NGC 3125−1: The Most Extreme Wolf-Rayet Star Cluster Known in the Local Universe1

Rupali Chandar and Claus Leitherer

Space Telescope Science Institute, 3700 San Martin Drive, Baltimore, Maryland 21218

[email protected] & [email protected]

and

Christy A. Tremonti

Steward Observatory, 933 N. Cherry Ave., Tucson, AZ, 85721

[email protected]

ABSTRACT

We use Space Telescope Imaging Spectrograph long-slit ultraviolet spec- troscopy of local starburst galaxies to study the massive star content of a represen- tative sample of “super star” clusters, with a primary focus on their Wolf-Rayet (WR) content as measured from the He II λ1640 emission feature. The goals of this work are three-fold. First, we quantify the WR and O star content for selected massive young star clusters. These results are compared with similar estimates made from optical spectroscopy and available in the literature. We conclude that the He II λ4686 equivalent width is a poor diagnostic measure of the true WR content. Second, we present the strongest known He II λ1640 emis- sion feature in a local starburst galaxy. This feature is clearly of stellar origin in the massive cluster NGC 3125-1, as it is broadened (∼ 1000 km s−1). Strong N IV] λ1488 and N V λ1720 emission lines commonly found in the spectra of individual Wolf-Rayet stars of WN subtype are also observed in the spectrum of NGC 3125-1. Finally, we create empirical spectral templates to gain a basic understanding of the recently observed strong He II λ1640 feature seen in Lyman Break Galaxies (LBG) at redshifts z ∼ 3. The UV field observed in local star- bursts provides a good overall match to the continuum and weak photospheric

1Based on observations with the NASA/ESA 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. – 2 –

features in LBGs in the spectral range λλ1300 − 1700, but cannot reproduce the He II λ1640 emission seen in the composite LGB sample of Shapley et al. An additional (ad hoc) 10-15% contribution from “extreme” Wolf-Rayet clusters similar to NGC 3125-A on top of the field provides a good match to the strength of this feature.

Subject headings: galaxies: individual (NGC 3125) — galaxies: starburst — galaxies: star clusters – galaxies: stellar content

1. INTRODUCTION

Understanding the formation and evolution of starbursts is a problem of fundamental importance to astrophysics. Massive bursts of star formation quickly enrich the host with metals, as gas processed in the interiors of massive stars is returned to the ISM by super- novae. Hubble Space Telescope ultraviolet/optical imaging has resolved the sites of active star formation in local (z < 0.1) UV-bright starbursts into numerous compact “super star” clusters (SSCs) contributing ∼ 20% of the total UV light, and extended diffuse UV “field” light, which makes up the remaining 80% (e.g., Meurer et al. 1995). A subset of starburst galaxies show signatures of Wolf-Rayet (WR) stars, which are the evolved descendants of

O stars more massive than M > 20 − 30M , and are characterized by high mass loss from fast, dense stellar winds. These may be ideal laboratories for studies of starburst properties, since they are known to harbor the most massive known stars (O stars and their evolved WR descendants), allowing us to probe the youngest stellar populations. The use of the UV spectral range (to study O and WR stars) has a number of advantages over similar studies conducted using optical spectroscopy: i) the massive stars which we are probing are the direct contributors to the UV continuum, which is essentially unaffected by an underlying stellar population, unlike the situation in the optical; ii) the UV spectral region gives more precise ages; iii) WR emission line equivalent widths (EW) in the UV are undiluted by an underlying stellar population; iv) in the UV, it is possible to measure the extinction intrinsic to the starburst directly from the same stars which contribute to the line flux, rather than inferring it from the nebular emission features; and v) a better quantification of the rest-frame UV is essential to understand the properties of high redshift galaxies. Rest-frame UV observations of high redshift galaxies are now beginning to directly reveal the dominant stellar populations forming in the early universe. One specific example is the discovery of strong, broad He II λ1640 emission in the composite spectrum of hundreds of – 3 – z ∼ 3 Lyman Break Galaxies (LBG) (e.g., Shapley et al. 2003). Broad stellar He II emission is primarily produced in fast, dense winds from WR stars. However, when compared with observations of local starbursts (e.g., Heckman et al. 1998), the He II λ1640 emission in the high redshift galaxy composite is significantly stronger. We are conducting an extensive study of the stellar content of local starburst galaxies, using longslit HST spectroscopy in the UV. This allows us to probe the properties of individ- ual clusters, thus isolating single age stellar populations, and relieving concerns arising when dealing with mixed age populations. In this work, we provide a census of the massive (O and WR) star content of a representative sample of SSCs in local starburst galaxies. In order to study the WR content in individual massive young star clusters formed in starburst galaxies, we focus on the He II λ1640 emission feature, which is observed in a number of objects. This paper is organized as follows: section 2 describes sample selection, data reduction, and basic measurements of the stellar content of our SSCs; section 3 compiles properties of individual Galactic and Large Magellanic Cloud (LMC) WR stars for comparison with stellar clusters; section 4 estimates the number of O and WR stars, and compares with existing optical data to form a coherent picture of the massive star content in individual SSCs; in section 5 we use empirical templates created from different local stellar populations (e.g., individual clusters at different ages, stellar field) to understand the composite rest-frame UV spectra of high redshift LBGs from the Shapley et al. (2001b) sample. Finally, in section 6 we present a summary and conclusions.

2. SAMPLE, DATA REDUCTION, AND MEASUREMENTS

2.1. Local starburst galaxy sample

As part of a larger project to quantify the stellar and interstellar properties of local galax- ies undergoing active star formation (Chandar et al. 2004, in prep.), we have obtained HST STIS long-slit far- and near- ultraviolet spectra for 15 local starburst galaxies; further infor- mation on the program can be found under proposal 9036 at http://www.stsci.edu.edu/public/propinfo.html. Target galaxies were selected to cover a broad range of morphologies, chemical composition, and luminosity. Most importantly, target selection was not based on WR content. We utilized the 5200 × 0.200 slit (projected onto the 25002 MAMA detector) and the G140L and G230L gratings to obtain the best compromise between throughput and spectral resolution. This combination of gratings provides continuous wavelength coverage from 1175 to 3100 A,˚ with a velocity resolution of 100 − 200 km s−1. The STIS MAMA detectors have a plate −1 scale of 0.02400 pix−1, and average pixel scales of 0.584 A˚ pix for the FUV/G140L and −1 1.548 A˚ pix for the NUV/G230L gratings. All target galaxies have been chosen to be – 4 – dominated by star formation; the nuclear starbursts in NGC 6764 and NGC 5102 show the clear signature of massive stars, with no evidence for Seyfert/Liner activity in the UV. To this sample we add three starburst galaxies previously observed with STIS: NGC 3049 (ob- served with the 0.500 slit), NGC 5253 (0.100 slit), and Tololo89 (hereafter Tol89; 0.500 slit) and available in the HST data archive. I Zw18 which was studied by Brown et al. (2002) is not included due to the low S/N of the observations. Basic information on each galaxy used for this study, such as distance, abundance and foreground reddening, as well as the physical area covered by the long-slit observations, are presented in Table 1. When available, we use direct distances from HST observations (e.g., tip of the red giant branch). The remaining distances were also assembled from the literature, but were determined from a variety of techniques. In this work we focus on data taken in the FUV, since the best diagnostics of stellar properties are found in this spectral region. Each of the starburst galaxies contains a number of individual SSCs, which are clearly visible on available HST WFPC2 and FOC images. We focused on bright SSCs in each galaxy by constraining the pointing and orientation of the STIS slit. Here we briefly describe target selection (we have concentrated on luminous clusters in this work; a future paper will present the entire cluster catalog). In NGC 5102 and NGC 6764 the nuclear starbursts were observed; in NGC 3310 the slit included a giant extra-galactic HII region, and in NGC 3049, NGC 4214, and NGC 5253 we study the UV dominant clusters. Starburst regions A in NGC 3125 and He 2-10 are included in this study. Targeting for the rest of the galaxies in our sample was made primarily from HST FOC UV imaging, where UV luminous, compact SSCs were included, with the goal of aligning the slit to cover as many such objects as possible. Cluster coordinates, ages and extinctions intrinsic to the starburst (age and extinction derivation are described below) are given in Table 2. For the purposes of this study, we wish to isolate coeval events of star formation. In general, we include the brightest (generally the targeted) cluster in each galaxy which fell in our slit, unless otherwise noted. We emphasize that these objects were not originally selected based on their WR features, but based on their UV luminosity. These clusters do however, tend to be the youngest/most massive objects, thus they are most likely to contain relatively large numbers of massive O and WR stars. For He 2-10, we discuss both the most luminous cluster, as well as our extraction of the entire starburst region A covered by the slit. In Chandar et al. (2003), we determined that the dominant stellar population in both the clusters and field has a burst age of 4 − 5 Myr, so this is in effect a single age template which provides a better spatial match to available (ground-based) optical spectroscopic observations. To facilitate comparison with results derived from optical spectroscopy, we have created a composite spectrum from the four brightest clusters observed in NGC 1741, and also study the most luminous cluster separately. Finally, for Mkn33 and – 5 –

NGC 4449, we combine the spectra of six and nine faint clusters, respectively, in order to increase S/N.

