=0.8$ and are most likely foreground
stars or background galaxies. We can safely assume that the 133 candidates with 
colors between $0.8<$~V--I~$<1.4$ are globular clusters and that sample is
practically free of contamination. 
   
\subsection{The Globular Cluster Luminosity Function: Distance and Age of 
the Clusters}

The absolute magnitude of the peak of the GCLF has been found to be remarkably 
constant over a wide range of galaxies and environments (Harris 1991).  While 
the theoretical basis for this remains unclear, the uniformity of the GCLF can 
be exploited to determine distances to distant galaxies (Jacoby et~al.\ 1992; 
Whitmore 1996). 

Fig.~6 is a plot of the histogram of the luminosity distribution of the cluster 
candidates in NGC~3115 with colors between $0.8<$~V--I~$<1.4$. As explained in 
the previous section, there is little or no contamination by foreground or 
background objects.  The distribution clearly shows a peak near V~$=22.5$~mag. 
However, the faint end of the distribution is affected by incompleteness and we 
must correct for this in the GCLF.  We added $\approx9000$ simulated clusters 
to the image using ADDSTAR to determine our completeness correction. The brightness 
of the objects was varied between $20<$~V~$<26$ with a random distribution of 
colors between V--I~$=0.8$~mag and V--I~$=1.3$~mag to roughly simulate the 
actual colors of the clusters. The 50\% completeness limit in the PC chip 
was found to be at V~$= 23.3$~mag while the completeness limit in the WF was 
found to be much fainter (24.5~mag) mostly due to the much lower galaxy 
background.  The dotted curve in Fig.~6 traces the completeness corrected GCLF. 
Assuming a completeness limit at V~$=24.4$~mag we find that the corrected GCLF 
is well fit by a Gaussian. Using NGAUSSFIT,  the Gaussian curve that best fits 
the corrected GCLF has a turnover luminosity at $<m_v>=22.37\pm0.05$ and 
$\sigma=1.29\pm0.06$. We use this as out best estimate of the turnover 
luminosity. The estimate of the turnover is fairly robust and  changes by less 
than 0.1~mag even when the 11~objects with colors outside the range 
$0.8<$~V--I~$<1.4$ are added.
\begin{figure*}%6
\centerline{\psfig{figure=pp1287f6.eps,width=4in,angle=-90}}
\caption{ The Globular Cluster Luminosity Function derived from the objects 
with colors between $0.8<$~V--I~$<1.4$. The dotted line traces  the completeness 
corrected distribution and the dashed lines give the best Gaussian fit.    
\label{fig6}}
\end{figure*}

The current estimates of the absolute magnitude of the peak of the GCLF 
describing the Milky Way and M31 range from V~$=-7.6$ to V~$=-7.2$ (Harris 1991; 
Whitmore 1996). Following Whitmore et~al.\ (1995) and Whitmore (1996)
we adopt a value of M$_V^0=-7.4\pm0.25$~mag for the turnover luminosity.
This yields a distance modulus of m--M~$=29.8\pm0.3$~mag, which corresponds to
a distance of $9.1\pm1.3$~Mpc to NGC~3115. This value is marginally smaller than 
the distance modulus of $30.2\pm0.3$~mag determined from the TRGB in the last 
section, and m--M~$=30.17\pm0.13$ determined from the planetary nebula luminosity 
function (Ciardullo, Jacoby \& Tonry 1993). However, it is larger than the 
previous estimate of m--M~$=29.65\pm0.25$ using surface brightness fluctuations 
(Ciardullo et~al.\ 1993). This has  since been updated to m--M~$=29.9$ (Tonry 
1998; private communication), and is in excellent agreement with our GCLF distance.  
Though our distance estimate is on the lower side of this spread, it is 
consistent, within the errors, with other distance estimates, and reaffirms our 
faith in the usefulness of the GCLF as a secondary distance indicator. In all 
subsequent sections we will use the distance modulus derived from the GCLF 
(m--M~$=29.8$~mag or distance~$=9.1$~Mpc) to calculate distant dependent 
quantities.