2.2. Reduction and pre-processing

The data were retrieved from the HST archive. When multiple exposures of a single object existed, we checked for spatial offsets between exposures (none were found), co-added the raw data files, and updated the header to reflect the total exposure time. These co-added raw files were then processed through the CALSTIS pipeline, which rebins the spectra and provides a global detector linearity correction, dark subtraction, flat fielding, wavelength calibration and conversion to absolute flux units. For NGC 4449, NGC 1741, and Mkn33 we present co-added cluster spectra. In these cases, clusters were detected as locations with a minimum 5σ detection above the local “background” when considering a cut along the spatial direction. Because our target galaxies have a range in distance (from ∼ 4 Mpc to ∼ 50 Mpc), cluster profiles also vary. We extracted all clusters assuming that they are fully extended objects, with no correction made for slit losses. The number of binned (spatial) pixels which form a given cluster spectrum were chosen by eye, and are given in column 4 of Table 2. The minimum extraction height is 11 pixels (the recommended extraction height for a point source; STIS Data Handbook). In Chandar et al. (2003) we studied the effect of varying the extraction height on He 2-10 clusters, and found little variation in the derived cluster age, EB−V , and mass for moderately luminous clusters. Our extraction technique provides a lower limit to the actual cluster fluxes, since we make no correction for slit losses even though star clusters are generally extended in HST observations out to ∼ 10 Mpc. Further processing steps on the extracted one-dimensional spectra included correction for slight relative wavelength shifts, deredshifting (values taken from NED) and rebinning to the wavelength scale of the Starbust99 models (0.75 A˚ pixel−1), which were utilized for comparison with the observations. We then corrected the extracted spectra for strong geocoronal emission at Lyα and OI λ1302 by assuming that an outer portion from the two dimensional spectrum represents a purely geocoronal spectrum near the Lyα and OI λ1302 features. Most of the clusters in our sample are neither point nor fully extended sources. Ac- cording to the STIS manual, the G140L grating has an instrumental resolution (FWHM) of 1.5 pixels for point sources, and 4, 8, and 20 pixels for fully extended sources filling the 0.100, 0.200, and 0.500 slits respectively. We use cluster NGC 3125-1 as an example to estimate – 6 – the resolution of a typical cluster observed with the 0.200 slit. Measurements of the surface brightness profile along the slit establishes that while this cluster is clearly resolved, the FWHM of 0.1100 is smaller than the 0.200 slit width. Therefore the resolution is intermediate between that for a point source (0.88 A)˚ and a fully extended source (4.67 A).˚ We estimate a value of ∼ 2.5 − 2.7 A˚ as an appropriate resolution for our cluster spectra observed with the 0.200 slit.

2.3. Cluster age and extinction determination

Cluster ages and extinctions were determined by comparing the extracted spectra with the Starburst99 models (Leitherer et al. 1999). The Starburst99 models have been optimized to reproduce many spectrophotometric properties of galaxies with active star formation, and the UV spectral library is available for two metallicities: solar (Z = 0.2) and LMC/SMC metallicity (∼ 1/4 Z ). Details of the input stellar parameters to Starburst99 can be found in Leitherer et al. (1999); here we briefly summarize the model parameters used in this work. For an instantaneous burst of star formation (appropriate for SSCs) we adopted a standard Salpeter (1955) initial mass function (IMF), with lower and upper masses of

1 M and 100 M , respectively. Arguments in favor of a Salpeter IMF are summarized in Leitherer (1998), and a discussion of various studies available in the literature can be found in Tremonti et al. (2001). Theoretical high mass loss rates due to stellar evolution were assumed (Meynet et al. 1994). The model most appropriate for the overall galaxy metallicity as measured from the nebular emission (taken from the literature and given in Table 1), was adopted for each cluster. These metallicities are strongly weighted towards bright HII regions, where we obtain our UV spectra. In order to determine intrinsic properties for a stellar population, the one-dimensional spectra must be corrected for the effect of (foreground) Galactic extinction, as well as for the dust obscuration intrinsic to the starburst itself. We use the Fitzpatrick (1999) Galactic extinction law to account for foreground reddening (EB−V values for each galaxy are taken from the Schlegel et al. 1998 maps). One important advantage of using the UV spectral region is that the reddening mea- surement is made directly on flux coming from the massive stars which we are interested in studying, not from surrounding ionized gas and the resulting nebular lines. We com- pared the observed continuum of each SSC with stellar evolutionary models to determine the reddening internal to the starburst. The intrinsic FUV spectral distribution of young, unobscured single stellar populations follows a power law, with effective spectral index β, β where Fλ ∝ λ . Any deviation of the power law exponent from that predicted by theoretical – 7 – models is assumed to be due to the effects of dust. Because intervening ISM lines due to the Milky Way sometimes end up in useful portions of the continuum, we adopted the procedure developed by Tremonti et al. (2003, submitted) to subtract these lines from the spectrum prior to continuum fitting. Briefly, the strength of each ISM line from our Galaxy is modeled as a Gaussian with the median equivalent width taken from Savage et al. (2000), centered at zero velocity, and having a width set by the instrumental FWHM given in the previous section. This intervening ISM spectrum is then subtracted from each cluster spectrum.

In order to determine the age and reddening, EB−V , of each cluster, we adopted the following procedure. First we compared rectified cluster spectra with the appropriate metal- licity Starburst99 model, where the ISM lines are eliminated from consideration, and wind lines are weighted 1.5 times higher than the continuum (for details of the age fitting pro- cedure, see Tremonti et al. 2001). For young clusters, this procedure essentially uses O star features as a chronometer. We then measure the slope of the best fit (age) model, and adopt this as the intrinsic slope, β. The fit for β was performed over the spectral region 1240 − 1600 A.˚ To account for the spectral features (which are mostly in absorption), we performed the fit iteratively with rejection thresholds set at 2σ for the lower bound and a 3.5σ upper bound. Typically, β = −2.6  0.2 for young (≤ 10 Myr) starbursts. Finally, the extinction intrinsic to the starburst was determined by dereddening each (foreground extinction corrected) spectrum by the Calzetti et al. (2000) starburst obscuration law until the measured slope matched the adopted intrinsic value. The UV slope β measured after correction for foreground extinction but before applying the starburst obscuration law are given in column 5 of Table 2. Uncertainty in the age results in a small uncertainty in the value of the intrinsic slope β, as measured from the models. As a conservative estimate, we assume that β is known to within 0.1. The routine returns the best fit age (minimum χ2) and EB−V , which are given in columns 6 and 7 of Table 2. Age uncertainties were determined as follows: for most clusters, utilizing the bootstrap technique described in Tremonti et al. (2001) resulted in errors ∼ 1 Myr. Three clusters, Mkn36-1, Mkn209-1, and Tol1924-416-1 have formal best fit ages of 0 Myr, but a plot of χ2 vs. age shows that there is significant degeneracy with 6 Myr, which also give local minima. NGC 5102-1 is much older than the rest of the clusters presented here, and its age is poorly constrained. As a consistency check, we compared our derived ages for five SSCs presented here (NGC 3125-1, He 2-10-all, NGC 3049-1, NGC 5253-5, and Tol89-1), plus NGC 3125-2 (which will be included in a forthcoming paper), with those derived in Schaerer et al. (1999). Schaerer et al. (1999) derived cluster ages by comparing measured Hβ equivalent widths with population synthesis model predictions. With the exception of Tol89-1, our age estimates – 8 – are consistently 1.0-1.5 Myr younger. To assess whether our age estimates are superior to those derived from optical spectroscopy (as expected from the reliability of each technique), we made a visual inspection of the morphology of emission lines grouped into the so-called blue WR bump (i.e. N III/N V λ4640, C III/C IV λ4650, and He II λ4686) for the clusters mentioned above (as shown in Schaerer et al. 1999, Figure 7), with the synthetic starburst spectra with ages between 2.5 Myr and 4.5 Myr presented by Norris, Smith, & Crowther (2003). These simulations, using new O star atmospheres and fully theoretical, non-LTE line-blanketed WR spectra, suggest that the relative strength of the He II λ4686 line to the N III λ4640 line is highest at an age of 3.0 Myr. Qualitatively, the strengths of the optical features arising in Wolf-Rayet stars depicted in Figure 1 of Norris, Smith, & Crowther (2003) show remarkable similarity with the morphology of the blue bump for He 2-10, NGC 3049, NGC 3125-1, and Tol89, if we assume the ages derived in this work, which come primarily from UV lines arising in O stars. The ages derived in the Schaerer et al. (1999) study give poorer agreement with the Norris et al. (2003) models.

2.4. Measurement of 1640A˚ feature

Broad He II λ1640 emission is seen in the UV spectra of individual Galactic and Mag- ellanic Cloud WR stars (e.g., Conti & Morris 1990), but is not prevalent in the integrated spectra of galaxies due to the extremely short lifetime of WR stars. Note that a feature at 1640A˚ is common in IUE spectra; however this wavelength region is compromised by IUE instrument blemishes (Crenshaw et al. 1996). Here we describe our technique to measure the He II λ1640 emission, which is used later to estimate the WR content. There appears to be an emission feature in many of the spectra at the correct wavelength; however obtaining an accurate measure of the local continuum is not trivial in most cases. In order to isolate the He II λ1640 emission from WR stars, we implemented the following procedure. First, the WR spectra which are input to the Starburst99 evolutionary code were edited to remove the He II λ1640 line, by setting the flux in the λλ1635 − 1655 bandpass to the value of the local continuum. Although weaker He II λ1640 is present in early Of supergiants, we left these templates untouched, as one goal of this paper is to use the He II λ1640 emission to quantify the WR content of starburst galaxies. Next, the Starburst99 evolutionary synthesis code was run, assuming all the same parameters given in §2.3, but using the new, “He II-free” WR spectral templates. This generated a new instantaneous burst model at solar metallicity (the sub-solar metallicity models do not extend redward of 1600A).˚ We then subtracted the He II-free Starburst99 model with the best fit age (as given in Table 2) from the observations. We defined the – 9 –

He II λ1640 bandpass used to measure the EW to be λλ1632 − 1654. Figure 1 shows the results of our age fitting technique and He II-free model subtraction applied to our entire sample. Determining the uncertainties on our EW measurements is difficult. There may be sys- tematics in our continuum placement, there are uncertainties arising from comparing sub- solar metallicity clusters with the solar metallicity models, etc. We defined a comparison bandpass (1600 − 1630A)˚ adjacent to the He II λ1640 portion of the spectrum, and esti- mated EW uncertainties as follows. The measured EW of the comparison bandpass is added in quadrature to the formal errors, under the assumption that the continuum bandpass traces any systematics in the continuum placement. We measure the λ1640 EW by fitting a Gaus- sian to the data minus He II-free model spectrum, and recorded the result in Table 3 if the EW measurement is larger than the uncertainty. Formally, we measure He II λ1640 emission in all but five clusters, although most measurements have large uncertainties. The Gaussian fit also gives flux (converted to luminosities) and FWHM measurements, which are recorded in columns 3 and 6 of Table 3 respectively. Luminosities have not been corrected for slit losses (i.e. the tabulated luminosities represents a measurement only of the light falling in the slit). Formal uncertainties on the luminosity measurements have the same percent error as the EW measurements. However, there are large systematic uncertainties which domi- nate cluster luminosity estimates, primarily uncertainties in the galaxy distances (which are variable and difficult to quantify) and aperture corrections (discussed in §2.6).

2.5. Morphology of He II λ1640

We note a few interesting features from this exercise. First, there are essentially two different line morphologies of He II λ1640 emission in local starburst galaxies. NGC 3125- 1 and NGC 5253-5 show narrow He II emission, while the other SSCs with measured He II λ1640 emission show a broader profile. For the purposes of this work, we consider the ˚ He II λ1640 feature to be “narrow” if it has a FWHM <∼ 8A. If massive stars are forming, He II λ1640 can sometimes appear as a nebular recombination emission line (e.g., Garnett et al. 1991), which has a characteristic, narrow line morphology. In all cases He II λ1640 emis- sion in our cluster sample was broader than the estimated intrinsic line profile. In the case of NGC 3125-1, the measured FWHM of 5.90.4A˚ is clearly broader than the estimated in- strumental profile (∼ 2.5 − 2.7A).˚ The corrected width of the He II λ1640 line in this cluster (instrumental profile subtracted in quadrature from the measured width of the λ1640 line) is ∼ 1000 km s−1. Similar line widths are typically found in the latest subtypes of individual WN stars (see §3.1). The broadest line is observed in NGC 5996-1, which has a width of – 10 –

2600 km s−1, although there are large uncertainties on this measurement. This large value is still well within the range seen for individual late-type WN stars, as discussed in the next section. We conclude that the underlying mechanism exciting He II λ1640 emission in our cluster sample is non-nebular. NGC 5253-5 appears to be the exception to this statement. In this spectrum, we also see the O III] nebular emission lines at λ1661 and λ1666, which are not present in any of the other cluster spectra. Campbell, Terlevich, & Melnick (1986) suggest that all narrow-lined He II λ4686 galax- ies (e.g., I Zw18) have relatively low metal abundances. Based on our UV spectroscopic sample, this statement also holds for the morphology of the He II λ1640 line. The two narrowest measurements are made in NGC 3125 and NGC 5253, both with abundances of ∼ 1/5 solar. However, the converse statement is not true: not all metal-poor galaxies with detected λ1640 emission show a narrow morphology. An inspection of Table 3 establishes that NGC 3125-1 has (by far) the largest He II λ1640 equiv- alent width in our sample. In Figure 2 we show the entire UV spectrum of this object. The λ1640 emission is quite prominent, and appears significantly stronger than the C IV emission at λ1550. In addition to this feature, we also see N IV] λ1488 and N V λ1720 emission typically found in WR stars, consistent with our interpretation that the strong He II λ1640 emission arises in the winds of massive stars. In part, the prominence of this line is due to its narrow morphology, but our equivalent width and flux measurements estab- lish that in our survey, NGC 3125-1 has the strongest He II λ1640 emission due to massive stars in a local massive young cluster. The strength of the He II λ1640 feature relative to the C IV λ1550 strength in particular, is unique among local starburst clusters. Potentially, this result can provide an important new tool for understanding the stellar populations of high redshift galaxies, as discussed in section 5.

2.6. Uncertainties in total cluster flux measurements

In Table 3 we recorded the total luminosity of the He II λ1640 emission for each target cluster. The largest systematic uncertainties in these values comes from uncertainties in galaxy distances. This applies however, only to the light which actually falls in the slit. Because the measurements presented here are based on longslit spectroscopy, the largest uncertainty in the derived stellar content (which is based on measured fluxes) for the entire cluster comes from aperture corrections. This issue is of particular concern for the nearest galaxies, where the clusters are resolved. In principle, the total flux for the more distant clusters can be determined by treating them as point sources, and correcting for slit losses under this assumption. A different procedure however, is required for resolved clusters. Here, – 11 – we attempt to quantify by how much we underestimate the total cluster He II λ1640 line intensity. We assume that the UV flux outside of our slit has the same spectral signature as that within our slit. For clusters residing in galaxies located more than 10 Mpc away (essentially point sources), we use the following technique. The He II λ1640 fluxes were extracted using two different techniques: 1) the light was summed up in 11 pixels (as recommended by the STIS handbook when extracting point sources) with no correction made for slit losses; and 2) using the point source extraction routine X1D in IRAF, which does make the correction for slit losses, assuming the PSF of the source is not extended. Comparison of the results suggests that the fluxes listed in Table 3 underestimate the total cluster flux by a factor of ∼ 1.3 for these distant, unresolved sources. For nearby clusters (particularly in NGC 5253-5 and NGC 4214-1), the correction is much larger, since these clusters are resolved. The most extreme example is NGC 5253-5, which is one of the closest objects in our sample, and also one of the most extended. We used available HST WFPC2 F255W imaging to estimate the correction factor between the STIS spectra (the 0.100 wide slit was used) and total cluster light, (by comparing the number of counts measured from the WFPC2 image in an aperture mimicking the STIS slit, and one which takes into account the total amount of light in a 1.000 radius). We find that the amount of light falling in the STIS aperture accounts for only 1/9 of the total cluster UV light in an 1.000 radius, which means that the flux in Table 3 is less by nearly an order of magnitude when considering the object. For cluster NGC 4214-1, which is located at approximately the same distance as NGC 5253-5, but was observed with the 0.200 slit, we estimate (from an available WFPC2 F336W image), that we have underestimated the UV flux by a factor of ∼ 4. A similar exercise for NGC 3125-1 using our STIS two-dimensional images under the assumption of circular symmetry, implies that the total UV flux is underestimated by a factor ∼ 2. Although we have discussed by how much we underestimate the UV light for different objects, it is important to mention that the most massive stars, including WR stars, are typically found in the innermost cluster regions, while O and in particular B stars are found further out. Given this potential bias and the large and uncertain correction factors due to the small slit width, we do not make any explicit correction to the measured He II λ1640 lu- minosities in Table 3. Because the less massive but more numerous cluster O and B stars contribute signifi- cantly to the UV continuum, but may not be co-spatial with the WR stars, they may affect the measured equivalent width of the He II λ1640 feature, particularly because the UV spec- tra presented here are taken with narrow slits (typically 0.200 wide). If this is the case, – 12 – we would expect to see a dependence of the measured He II λ1640 equivalent width with distance, since the same physical slit will cover a larger intrinsic portion of a more distant cluster. We find that no such correlation exists, as exemplified by the relatively large equiv- alent width measurements for NGC 3125-1 and He 2-10, which are at intermediate distances in our sample (roughly 10 Mpc).