We now turn to the question of the difference between the luminosity functions 
of the red and blue clusters and the derived ages of the clusters. We find that 
the blue clusters are $0.16\pm0.20$~mag brighter than the red clusters in
the V~band, and $0.06\pm0.20$~mag  dimmer than the red clusters in the I~band. 
For the metallicities of the blue and red clusters derived in \S3.1, Ashman, 
Conti \& Zepf (1995) (ACZ) predict that the blue clusters should be 0.3~mag 
brighter than the red clusters in the V~band  and 0.1 brighter than the red 
clusters in the I~band if the mass to light ratio for all clusters is the same. 
This discrepancy is very similar to the results for M~87 (Kundu et~al.\ 1998) 
and can be resolved if we relax the ACZ assumption of coeval formation of the 
two populations, allowing the red population to be younger. In this context,
Whitmore et~al.\ (1997) have traced the evolution of the color and magnitude 
difference between an old metal poor system and a younger more metal rich one, 
using Bruzual \& Charlot (1993) models. We estimate, from these  tracks, that 
the age of the red clusters is $11\pm3$~Gyr, or, $4\pm3$~Gyr younger than the 
blue population on a scale in which the blue, metal poor population is 15~Gyr 
old.

\subsection{The Specific Frequency of Globular Clusters} 

The WFPC2 covers only a small portion of the inner region of NGC~3115, hence we 
cannot calculate the total specific luminosity of the entire galaxy directly. 
We can however, derive the local specific frequency and place a lower limit on 
S$_N$. In order to do this we must first  calculate the underlying galaxy 
luminosity within our field of view. By fitting our data to published large 
aperture integrated luminosities along the minor axis of the galaxy (Strom et~al.\ 
1974; Persson et~al.\ 1979; Stone 1982), we derived a background of 7.3~DN per 
WF~pixel. After correcting for this offset, the integrated magnitude of the 
galaxy within our field of view is M$_V=9.7$~mag. Assuming the distance modulus 
derived from the GCLF (m--M~$=29.8$~mag), we calculate a specific frequency of 
$1.3\pm0.1$ in the inner region of the galaxy (within a radius of $\approx3$~Kpc
of the center of the galaxy). For comparison, the Harris (1996) and Hanes \&
Harris (1986) estimate of the specific frequency corrected for a distance 
modulus of m--M~$=29.8$~mag is S$_N=3.1\pm0.5$. The lower value of specific 
frequency in the inner region of the galaxy implies that the cluster density
distribution near the center of the galaxy if flatter than the stellar light 
distribution. This is consistent with observations of the clusters systems in 
other galaxies which show a similar effect (Forbes et~al.\ 1996).  The low 
value of S$_N$ is consistent with the general trend of early type galaxies 
having higher specific frequency than late types (Harris 1991).

\subsection{Shape of the cluster distribution}
   
A cursory glance at Fig.~2 reveals that the cluster distribution is flattened 
and mimics the luminous component of the galaxy itself. This is in  contrast to 
Hanes \& Harris (1986), who find a lower surface density of clusters along the 
plane of the disk. This apparent discrepancy is probably due to their lower 
completeness limit along the major axis where the higher background has a larger 
effect on ground based observations than on HST observations.  