3. Properties of individual Wolf-Rayet stars

Our measurements of the He II λ1640 line for massive clusters in nearby starbursts indicate that cluster NGC 3125-1 has both an extremely large He II λ1640 EW and flux. Can this very large value be reconciled with what we know about the fluxes of various emission lines of individual WR stars? This galaxy has certainly been studied previously in the optical – have these studies failed to reveal the extremely large number of implied WR stars? We begin to address these issues by tabulating (and measuring when necessary) known properties of individual WN stars, since the presence of strong N IV] λ1488 and N V λ1720 emission indicates that NGC 3125-1 is dominated by WR stars of WN subtype.

3.1. Line widths

In Table 4 we compile the mean properties of various WN type stars. The FWHM measurements of the He II λ4686 line come from Tables 6 and 7 in Smith, Shara, & Moffat (1996). In that work, they reclassify known Galactic and LMC WR stars according to their new three-dimensional WN classification system, and provide a compilation of FWHM measurements for the λ4686 feature (see references in Smith et al. 1996). We grouped stars by WN subclass from WN2 to WN9, and determined the mean FWHM for the λ4686 line; these are presented in column 9 of Table 4. Stars with known companions (Breysacher, Azzopardi, & Testor 1999; van der Hucht 2001) were not used in this calculation. These values show clearly that as one moves to later-type WN stars, the FWHM of the λ4686 line decreases dramatically. We were unable to find line width measurements of the λ1640 feature in the literature. There may be concern that this feature shows P Cygni absorption+emission profiles in some individual WR stars, making a FWHM measurement more complicated and ambiguous. Nonetheless, in order to make a qualitative comparison with our cluster spectra (none of which show an obvious absorption signature in the He II λ1640 line), we downloaded available IUE SWP spectra for a number of WN stars. Because very few of these were taken in high – 13 – resolution mode, we use the low resolution observations taken in the large aperture. The range of measured FWHM values for one to three stars in each subclass are presented in column 5 of Table 4. λ1640 FWHM measurements show a similar trend as observed for the λ4686 line, where later WN subtypes have narrower line widths. However there is an offset from the λ4686 widths, such that the λ1640 lines at all subtypes appear narrower than their optical counterparts. A comparison of the values in Tables 3 and 4 suggests that our spectra are dominated by later-type WN stars, and the spectrum of NGC 3125-1 is dominated by the latest WN subclasses.

3.2. Line strengths and equivalent widths

From a study of high resolution IUE spectra of 14 WR stars, Willis et al. (1986) conclude that strong He II λ1640 emission is present in all WN subclasses, and is particularly strong in WN5 and WN6 stars. For individual stars, the λ1640 feature has a definite, but unsaturated P-Cygni absorption component in subclasses WN4-7, but appears only in emission in HD104994 (subtype WN3) and at WN8. Overall, the appearance of the UV spectra of individual WN stars as a function of WN subclass can be broadly understood in terms of a combination of ionization, elemental abundance and wind density effects. Here we compile average line strengths and equivalent widths for individual WN stars of different subclasses from the literature. Because of the effect that metallicity has on the wind parameters, we only consider stars residing in the LMC (not Galactic stars), since these have abundances very similar to those of NGC 3125, the starburst galaxy with the most secure He II λ1640 measurement. Additionally, by only using LMC stars, we avoid issues of distance uncertainties which dominate the intrinsic measurements for Galactic stars. The equivalent widths of λ1640 and λ4686 for individual WN stars are taken from Table 1 of Conti & Morris (1990). While the equivalent widths are distance and extinction independent (they are measured normalized to local continuum levels), they are affected by the possible presence of a binary star. Therefore, we ignore WN stars with listed companions. Additionally, we have revised the classification given in Conti & Morris, substituting the new WN subclasses from Smith et al. (1996) when available. Typical measurements show that regardless of WN subclass, the equivalent widths of λ1640 is ∼ 1/3 − 1/4 that of the λ4686 line. Because intrinsic line fluxes are distance and extinction dependent, we use the distance and foreground extinction corrected values for LMC WN stars, as presented in Table 2 of Conti & Morris (1990). Average values per subclass are presented in our Table 4. WC stars may also contribute to the He II λ1640 emission in WR galaxies. The emission – 14 – feature is strongest in the WC7 subtype for this class. However, NGC 3125-1 has strong N V λ1720 and N IV] λ1488 emission, and its C IV λ1550 emission is much weaker than the He II λ1640 emission line strength, suggesting that the contribution from WC stars is significantly lower than that for WN stars in this object.

4. THE MASSIVE STAR CONTENT OF LOCAL STARBURST GALAXIES

In this section, we combine massive star contributions derived from optical spectroscopy and available in the literature, with our estimates of the total number of WR and O stars from HST UV spectroscopy. We then assess whether the UV and optical results form a consistent picture of the massive star content of SSCs in local starbursts.

4.1. Complementary optical data

For our SSC sample we collect information on the WR and O star populations as measured from optical spectroscopy; these numbers are collected in columns 7−11 of Table 3. The primary references for these values are Schaerer, Contini, & Kunth (1999) and Guseva, Izotov, & Thuan (2000), with values for NGC 4214 provided from Sargent & Fillipenko (1991). In general, the slit widths used for optical spectroscopy are significantly larger than those used for our UV survey. For example, Schaerer et al. (1999) studied He 2- 10, NGC 3049, NGC 5253, Tol89, and NGC 3125 using a 1.600 wide slit and Guseva et al. (2000) studied NGC 1741 and NGC 3049 with a 200 wide slit, compared with the 0.200 slit typically used in the UV. In He 2-10 observations in the optical cover all 4 clusters studied in Chandar et al. (2003), as well as additional nearby clusters. We found from our detailed UV spectroscopic study of four He 2-10 clusters that these all formed coevally ∼ 4 − 5 Myr ago, and available optical colors suggest that the other clusters near the center of starburst A have similar ages. Thus, we also present results for He 2-10 where a 0.200 × 5.2500 area is extracted (called “He 2-10-all”). Due to the mismatch in aperture sizes, there are likely offsets in terms of the total flux levels between the two wavelength regimes. In §2.6 we estimated the amount of flux typically missed due to our slit width, and found for the nearest objects, this ranges from a factor of 4 to 9, whereas for more distant clusters, most of which are point-like, the UV fluxes are likely underestimated by a factor of 1.3 for point sources. The WR content from the optical spectra is usually determined from measurements of the He II λ4686 emission line, which can arise from WR stars of WN and/or WC subtypes. Note however, that many authors use the “WR bump”, which is the N IIIλ4640+ N III λ4650 – 15 – feature (which includes He II λ4686 in some cases). The He II λ4686 feature is often blended with nearby nebular emission lines (He, Fe, or Ar) and can show several broad stellar emission components (N III λ4640, C III λ4650), which can be difficult to separate in most low or medium-resolution spectra. Columns 8 and 9 in Table 3 give available He II λ4686 equivalent width and dereddened flux measurements from the literature for the SSCs studied in this work. We noted above that these measurements are based on observations taken through a significantly larger aperture than our UV data, which can result in some offset in the flux scale between the optical and UV. The extraction apertures used in the UV were optimized individually for each cluster, and chosen to retain the maximum amount of flux; thus these values in principle should at least asymptotically approach the total flux (although they actually represent lower limits).

4.2. Estimating the number of WR stars

In general, the number of WR stars is derived from the luminosity of a WR line; here we use the UV He II λ1640 line. Our dereddened He II λ1640 line fluxes are presented in column 3 of Table 3. Because we make no corrections for slit losses, our calculations represent a lower limit to the number of WR stars derived for each cluster. In principle, as long as we know the luminosity of one WR star in a specific line, we can derive the number of WR stars. We assume a mean WNL line luminosity of 1.6  1.5 × 1036 ergs s−1 for the λ4686 feature, as given in Schaerer et al. (1999) and Guseva et al. (2000). In a pure recombination spectrum (as expected in an optically thin nebula), the theoretical recombination value for the λ1640/λ4686 flux ratio is 6.5 − 7.5. Because the helium emission line spectra of WN stars are not necessarily optically thin, we adopt an intrinsic value of 7.6 for this ratio, as empirically derived by Conti & Morris (1990) for LMC WN stars (regardless of subclass), rather than using the theoretical value. The mean WNL λ1640 luminosity corresponding to the λ4686 value given above (assuming an intrinsic line ratio of 7.6), is 1.2  1.1 × 1037 ergs s−1 per “average” WNL star. Within the uncertainties, this value is consistent with the mean LMC WN4-WN6 star luminosity (based on individual stellar measurements from Conti & Morris 1990) of 1.4  1.2 × 1037 ergs s−1. Although there are large uncertainties in the WNL equivalent flux, resulting in systematic uncertainties on the total number of WN stars present, this does not affect the relative WN star numbers within our sample. Columns 4 and 10 of Table 3 present the number of WNL equivalents in our galaxy sample based on the UV measurements presented here, and on the λ4686 line flux published in Schaerer et al. (1999). For consistency between the optical and ultraviolet measurements we have adopted the distances of He 2-10, NGC 3049, NGC 3125, NGC 5253, and Tol89 – 16 – used in Schaerer et al. (1999). Also, the optical WR flux measurements are sensitive to the S/N of the observations.

4.3. Estimating the number of O stars

The number of O stars is calculated using different techniques in the UV and optical. Measurements from optical spectra are taken from the literature (Schaerer et al. 1999; Guseva et al. 2000). Here we briefly summarize the techniques used by these authors. The number of O stars were derived under the condition of case B recombination and assuming that all the ionizing photons emitted by the stars are absorbed by the gas. The total number of Lyman photons is derived from the observed luminosity of the Hβ emission line. To then estimate the number of O stars, the authors made a correction for the ionizing photon contribution from WR stars, and accounted for the age and IMF of the stellar population. The number of O stars derived in this manner are given in column 11 of Table 3 for galaxies with available observations in the literature. From our ultraviolet spectra, we estimate the number of O stars by scaling the best fit Starburst99 model to the dereddened cluster flux at λ1500. These values are given in column 5 of Table 3. Note also that the large number of WR stars derived for NGC 3125-1 means that these objects contribute significantly to the UV continuum, and in this case the number of estimated O stars from the UV observations must actually be the number of O + WR stars.

4.4. Comparison with optical spectroscopy results

In this section, we compare the estimated numbers of WR and O stars from our UV and available optical spectroscopy. This provides an important consistency check on the derived values. One question we wish to address is: could the strong He II λ1640 emission found in NGC 3125-1 have been predicted from optical results? Schaerer et al. (1999) fit multiple gaussians to measure individual emission lines, including He II λ4686, the feature typically used to estimate the contribution from WNL stars. Due to their limited spectral range however, they have used values from the literature for extinction corrections, based on the Balmer decrement. Guseva et al. (2000) adopt a different measurement procedure. They estimate the WR content by measuring the flux of the entire “blue bump”, after fitting and subtracting narrow nebular emission lines. One object presented here, NGC 3049, is common to both optical studies. Both works use comparable slit widths, although the exact – 17 – pointings cannot be assessed. The WN and WC number estimates for NGC 3049 between these two works is quite similar, suggesting that independent optical studies give consistent results. Not surprisingly, there is a larger discrepancy for the derived numbers of O and WR stars when comparing estimates made from optical vs. UV spectroscopy. Here, we discuss the results for a few galaxies individually. We note that the extremely large number of WN stars estimated by Schaerer et al. (1999) for He 2-10 starburst region A are significantly affected by their assumption of a very large extinction value. In fact from the slope of the

UV continuum, we found a much lower EB−V intrinsic to the starburst (0.24 vs. 0.56) for He 2-10 than that found by Vacca & Conti (1992). Given that C(Hβ) = 1.46 × EB−V , He 2- 10 shows the largest difference in reddening derived in previous works from the (nebular) Balmer decrement, when compared with our derivation based on the (stellar) continuum UV slope. Based on our NGC 3125-1 λ1640 emission line measurement, we estimate that ∼ 6100 WNL stars reside in this SSC (with no correction made for light which falls outside the slit), making it the most WR-rich known example of an individual SSC in the local universe. NGC 6764-1 and NGC 1741-1 also have large WN equivalent estimates; however there are large uncertainties associated with these measurements. The number of NGC 3125-1 WN stars estimated from the λ1640 line is roughly an order of magnitude larger than estimated from the λ4686 line. Assuming that all of the measured line flux in λ1640 and λ4686 comes only from WR stars (i.e. no nebular contamination), we are sensitive to the same stars in both lines, and we would naively expect to obtain a ratio of 1/3 − 1/4 (as measured for individual WR stars) for the λ1640/λ4868 equivalent width ratios. However, for NGC 3125-1 we find that the λ1640 EW is roughly equal to that measured for λ4686, despite the expected ratio of 1/3. In fact, none of the cluster equivalent width ratios approach the intrinsic value of 1/3 measured for individual WR stars. This is due to the dilution of the optical λ4686 equivalent width from an older underlying stellar population in these starburst galaxies (since equivalent width measurements are normalized to the local continuum), and makes the use of the λ4686 equivalent width a poor diagnostic measure of the true WR stellar content. Regardless, we note that NGC 3125-1 does have the largest He II λ4686 equivalent width measurement of all galaxies in the Schaerer et al. (1999) work; clues to the WR content in NGC 3125-1 were available from optical measurements. In the UV, where there is essentially no continuum contribution from an underlying stellar population, both the continuum and line fluxes are dominated entirely by the starburst itself, and we get a more robust estimate of the WR feature equivalent width. From the few objects which have both optical and UV based estimates of the massive – 18 – star content, we conclude that within a factor of a few, there tends to be reasonable agree- ment between the numbers of WR and O stars derived from the line luminosities in each wavelength regime. However, the difficulty of isolating the λ4686 line combined with signif- icant underlying continuum from lower mass stars in the optical, makes the measurements “cleaner” in the UV. We cannot draw any conclusion on whether the morphology of the He II λ1640 line can be accurately predicted from optical spectroscopy, since FWHM mea- surements of the λ4686 feature in composite stellar populations are not currently available in the literature.

4.5. Use of λ1640/λ4686 line fluxes as a reddening measurement

Conti & Morris (1990) empirically derived an intrinsic value of 7.6 for the λ1640/λ4686 flux ratio in individual LMC WR stars, very close to the predicted recombination value. In principle, this line ratio provides an extremely robust method to determine the extinction due to the ISM along the line of sight to a starburst galaxy. Because the λ1640 and λ4686 lines arise in the exact same stars (which are very short lived), if these features are present, then there is essentially no age dependence. By contrast, the commonly used optical Balmer lines give a measure of the nebular reddening rather than the stellar reddening. The HeII ratio should be independent of the underlying continuum, since it is the flux ratios which are used. The λ1640/λ4686 equivalent width ratio is larger in starburst galaxies when compared with individual WR stars, because there is a diluting underlying stellar continuum from lower mass stars, which should not affect the UV measurement, but does affect the optical (this stellar population is likely not from the starburst itself). However, the flux ratio λ1640/λ4686 should be a secure value, independent of the un- derlying stellar population, with an intrinsic value of 7.6 (Conti & Morris 1990). While we cannot precisely make this measurement due to mismatches in aperture size and pointing, here we simply check to make sure that the He II λ1640/λ4686 flux ratios are within the right ballpark for our SSCs. We expect that due to our significantly smaller aperture size in the UV, UV fluxes are underestimated relative to optical ones of the same objects, by a factor of ∼ 1.3 up to an order of magnitude, which should make the observed (dereddened) ratio smaller than 7.6. For all SSCs studied in this work, with the exception of NGC 3125-1, the flux ratio λ1640/λ4686 (after dereddening) is less than the intrinsic value of 7.6, as expected. For NGC 3125-1 however, even with no flux correction made to our 1640 A˚ flux, the ratio is already larger than 7.6. As mentioned in the previous section, this unusual cluster is also the site of the strongest known He II λ1640 stellar emission in a local starburst. In the next section, we take a closer look at this object. – 19 –

4.6. The unusual nature of NGC 3125-1

As noted above, the λ1640/λ4686 flux ratio in NGC 3125 already exceeds the expected recombination value of 7.6, prior to accounting for slit losses in the UV. In this respect, it is unique to our sample. One potential explanation for the strong He II λ1640 emission is the presence of a bright X-ray point source (J100633.9-295612) coincident (within the astrometric uncertainties) with NGC 3125-1 (from the on-line ROSAT HRI catalog). The high energy radiation field from an intermediate mass black hole theoretically could ionize the local ISM and thus produce the observed λ1640 emission. However, the line width of ∼ 1000 km s−1measured in NGC 3125-1 is much broader than expected for lines arising in the ISM, but entirely consistent with a stellar origin. In general, the observed UV/optical spectrum of this object is consistent with a stellar population dominated by WN stars, and doesn’t show any evidence for weak AGN-like activity. A second possibility for the λ1640 strength could be contamination from nebular emission. Indeed, Schaerer et al. (1999) see some weak, high-excitation Ar IV emission lines (λ4711, λ4740) in the optical spectrum, which is typical of nebular emission. They suggest that this may imply a nebular contribution to the He II λ4686 line. However, nebular contamination would affect both the λ1640 and λ4686 fluxes. Further, the UV spectrum of NGC 3125-1 clearly shows the presence of λ1488 and λ1720 emission lines, which arise in WN stars. A qualitative comparison of the FUV+NUV spectrum of NGC 3125-1 (shown in Figure 2) with that of individual WNL stars (from St. Louis 1990; PhD thesis) establishes that the relative strengths and widths of N IV λ1488, C IV λ1550 , He II λ1640 , and N V λ1720 are nearly identical between this cluster and individual late WN stars. Previously, we noted the presence of the λ1488 emission feature in the He 2-10 spectrum as well, although the λ1720 emission feature is not unambiguously detected in that galaxy (Chandar et al. 2003). In fact, NGC 3125-1 has the only detected N V λ1720 emission in our entire sample of SSCs. The galaxy NGC 3125 hosts another massive starburst cluster (NGC 3125-2), which unfortunately falls near the edge of the MAMA. However, this second cluster does have enough signal for age-dating. The low metallicity instantaneous Starburst99 model gives a best fit age of 3 Myr, identical to that of NGC 3125-1. Interestingly, in the Schaerer et al. (1999) paper, they estimate an even larger number of WN stars in NGC 3125-2 than NGC 3125-1, as well as a significantly larger number of WC stars, while still finding a similar number of massive O stars as NGC 3125-1. Potentially, this cluster has broad He II λ1640 emission, which is more difficult to observe at a given S/N. However, based on our UV spectroscopy, we conclude that the λ1640 flux from NGC 3125-2 is substantially lower than that for NGC 3125-1. Have we just happened to catch NGC 3125-1 at the correct age of 3 Myr, when the – 20 – number of WR stars (and particularly the number of WN stars) is at a maximum? There are other massive 3 Myr old SSCs in our sample (besides NGC 3125-2), and none of these approach the large number of WR stars detected in NGC3125-1. For example, NGC 1741-1 5 is a massive (> 10 M ), 3 Myr cluster which shows some evidence for He II λ1640 emission but with very large uncertainties (see Figure 1). Can we reconcile these observations and construct a view which accounts for the strange λ1640/λ4686 flux ratio of NGC 3125-1? The UV continuum slope is normally dominated by OB stars; however in this case the large number of WN stars will also make a significant contribution. One possible explanation is that we are seeing the WR stars through a “hole” where their energetic winds have blown out the natal cocoon earlier than have the OB stars. The relatively large extinction value derived from the UV slope is consistent with the gener- ally accepted scenario where very young clusters remain embedded in their natal material, until energetic stellar winds from evolving massive stars blow out the surrounding gas and dust. Clusters NGC 5253-5 and NGC 3310-1, which are the most similar to NGC 3125-1 in terms of age and He II λ1640 emission, are also young and heavily extincted. If WR stars are preferentially less extincted than OB stars in NGC 3125-1, then the