In the top panel of Fig.~7 we plot the number of cluster candidates within 65~arcsecs 
of the center of the galaxy, binned in 25$\arcdeg$ sectors, as a function of 
angle from the major axis. We can clearly see that the spatial distribution of 
globular clusters is flattened and that the axes of the cluster density 
distribution are roughly colinear with the light profile.  In order to quantify 
the ellipticity of the clusters distribution we created a fake image, in which 
we added a Gaussian source at the location of each cluster, and then smoothed it 
by convolving with a wide Gaussian of FWHM $\approx30''$. We find that while 
galaxy light follows an E7 (Hubble type)  distribution, the globular cluster 
distribution is less flattened. At a distance of 60$''$ from the center of the 
galaxy along the major axis the cluster distribution has an E$5.1\pm0.3$ shape.
One concern is that smoothing of the fake image by a broad Gaussian may lead us 
to significantly underestimate the ellipticity of the clusters distribution. In 
order to test this we generated fake cluster distributions with varying 
ellipticities and position angles and approximately the same radial density 
profile as the NGC~3115 clusters. We created smoothed images of the cluster
distributions using the exact same parameters as for the NGC~3115 clusters,
and measured the ellipticity of the resulting image. We found that while an
extremely flattened  input E8 distribution was measured as a $\sim$E6.5 
distribution after smoothing, the difference between the ellipticity of the
input distribution and the smoothed image decreased with ellipticity, with
an input E2 distribution of fake clusters appeared unchanged by the smoothing
process. We also tested simulated cluster distributions in a wide range of 
position angles on various locations on the image and found that we could 
measure the position angle of the simulated distributions to within 5$\arcdeg$.
\begin{figure*}%7
\centerline{\psfig{figure=pp1287f7.eps,width=5in,angle=-90}}
\caption{ The top panel plots the histograms of the cluster position angles within 
65~arcsecs of the center of the galaxy. The two bottom panels show the locations 
of the red and blue clusters within the WFPC2 chip. Both the blue and red cluster 
populations are found to be flattened, though the red clusters distribution appears  
to be much flatter than the blue cluster distribution.
\label{fig7}}
\end{figure*}

Significantly, we find that both the red and the blue clusters have flattened 
density distributions. Elliptical fits to the smoothed images of the red and 
blue cluster distributions reveal that the red clusters roughly follow an 
E$6.4\pm0.3$ Hubble profile at a major axis distance of 60$''$, while the 
blue clusters follow a less flattened E$3.2\pm0.3$ profile at the same major 
axis distance. Taking into consideration the effect of smoothing by a Gaussian,
we estimate that the actual shape of the red cluster distribution is 
$\sim$E$7.5\pm0.5$, while that of the the blue clusters is $\sim$E$3.7\pm0.5$.
This difference in ellipticities can also be seen in the bottom panels of 
Fig.~7, in which the positions of inividual red and blue cluster candidates on 
the WFPC2 chip are shown. In the few galaxies in which the shapes of the red and 
blue clusters have been measured (NGC~1380---Kissler Patig et~al., NGC~4472---Lee, 
Kim \& Geisler 1998) similar effects have been observed e.g.\ the metal rich, 
red clusters  appear to closely follow the galactic light distribution while the 
blue clusters are less flattened. 

The large ellipticity of the red clusters suggests that it is most likely 
associated with the thick disk component of NGC~3115, while the shape of the 
blue cluster distribution is closer to the $\sim$E4 distribution of the 
spheroidal bulge light distribution (Silva et~al.\ 1989). While it seems fairly 
certain that the blue clusters are linked to a spheroidal component, we cannot 
say with certainty whether they are linked to the bulge or the halo due to the
limited field of view of the WFPC2 camera. The low metallicity of the blue 
clusters may lead one to presume that they are linked to the halo, however 
NGC~3115 has such a strong radial color gradient that the inferred color of the 
bulge is very dependent on bulge-disk decomposition method. We shall therefore 
refer to the blue clusters as a bulge/halo component.   It seems very likely 
that the processes that formed the thick disk and the bulge/halo of the galaxy 
are also responsible for the formation of the two populations of globular 
clusters.

Silva et~al.\ (1989)  found that the disk of NGC~3115 is bluer than the bulge 
and that the age of the disk is intermediate between those of typical S0~galaxies 
and  typical spirals. The $4\pm3$~Gyr difference in age between the red and the 
blue clusters is consistent with this supposition and further bolsters the case 
for  the red clusters being associated with the younger disk and for the blue 
ones to be linked with the older bulge/halo of the galaxy. 