OB stars need to have a larger reddening than the EB−V = 0.52 given in Table 2, in order to compensate for the very low extinction of WR stars. Note that even if most of the UV continuum comes from WR stars seen through a hole in the intervening gas and dust, the intrinsic UV slope expected for such a population should still be close to −2.7. The scenario suggested above simultaneously explains the extinction value inferred from the UV continuum, while allowing for the flux ratio of the He II lines to approach the recombination value (since in effect the He II λ1640 flux would have been dereddened by too large a value for the WR stars). Although this “contrived” scenario is not very satisfying, we do not yet have a plausible alternative explanation. We conclude this section by emphasizing the empirical result of a very large equivalent width and flux measurement for stellar He II λ1640 emission in the massive young cluster NGC 3125-1, and suggest that this cluster is the most “extreme” WR cluster known in the local universe. The large number of inferred WR stars is the result of assuming that late-type WN stars dominate the spectrum; we note however, that the observed narrow line width gives firm constraints on the types of WN stars which contribute to the spectrum.

5. Comparison with z ∼ 3 Lyman Break Galaxy spectra

Recent high S/N composite spectra of hundreds of z ∼ 3 LBGs have shown relatively strong rest-frame He II λ1640 emission (Shapley et al. 2003). The width of the emission – 21 – feature (1500 km s−1) indicates that it is stellar in origin, likely from the strong winds of WR stars. The fact that the strength of this feature does not track the Lyα equivalent width in a systematic way (Shapley et al. 2003; Figure 9) supports a non-nebular origin for the He II λ1640 emission in the LBGs. While the He II λ1640 emission in z ∼ 3 LBGs is reproducible by a 3 Myr instantaneous burst model (for example from Starburst99), this is an unrealistic characterization of the aggregate properties of a large number of galax- ies, some of which are known individually to have a wide range of ages (Shapley et al. 2001a; Papovich, Dickinson, & Ferguson 2001). Currently available continuous star forma- tion models meanwhile, which should be much better at reproducing these composite galaxy spectra, cannot reproduce the strength of the He II λ1640 feature without an IMF slope which is significantly flatter than Salpeter. Models using these flat IMF slopes overproduce the observed C IV λ1550 P Cygni emission by a factor ≥ 4. Thus, current continuous star formation population synthesis models cannot simultaneously reproduce the C IV λ1550 and He II λ1640 stellar wind features in high redshift LBGs. Similarly, current synthesis models of instantaneous star formation cannot reproduce the strong He II λ1640 emission seen in NGC 3125-1. While there are well known deficiencies in the accounting of WR populations in these models, we note that we have established that NGC 3125-1 has unique properties when compared with other SSCs in local starbursts. Could these high redshift objects have formed large numbers of massive clusters with properties similar to NGC 3125-1? Although we do not yet understand what underlying mechanism is responsible for the very large observed WR population in this cluster, we use this empirical result to “interpret” the composite spectra of z ∼ 3 LBGs. To set limits on the types of stellar populations contributing at high redshift, below we describe an exercise aimed at quantifying the general rest-frame UV contribution of young SSCs in high redshift galaxies. We compare combinations of various empirical low redshift UV templates (described below), with composite LBG spectra from the Shapley et al. (2001b) study. We emphasize that this is a simple analysis, not meant to over-interpret the results, but to place limits on the possible contribution made to the integrated spectra of LBGs from clusters with properties similar to those found in local starbursts. The primary disadvantage of our study is that the small “slices” of starburst regions (covering ∼ tens to hundreds of pc) observed locally are not well matched to the kpc scales covered at higher redshifts. We alleviate this disparity by summing together large swaths of “field” from low metallicity sample galaxies. We consider the field to be the diffuse UV light seen in starbursts, and as studied via UV spectroscopy in NGC 5253 (Tremonti et al. 2001) and in He 2-10 (Chandar et al. 2003). We used four reasonably high (≥ 15) S/N templates as input:

• a high S/N template of a low metallicity “field” spectrum where we have summed up the – 22 –

diffuse UV light between luminous clusters in five (MKN33, NGC 4214, NGC 4449, NGC 4670, and NGC 5253) low metallicity starbursts. Typical abundances for the galaxies are ∼ 1/3 − 1/5 solar, which is a good match to abundances measured for high redshift galaxies (Pettini et al. 2000). This has a characteristic spectrum which is “older” than individual clusters (i.e. lacking the signature of massive stars), and is well represented by continuous star formation models.

• Super star cluster NGC 3125-1, representing 3 Myr “extreme” WR clusters. This object has the strongest known He II λ1640 stellar emission known in a local cluster.

5 • Cluster NGC 1741-1, which is a massive (> 10 M ) 3 Myr old cluster showing strong C IV λ1550 emission but weak Si IV λ1400. Unlike NGC 3125-1, which has a similar age and mass, the He II λ1640 emission feature is much less prominent in this cluster.

• Older high S/N individual SSCs with ages of 5, 7, and > 20 Myr. For example, we used He 2-10-1, NGC 4670-1, and NGC 5102-1.

In Figure 3 we show a few different linear combinations of our normalized template spectra compared with the Lyα “emission” z ∼ 3 LBG composite presented in Shapley et al. (2001b). They have made different composites based on whether Lyα is seen in emission, absorption, or has both components, in a given spectrum (see also their more re- cent composite spectra presented in Shapley et al. 2003, which qualitatively show the same He II λ1640 strength). The results differ only slightly if the “absorption” or “emission+absorption” Lyα composites are used, suggesting little difference in the massive cluster/WR content of the high redshift LBGs galaxies used to make these different spectra. The top panel shows that our normalized field spectrum, when compared with the LBG composite is a reason- ably good match to the continuum and weak absorption features, but cannot reproduce the strong He II λ1640 and C IV λ1550 emission features. The second panel shows a com- posite comprised of 90% field + 10% NGC 3125-1. While the C IV λ1550 emission line is still underpredicted, the He II λ1640 feature strength is well reproduced. The third panel shows 80% field + 20% NGC 3125-1, establishing that this fraction now overpredicts the He II λ1640 emission. The bottom panel shows a composite of 70% field + 15% NGC 3125-1 + 15% NGC 1741-1, which gives a reasonably good fit to the C IV λ1550 P Cygni profile and the He II λ1640 emission; however, the “blue edge” of the C IV λ1550 emission shows a larger and larger mismatch as one moves down the panels. This can be alleviated somewhat (although not completely) if rather than NGC 1741-1 (a 3 Myr cluster), we use a SSC which is 7 or > 20 Myr old (NGC 4670-1 or NGC 5102-1), although this combination makes it somewhat more difficult to match the stellar wind features. While the overall abundances for LBGs at high redshift are comparable to that in local starbursts, one could speculate – 23 – that this discrepancy in the C IV blue edge is related to differing O/C ratios. Carbon is produced primarily in low and intermediate mass stars, while oxygen comes primarily from high mass stars (e.g., Wheeler, Sneden, & Truran 1989). The longer timescale for carbon enrichment could then potentially result in a lower C/O abundance in the early universe relative to that observed locally. Based on the comparisons in Figure 3, we conclude that “extreme” WR clusters similar in properties to NGC 3125-1 cannot contribute more than ∼ 15% of the total UV flux in LBGs. This upper limit is of course substantially larger than found in the local universe, where this object thus far, appears to be unique in terms of its WR content. The implied 15% flux is similar to the total UV light contributed by all observed SSCs in local starbursts (e.g., Meurer et al. 1995). Of course, there are a number of other combinations/templates which might reproduce the observed LBG spectral features. If instead of NGC 3125-1, we use local 3 Myr old SSCs showing narrow but weak He II λ1640 emission (e.g., NGC 3310-1), such objects would have to comprise between 40-70% of the total UV flux in z ∼ 3 LBGs in order to match the P Cygni C IV λ1550 profile and He II λ1640 emission features. However, a larger contribution from SSCs implies a smaller overall contribution from the field. As the field contribution is lowered below ∼ 60 − 70%, the mismatch in the “blue edge” of C IV λ1550 increases, even if older SSCs are included.