The metal-poor blue clusters are generally assumed to have formed during the 
initial collapse phase of a galaxy (which may merge with another galaxy to
form a new one of a different Hubble type). The more interesting question is
what triggered the formation of the red clusters at a later epoch. One possibility 
is that they formed during a merger event. Though there are no tell tale 
signatures of a merger, e.g.\ loops, shells etc., in broad band images of NGC~3115 
(Silva et~al.\ 1989; Scorza \& Bender 1995), if the  red clusters were formed 
during a merger event $11\pm3$~Gyr ago, it is likely that any anomalies in the 
structure of the galaxy have subsided since then. A major merger between two 
galaxies of roughly equal mass is unlikely since it would most probably have 
destroyed the disk. However, a minor merger with a smaller companion might 
trigger the formation of the second generation of red, metal poor clusters and 
result in the formation of the thick disk, with the same ellipticity as the 
red clusters, at the same time. Schweizer (1990) proposed that S0~galaxies are 
likely to have formed by the merger of two gas rich galaxies with mass ratios 
around 0.3 (1:3 mergers) or slightly less. Recent numerical simulations by 
Bekki (1998) support this model of S0~galaxy formation. One possible scenario 
is that the merger event causes a starburst in the disk which heats the gas
and induces  star formation in Giant Molecular Clouds, most of which which are 
likely to be associated with the disk of the larger of the two merging galaxies. 
This disk may be thicker due to the interaction, resulting in the E7 shape of 
the red clusters that we observe.
	
Interestingly, the characteristics of the red and blue clusters are very
reminiscent of the globular clusters in the Milky Way. As in NGC~3115, the 
metal rich, red clusters in the Milky Way are known to lie in a flattened 
distribution  are linked to the thick disk of the galaxy (Zinn 1996), while the 
metal poor blue cluster distribution resembles the halo. Extending the analogy 
further, the red disk population of the  Milky Way globular clusters is estimated 
to be $\sim2$~Gyr younger than the blue population, which is comparable to the 
$4\pm3$~Gyr age difference between the corresponding clusters in NGC~3115 (see 
Ashman \& Zepf 1998 for a detailed discussion on the relative ages of Milky Way 
clusters). Even though NGC~3115 and the Milky Way have significantly different 
morphological properties, the similarities between  their globular cluster systems 
leaves open the possibility that the processes that created the two populations 
of globular clusters are the same in both galaxies.   	

We also note that  the position angle of the major axis of the stellar component
is $45\pm2\arcdeg$ East of North, while the major axis of the cluster density 
distribution is slightly offset, and lies $32\pm5\arcdeg$ East of North. 
Furthermore, both  the red and the blue clusters distributions appear to be 
misaligned with the galaxy by approximately the same amount. It seems unlikely 
that the axes of symmetry of the globular clusters and the stellar component are 
inclined to each other, moreover, no current theory of GCS formation can explain 
such a misalignment. On the other hand, it is possible that the apparent angular 
separation is caused by projection effects. In other words, the globular cluster 
system of NGC~3115, and possibly the stellar population of the galaxy is 
tri-axial. There is some support for this conjecture in the isophotal analysis 
of  Capaccioli, Held \& Nieto (1987). They found that the bulge component of 
NGC~3115 has twisted isophotes although the position angle of the major axis 
rotated by only $2\arcdeg$ over 4$'$.      

\subsection{Sizes of the Clusters}

As we noted in \S2.1, the profiles of the globular clusters in NGC~3115 are 
broader than stellar profiles, indicating that they are spatially resolved.
The cumulative magnitude  profiles of an unresolved object and selected cluster
candidates from the GC field are plotted in Fig.~8. We can see from the figure
that the  profiles of globular cluster candidates can be clearly distinguished 
from those of stars and that we can attempt to measure the sizes of the clusters. 
The technique used to determine the size of an individual cluster candidate is 
described below.
\begin{figure*}%8
\centerline{\psfig{figure=pp1287f8.eps,width=4.5in,angle=-90}}
\caption{ The curve of growth for stars and selected cluster candidates in the 
PC and WF chips. The sizes of individual clusters have been determined by fitting 
them to PSF convolved king models. Note that the PC pixels are smaller than the 
WF pixels and that the PC profile is shallower than the WF profile because we are 
comparing them at similar aperture radii. 
\label{fig8}}
\end{figure*}

In order to model the light distribution of the globular clusters in NGC~3115 
we assume that they are similar to the Milky Way clusters and that the surface 
brightness profile can described by King models (King 1962). The size of a 
cluster can then be defined by two parameters, the King radius (r$_0$), and the 
concentration parameter c~$=\log_{10}(\frac{r_{tidal}}{r_0}$). However, there 
are two important considerations to be made before we model the observed light
distribution of the clusters. We must take into account the fact that the WFPC2 
point spread function varies across each of the chips. Also, the undersampling 
of the PSF by the WFPC2 camera has the unpleasant effect of modifying the 
observed light profile depending on the location of the peak  within a pixel. 