6. SUMMARY AND CONCLUSIONS

We have compiled STIS FUV long-slit spectra of individual SSCs from 18 local starburst galaxies, and estimated the massive (WR and O) star content. The He II λ1640 line in par- ticular provides important information on the most massive stars in these clusters. Cluster NGC 3125-1 is found to have the strongest known He II λ1640 emission of stellar origin. We compile mean properties of individual Galactic and LMC WN stars for comparison. The narrow morphology of NGC 3125-1, when compared with individual WN stars, indicates that WR stars of late WN subtype dominate this spectrum, and we estimate that ∼ 6100 WNL-equivalent stars reside in this single cluster, roughly an order of magnitude larger than predicted by optical observations. Based on the measured λ1640/λ4686 equivalent widths and flux ratios, we find that NGC 3125-1 is atypical in our sample, with a measured (dereddened) flux ratio larger than the recombination value. A scenario where the WR stars have preferentially blown out their natal gas and dust cocoons relative to the OB star population can account for the measurements. – 24 –

We use the observations of local SSCs combined with an empirical stellar field template to constrain the stellar populations in Lyman Break Galaxies at z ∼ 3. The comparison suggests that if we just summed up data from local low metallicity starbursts, we can- not reproduce the z ∼ 3 LBG composite. However, if clusters with properties similar to NGC 3125-1 were much more common in the early universe than now, we are able to repro- duce the stellar emission features in the LBG composite. This suggests that “extreme” WR clusters similar to NGC 3125-1 (which is unique in our sample) were much more common in the early universe, contributing up to 15% of the total UV flux. This conclusion is strength- ened when other massive, 3 Myr old local SSCs showing weaker He II λ1640 emission are used as templates, since these would have to comprise 40 − 70% of the total UV light bud- get, which is extremely unlikely. We note that the “blue edge” of C IV λ1550 absorption feature is not well matched by linear combinations including SSCs with ages between 5 − 8 Gyr at a fraction higher than ∼ 70%. This mismatch can be explained by different element abundance ratios at high redshift and at z = 0.

We thank Sandra Savaglio for her help on this paper, and an anonymous referee for his/her careful reading which significantly improved the presentation of this work. We are grateful for support from NASA through grant GO-09036.01-A from the Space Telescope Science Institute, which is operated by the AURA, Inc., for NASA under contract NAS5- 26555. – 25 –

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This preprint was prepared with the AAS LATEX macros v5.0. – 28 –

Table 1. Properties of Sample Galaxies

a b Target Distance logO/H + 12 EB−V Size Projected Ref. Galaxy (Mpc) in Slit (pc)

He 2-10 9 8.9 0.111 1100 × 9 1,12 MKN33 22 8.4 0.120 3150 × 25 2,13 MKN36 8 7.9 0.030 900 × 7 3,14 MKN209 4.9 7.9 0.015 600 × 5 4,15 NGC 1741 57 8.2 0.051 6200 × 49 5,16 NGC 3049 18.3 9.0 0.038 2200 × 44 1,1 NGC 3125 11.5 8.3 0.076 1400 × 11 1,6 NGC 3310 18 8.2 0.022 2300 × 18 1,17 NGC 4214 2.9 8.2 0.022 350 × 3 7,12 NGC 4449 3.9 7.8 0.019 550 × 4 8,8 NGC 4670 16 8.2 0.015 1900 × 16 5,6 NGC 5102 3.1 9.0 0.055 400 × 3 9,18 NGC 5253 4.1 8.2 0.066 500 × 2 10,19 NGC 5996 47 8.9 0.034 6200 × 49 6,6 NGC 6764 32 8.7 0.067 4800 × 39 11,20 NGC 7552 21 9.2 0.014 2500 × 20 5,6 Tol1924-416 37 8.1 0.087 5000 × 40 6,6 Tol89 14.7 8.3 0.066 1800 × 36 1,1

References. — (1) Schaerer et al. 1999 (2) Kennicutt et al. 2003 (3) Conti 1991 (4) Viallefond & Thuan 1983 (5) Tremonti et al. 2003 (6) Heckman et al. 1998 (7) Maiz-Apellaniz, Cieza, & MacKenty 2002 (8) Boeker et al. 2001 (9) McMillan, Ciardullo, & Jacoby 1994 (10) Saha et al. 1995 (11) Schinnerer, Eckart, & Boller 2000 (12) Kobulnicky, Kennicutt, & Pizagno 1999 (13) Davidge 1989 (14) Garnett 1990 (15) Hunter & Hoffman 1999 (16) Guseva, Izotov, & Thuan (17) Pastoriza et al. 1993 (18) D. Calzetti, private communication (19) Izotov & Thuan 1999 (20) Contini et al. 1997 aForeground, Milky Way extinction from the Schlegel et al. (1998) maps bBased on Distance given in Column 2, assumed slit length projected onto (2500) STIS MAMAs and given the slit width of the observations – 29 –

Table 2. Basic Properties of Super Star Cluster Sample

a a b c d Cluster RA DEC # pixels βi Age E(B − V ) (J2000) (J2000) (Myr) (mag)

He 2-10-1e 08:36:15.1 -26:24:33.7 15 -0.90 5  1 0.25 He 2-10-alle 08:36:15.1 -26:24:33.7 48 -1.01 5  1 0.23 MKN 33-all 10:32:31.9 +54:24:02.2 66 -1.84 5  1 0.09 MKN 36-1 11:04:58.4 +29:08:16.2 19 -1.83 < 1f 0.10 MKN 209-1 12:26:15.9 +48:29:38.8 31 -2.45 < 1f 0.00 NGC 1741-1 05:01:37.6 -04:15:32.6 15 -0.54 3  1 0.34 NGC 1741-all 05:01:37.6 -04:15:32.6 54 -1.35 4  1 0.19 NGC 3049-1 09:54:49.4 +09:16:15.9 25 -1.18 3  1 0.18 NGC 3125-1g 10:06:33.3 -29:56:07.8 15 0.62 3  1 0.52 NGC 3310-1 10:38:44.9 +53:30:04.7 11 -0.19 5  1 0.35 NGC 3310-all 10:38:44.9 +53:30:04.7 117 -0.98 6  1 0.20 NGC 4214-1 12:15:39.5 +36:19:34.8 23 -1.20 4  1 0.21 NGC 4449-all 12:28:09.7 +44:05:15.9 83 -2.74 5  1 0.00 NGC 4670-1 12:45:17.4 +27:07:31.8 53 -1.76 7  1 0.12 NGC 5102-1 13:21:57.5 -36:37:49.5 13 1.50 55−37 0.44 NGC 5253-5h 13:39:56.0 -31:38:25.0 27 -0.01 2  1 0.42 NGC 5996-1 15:46:58.9 +17:53:04.2 21 -1.29 4  1 0.20 NGC 6764-1 19:08:16.4 +50:55:59.3 17 0.70 3  1 0.48 NGC 7552-1 23:16:10.8 -42:35:02.7 15 -0.71 5  1 0.28 Tol 1924-416-1 19:27:58.3 -41:34:29.8 41 -2.67 < 1f 0.00 Tol 89-1 14:01:19:9 -33:04:10.7 81 -2.01 4  1 0.08

Note. — Nomenclature for “all” observed clusters within a galaxy is described in section 2.1 of the text. The naming convention for the remaining clusters matches that from our larger study (Chandar et al. 2004, in prep.). aCluster coordinates are measured from available HST WFPC2 and FOC images. If only a co-added cluster spectrum for a particular galaxy is presented, we give the specified central pointing of the slit bNumber of pixels summed up in the spatial direction for each cluster spectrum cPower law index of the UV continuum (F ∝ λβ) over the wavelength range 1240−1600A˚ af- ter correcting for foreground Milky Way reddening. β is measured to an accuracy of 0.1. dExtinction internal to the starburst. This value has been derived from the UV slope (as described in §2.3), after correcting for the foreground EB−V value given in column 4 of table 1, and assuming an intrinsic slope given by the best fit age. We estimate that the intrinsic UV slope is known to 0.1, due to uncertainties in the age. This translates into an uncertainty of 0.02 in EB−V . eHe 2-10-1 refers to cluster 1 presented in Chandar et al. (2003). He 2-10-all corresponds to a slice through starburst region A fThe χ2 vs. age plot shows a strong degeneracy between 0 Myr and 6 Myr. gNGC 3125-1 presented here corresponds to cluster NGC 3125-A presented in Schaerer et al. (1999). hNGC 5253-5 is cluster 5 from the Calzetti et al. (1997) and Tremonti et al. (2001) works. – 33 –

Table 3. Massive Star Content in Local Starburst Super Star Clusters

Object HeII 1640 HeII 4686 a b EW Luminosity # WNL # O FWHM EB−V EW Luminosity WNL # O (A)˚ (erg s−1) (A)˚ (mag) (A)˚ (erg s−1)

He 2-10-1 2.8  0.8 2.6E+39 213  61 1270 9.5  6.1 2.3 1.9E39 ...... He 2-10-all 3.0  0.9 6.6E+39 548  164 3033 11.1  2.3 0.56 2.3 1.9E39 1200 2450 − 4900 MKN 33-all 3.5  2.9 4.9E+39 407  337 1643 · · · ...... MKN 36-1 <1.6 <9.5E+37 <8 78 · · · ...... MKN 209-1 1.4  0.6 3.6E+36 0 3 11.2  3.8 ...... NGC 1741-1 4.0  3.3 6.3E+40 5252  4333 11561 · · · ...... NGC 1741-all 3.4  2.5 7.0E+40 5794  4260 19697 · · · 0.28 ...... 2000 4E5 NGC 3049-1 2.0  0.2 3.4E+39 283  28 1444 16.9  9.6 0.16 5.2 8.0E38 500 ... NGC 3125-1 7.4  2.2 6.2E+40 5184  1540c 6144 5.9  0.4 0.27 7.3 7.9E38 500 3240 − 6470 NGC 3310-1 2.8  2.5 1.3E+39 111  99 524 · · · ...... NGC 3310-all <2.2 <2.2E+39 <181 1128 · · · ...... NGC 4214-1 3.2  1.3 9.1E+38 76  31 278 13.4  10.0 ...... NGC 4449-all <2.6 <1.5E+37 <1 6 · · · ...... NGC 4670-1 <1.0 <1.3E+39 <105 1144 · · · ...... NGC 5102-1 <0.4 <9.8E+37 <8 0 · · · ...... NGC 5253-5 2.9  1.4 4.6E+38 38  18 155 4.8  0.2 0.30 3.0 5.0E37 30 960 − 1080 NGC 5996-1 2.9  0.4 3.9E+40 3285  453 14773 19.5  16.8 ...... NGC 6764-1 2.1  1.5 1.1E+41 9100  6500 44307 · · · ...... NGC 7552-1 2.5  2.1 6.1E+39 506  425 3397 · · · ...... Tol 1924-416-1 1.4  0.8 4.4E+39 363  207 3894 · · · ...... Tol 89-1 3.5  2.5 1.1E+39 95  68 317 · · · ......

Note. — All massive star properties and extinction derived in the optical comes from Schaerer et al. (1999), except for NGC 1741, where the values have been taken from Guseva et al. (2000). We use a mean WN flux of 1.2 × 1037 ergs s−1for the λ1640 line. For the λ4686 feature, Schaerer et al. (1999) and Guseva et al. (2000) used a flux of 1.6 × 1036 ergs s−1. aUncertainties on the number of WNL-equivalent stars is determined from the fractional uncertainty on the measured λ1640 EW.

b CHβ EB−V value determined as EB−V = 1.46 cDue to the large WR population, this number refers to the O + W R contribution. – 31 –

Table 4. Mean properties of WN stars

WN HeII 1640 HeII 4686 Subclass EW Luminosity FWHMa EW Luminosity FWHM (A)˚ (erg s−1) (A)˚ (A)˚ (erg s−1) (A)˚

WN2 ...... 23 ...... 63b WN3 11050 2.5E36b 11.5 − 14.5 30040 3.6E35b 38  10 WN4 10060 6.93.6E36 10.0 − 14.8 215100 7.53.0E35 30  5 WN5 143 1.2E37b ... 4030 2.8E36b 23  3 WN6 202 3.7E37b 8.9 − 10.0 808 7.8E36b 20  6 WN7 17b 2.1E36b 9.5 55b 2.9E35b 17  6 WN8 ...... 7.0 ...... 12  4 WN9 ...... 6  2

aλ1640 FWHM are given for 1 − 3 stars in each subtype. We give the range of measured values (a single value represents the measurement for one star of the given subclass). These are uncorrected for instrumental profile binformation available for three or less stars, so numbers may be unreliable. The value given is the mean value for the available data. – 32 –

2.0 He 2−10−1 1.5 Solar

1.0

0.5

0.0

2.0 He 2−10−all 1.5 Solar

1.0

0.5

0.0

2.0 MKN 209−1 1.5 LMC Rectified Flux 1.0

0.5

0.0

2.0 MKN 33−all 1.5 LMC

1.0

0.5

0.0

1200 1300 1400 1500 1600 1700 Wavelength (Å)

Fig. 1.— Far UV spectra of our sample clusters are presented with the best fit “He II-free” Starburst99 model (described in the text) overplotted. The data and models have been smoothed by three pixels for clarity. Each panel shows the cluster spectrum (thick black line), the best fitting model spectrum (gray line), their residual (thin black line) and the formal errors (dotted lines). The rest wavelength of the He II 1640 line is marked with a dashed line. Note that for galaxies where the LMC/SMC metallicity Starburst99 model was used to determine the age and extinction, we show the solar model (at the same age) redward of λ1600. – 32 –

2.0 MKN 36−1 1.5 LMC

1.0

0.5

0.0

2.0 NGC 1741−1 1.5 LMC

1.0

0.5

0.0

2.0 NGC 1741−all 1.5 LMC Rectified Flux 1.0

0.5

0.0

2.0 NGC 3049−1 1.5 Solar

1.0

0.5

0.0

1200 1300 1400 1500 1600 1700 Wavelength (Å)

Fig. 1.— continued – 32 –

2.0 NGC 3125−1 1.5 LMC

1.0

0.5

0.0

2.0 NGC 3310−1 1.5 LMC

1.0

0.5

0.0

2.0 NGC 3310−all 1.5 LMC Rectified Flux 1.0

0.5

0.0

2.0 NGC 4214−1 1.5 LMC

1.0

0.5

0.0

1200 1300 1400 1500 1600 1700 Wavelength (Å)

Fig. 1.— continued – 32 –

2.0 NGC 4449−all 1.5 LMC

1.0

0.5

0.0

2.0 NGC 4670−1 1.5 LMC

1.0

0.5

0.0

2.0 NGC 5102−1 1.5 Solar Rectified Flux 1.0

0.5

0.0

2.0 NGC 5253−5 1.5 LMC

1.0

0.5

0.0

1200 1300 1400 1500 1600 1700 Wavelength (Å)

Fig. 1.— continued – 32 –

2.0 NGC 5996−1 1.5 Solar

1.0

0.5

0.0

2.0 NGC 6764−1 1.5 Solar

1.0

0.5

0.0

2.0 NGC 7552−1 1.5 Solar Rectified Flux 1.0

0.5

0.0

2.0 Tol 1924−416−1 1.5 LMC

1.0

0.5

0.0

1200 1300 1400 1500 1600 1700 Wavelength (Å)

Fig. 1.— continued – 32 –

2.0 Tol 89−1 1.5 LMC

1.0

0.5

0.0 Rectified Flux

1200 1300 1400 1500 1600 1700 Wavelength (Å)

Fig. 1.— continued – 33 –

2 NGC3125−FUV

1 N IV He II

1200 1300 1400 1500 1600 1700 2 NGC3125−NUV

1

He II N IV 1600 1800 2000 2200 2400 2600 2800 Wavelength (Å) Rectified Flux Fig. 2.— Rectified STIS FUV and NUV spectra of NGC3125-1. Note the stellar wind lines at λ1488, 1640, and 1720, which indicate the presence of WR stars. – 34 –

1.0 0.8

0.6 He II

0.4 Si IV 100% Field C IV

1.0 0.8 0.6 0.4 90% Field + 10% NGC 3125−1

1.0 0.8 0.6 0.4 80% Field + 20% NGC 3125−1

1.0 0.8 Rectified Flux 0.6 0.4 70% Field + 15% NGC 3125−1 + 15% NGC 1741−1 1300 1400 1500 1600 1700 Wavelength (Å)

Fig. 3.— This figure compares the Lyman alpha Absorption+Emission spectrum from Shap- ley et al. (2001b) with empirical rest-frame UV templates created from local starbursts. All panels show the rectified high redshift LBG spectrum (gray) with that from different com- binations of our templates (black). Top Panel: 100% field spectrum; Second Panel: 90% “field” + 10% NGC 3125-1 contribution; Third Panel: 80% field + 20% NGC 3125-1 contri- bution; Bottom Panel: 70% field + 15% NGC 3125-1 + 15% fraction of cluster NGC 1741-1.