We created 4-fold oversampled Tiny Tim models (Krist 1995) of the PC and WF PSFs 
at five evenly spaced locations on each of the chips. Each of these PSFs were 
convolved  by a range of King models, varying c between 0.5 and 2.5, and r$_0$ 
between .5 and 16. We then resampled the models to normal size for eight different 
locations of the peak within a pixel. For each of the models, we performed 
photometry for a range of aperture radii to create a cumulative light profile 
and normalized all the profiles in our library to the observed light within 
3.3~pixels. The profiles of the candidate objects were created using the same 
aperture photometry parameters.  We  deduced the structural parameters of an 
individual cluster by  finding the best fit (in a  least square sense) within 
the set of models that were closest to the  candidate cluster with respect to 
the chip location and centering within a pixel. We find that most of the clusters 
have a concentration index c between 0.5 and 2.0 with a median of 1.25. This is 
consistent with the median King model concentration parameter of old globular 
clusters in the Milky Way (Harris 1996), M31 (Crampton et~al.\ 1985; Grillmair 
et~al.\ 1996), and the LMC (Elson \& Freeman 1985) all lie within this range.  

At this point we conducted some tests to ascertain how reliably we can measure 
the concentration index and core radius of individual clusters. We added model
clusters with various values of c and r$_0$ to the NGC~3115 field and used our 
size fitting algorithm to determine the structural parameters. We discovered
that while the mean measured structural parameters of identical models is close
to the input values, there are significant correlated errors when when we fit
c and r$_0$ simultaneously. On the other hand, when we restricted the fits to a single 
concentration index, we were able to measure the relative sizes of the clusters 
candidates reliably even in regions with a strong galaxy background. The mean 
half light radii (r$_h$) of the cluster candidates in the PC and WF, for c~$=1.25$ 
and c~$=1.5$, are shown in Table~1. The measured value of $r_h\approx2$~pc is 
comparable to the mean half light radius of $\approx3$~pc for the brightest 
Milky Way clusters (van den Bergh 1996). As in the Milky Way, we find no obvious 
correlation between the half light radii of the clusters and the brightness 
(mass). This implies that the clusters with larger r$_h$ are more diffuse.

Table~1 also shows the average sizes of the blue and red clusters in the PC and 
WF chips. Interestingly, we find that the blue clusters are statistically larger 
than the red clusters in both the WF and the PC.  We have observed a similar 
effect in the M87 globular cluster system (Kundu et~al.\ 1998), though it was 
less obvious due to the larger distance to M87. To verify the reliability of the 
size difference we developed a different algorithm  to  estimate the size of the 
clusters.  First, we  obtained  counts in a set of radial bins around the 
candidate objects and calculated the FWHM of  Gaussian curve that best fit the 
intensity distribution of each individual cluster.  WFPC2 observations of a set 
of program stars from a  field in $\Omega$~Cen were then convolved with Gaussian 
distributions of varying widths and fitted with the same routine to establish a 
relationship between the convolved width and resulting FWHM. The calculated sizes 
of the blue clusters in NGC~3115 are larger than the red clusters in all the 
chips, in both the V and I~filters using this method.    
 
We tested 5 independent star fields in $\Omega$~Cen to see whether this might 
be an instrumental effect. In each of the test images we selected 
unresolved/barely resolved objects with colors between $0.8<$~V--I~$<1.4$ and 
used our algorithms to determine the size. We find that the size of objects with 
V--I~$>1.06$~mag is statistically identical to the size of objects with V--I~$<1.06$ 
in both the V and I~filters. We also tested each quadrant of all four WFPC2 
chips and discovered no position dependent effect that may be inducing the difference 
in size between red and blue objects.  We therefore conclude that the observed size 
difference between the red and blue clusters is real and not due to instrumental 
biases.  van den Bergh (1994) showed that the half-width of the Milky Way globular 
clusters increases with galactocentric distance, consequently the red clusters 
which are nearer to the center of the Milky Way are smaller than the blue ones. 
Since the blue clusters in NGC~3115 are associated with the spheroidal bulge/halo 
and are on average more distant from the center of the galaxy than the red disk 
clusters, the difference in size between the red and the blue clusters may be 
indicative of a similar relationship in NGC~3115 and M87.   The physical reason 
for the difference in sizes between the red and blue clusters remains unclear. It may 
either be a relic of  the formation mechanism of the  clusters (e.g.\ if the 
initial size of the clusters depends upon its galactocentric distance), or a 
signature of the relative efficiency of various destruction mechanisms e.g.\ 
since the red clusters stay within the gravitational potential of the disk for 
a longer time than the blue ones, they may lose more stars from the outer regions 
of the clusters due to more efficient evaporation which in turn could induce 
structural changes in these clusters. 
\begin{figure*}%9
\centerline{\psfig{figure=pp1287f9.eps,width=6in,angle=-90}}
\caption{ The top panels plot the half light radius and  magnitude (mass) of 
the clusters as a function of the distance from the major axis while the bottom 
panels show the same quantities as a function of the distance from the minor 
axis. There is no spatial variation in the brightness (mass) of the clusters 
that would support dynamical friction. The sizes show no radial trend with 
respect to distance from the minor axis. However, most of the clusters with 
r$_h>3$~pc are within 40~arcsecs of the major axis. This {\it might\/} be a 
signature of bulge and/or disk shocking.  
\label{fig9}}
\end{figure*}

\subsection{Radial Variation in the Sizes and Magnitudes}

In the upper panels of Fig.~9 we plot the cluster sizes and magnitudes as a 
function of the distance from the semi-major axis while the lower panels track 
the sizes and magnitudes with respect to the minor axis. There is no obvious 
radial variation in the magnitudes of the clusters with respect to distance from 
the center. Theory predicts that cluster destruction mechanisms like dynamical 
friction, which preferentially destroys high mass clusters, and tidal shocking, 
which acts more efficiently on low density, low mass clusters, act most efficiently 
in the inner regions of galaxies.  Since we see no systematic changes in the 
luminosity distribution (or mass distribution for constant mass to light ratio),  
we conclude that we see no clear evidence of dynamical evolution of the clusters 
and that the mass distribution is similar everywhere.  

As an aside, we note there are two abnormally large objects with half light 
radii greater than 6~pc. On close inspection the smaller of the two objects 
appears to have some nebulosity and it may be a background galaxy. On the other 
the candidate with a  half light radius of 13~pc has a very well defined 
circular shape and fulfills all the criteria for a globular cluster (M$v=-10$, 
V--I~$=1.05$). It would be very interesting to obtain the velocity of this object 
to ascertain whether it is indeed a part of the NGC~3115 globular cluster system 
or a background object. The candidate is offset from the center of NGC~3115 by
 $\Delta$ R.A.~$= -1.3664$~secs, $\Delta$~Dec~$=-1'$~0.884$"$. 

On the face of it, the size distribution of the clusters appears to be evenly 
distributed in space. However, a closer inspection of the size vs.\ distance 
from major axis plot reveals that most of the clusters with a half light radius 
greater than 3~pc are located within 40~arcsecs of the major axis of the galaxy. 
Though this result is not very statistically very significant ($\approx2$ sigma), 
it suggests that clusters that lie along the major axis of the galaxy are smaller 
than those that are located along the minor axis could be indicative of the 
relative efficiency of cluster  destruction mechanisms along the two axes. 

\section{SUMMARY}
 
\begin{itemize}
\item We have identified 144 globular cluster candidates in a WFPC2 field of the 
inner regions of NGC~3115. 133 of the objects lie within a very narrow range of 
color between $0.8<$~V--I~$<1.4$.
\item  The color distribution of the globular clusters is bimodal with a blue 
peak at V--I~$= 0.96$~mag and a red peak at V--I~$=1.17$~mag. There are roughly 
equal number of blue and red clusters.
\item The metallicity of the blue clusters is Fe/H~$=-1.36$~dex while that of 
the red clusters is Fe/H~$=-0.37$~dex
\item Assuming that the blue peak is 15~Gyr old, we derive an age of
$11\pm3$~Gyr for the red clusters.
\item The age and metallicity of the globular clusters are consistent with the 
red giant branch stars in the field suggesting that the stars and clusters 
formed roughly coevally.
\item The peak of the globular cluster luminosity function is at 
V~$=22.37\pm0.05$~mag. This yields a distance magnitude, m--M~$=29.8\pm0.3$~mag. 
This distance is in fair agreement with the distance estimate from the tip of 
the red giant (m--M~$=30.2\pm0.3$~mag) and other previous measurements.
\item The local specific frequency of the galaxy is $1.3\pm0.1$ which is lower 
than the global specific frequency of $\sim3$, implying that the radial density
profile of the globular cluster system is flatter than the stellar light profile 
in the inner regions of NGC~3115.
\item We find that both the red and the blue cluster distributions are flattened 
at all radii and roughly colinear with the light distribution.  While the blue 
clusters follow a E$3.7\pm0.5$ profile at a major axis distance of 1$'$, and 
appear to be related to the bulge/halo, the red clusters follow a more flattened  
E$7.5\pm0.5$ profile at the same major axis distance, and are most likely related 
to the thick disk component of the galaxy.
\item The red clusters may have formed during the accretion of a companion/a minor 
merger event $\sim11$~Gyr ago. 
\item The half light radius of the clusters is $\approx2$~pc which is similar
to the Milky Way globular clusters. The red clusters are $\approx20$\% smaller 
than the blue clusters.  
\end{itemize}

\acknow
We would like to thank Laura Fullton for useful discussions on Milky Way 
globular clusters and for providing us access to some of her recent 
color-magnitude plots of them, and the anonymous referee for many 
useful comments. 

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\onecolumn

%\documentstyle[11pt,apjpt4]{article}
%\renewcommand{\baselinestretch}{2}
%\begin{document}
%\newpage
\begin{center}
\begin{tabular}{ccccccc}
%\tablenum{1}
%\tablewidth{0pt}
\multicolumn{7}{c}{\bf TABLE 1. EFFECTIVE RADIUS OF THE CLUSTER CANDIATES}\\ 
\multicolumn{7}{c}{\bf USING KING MODEL FITS}\\
%\startdata
\\
\hline
\hline
&\multicolumn{6}{c}{Average half light radius of the cluster candidates}\\
&\multicolumn{3}{c}{Planetary Camera} &\multicolumn{3}{c}{Wide-Field Camera}\\
King Profile &r$_{h}$ All   &r$_{h}$ blue$^{\rm a}$  &r$_{h}$ red$^{\rm b}$  
&r$_{h}$ All &r$_{h}$ blue$^{\rm a,c}$ &r$_{h}$ red$^{\rm b}$\\
c~$=\log10(\frac{r_{tidal}}{r_0}$) &(pc) &(pc) &(pc) &(pc) &(pc) &(pc)\\
\hline
1.2\rlap{5} &$2.1\pm0.2$ &$2.7\pm0.4$ &$1.7\pm0.2$ &$2.13\pm0.07$ &$2.22\pm0.11$ &$2.05\pm0.10$\\
1.5         &$2.6\pm0.4$ &$3.5\pm0.8$ &$1.9\pm0.4$ &$2.02\pm0.09$ &$2.08\pm0.13$ &$1.97\pm0.11$\\
%\enddata
\hline
\noalign{\smallskip}
\end{tabular}
\end{center}
\begin{itemize}
\item[$^{\rm a}$] $0.8<$~V--I~$<1.06$
\item[$^{\rm b}$] $1.06<$~V--I~$<1.4$
\item[$^{\rm c}$] The two anomalously large sources in the WF (with half light 
radius greater than 6~pc) have not been included in the average. Both of these 
objects are blue in color and the difference between the sizes of the blue and 
the red clusters would be larger if they were to be included.
\end{itemize}
\smallskip
\hspace{.2in}Note.---A distance of 9.1~Mpc to NGC~3115 has been adopted in order 
to calculate the size of the clusters.

\end{document}
