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The R136 star cluster dissected with Hubble Space Telescope/STIS. I. Far-ultraviolet spectroscopic census and the origin of HeII 1640 in young star clusters

We introduce a HST/STIS stellar census of R136a, the central ionizing star cluster of 30 Doradus. We present low resolution far-ultraviolet STIS/MAMA spectroscopy of R136 using 17 contiguous 52x0.2 arcsec slits which together provide complete coverage of the central 0.85 parsec (3.4 arcsec). We provide spectral types of 90% of the 57 sources brighter than m_F555W = 16.0 mag within a radius of 0.5 parsec of R136a1, plus 8 additional nearby sources including R136b (O4\,If/WN8). We measure wind velocities for 52 early-type stars from CIV 1548-51, including 16 O2-3 stars. For the first time we spectroscopically classify all Weigelt&Baier members of R136a, which comprise three WN5 stars (a1-a3), two O supergiants (a5-a6) and three early O dwarfs (a4, a7, a8). A complete Hertzsprung-Russell diagram for the most massive O stars in R136 is provided, from which we obtain a cluster age of 1.5+0.3_-0.7 Myr. In addition, we discuss the integrated ultraviolet spectrum of R136, and highlight the central role played by the most luminous stars in producing the prominent HeII 1640 emission line. This emission is totally dominated by very massive stars with initial masses above ~100 Msun. The presence of strong HeII 1640 emission in the integrated light of very young star clusters (e.g A1 in NGC 3125) favours an initial mass function extending well beyond a conventional upper limit of 100 Msun. We include montages of ultraviolet spectroscopy for LMC O stars in the Appendix. Future studies in this series will focus on optical STIS/CCD medium resolution observations.


INTRODUCTION
The formation of massive stars remains an unsolved problem in astrophysics. This is especially challenging since very massive stars (VMS, ≥ 100M⊙) are usually located within young massive star clusters (≥ 10 4 M⊙) with small core radii (rc ∼ 0.1-0.2 pc, Portegies Zwart et al. 2010). The short lifetime of such stars (2-3 Myr, Yusof et al. 2013) restricts suitable Milky Way star clusters to only a handful, including the Arches (Martins et al. 2008) and NGC 3603 (Schnurr et al. 2008), both of which are highly reddened. Therefore, observations of young star clusters at ultraviolet wavelengths -corresponding to the peak output of hot, luminous starsare largely restricted to extragalactic cases.
R136 was included in the census of optically bright Magellanic Cloud stars by Feast et al. (1960), and subsequently resolved into three components, a, b and c (Feitzinger et al. 1980), with component a subsequently shown to be multiple (Weigelt & Baier 1985). R136a is the closest example of an extragalactic young massive star cluster (Hunter et al. 1995). It is the central cluster of the NGC 2070 complex within the 30 Doradus star forming region in the Large Magellanic Cloud (LMC, Massey & Hunter 1998;Doran et al. 2013), and is relatively lightly reddened, facilitating study at ultraviolet wavelengths (de Marchi et al. 1993). However, the 0.1 pc core radius of R136a (Portegies Zwart et al. 2010;Selman & Melnick 2013) subtends only 0.41 ′′ at the 50 kpc distance to the LMC (Pietrzynski et al. 2013), which led to the cluster at one time being mistaken for a ≥1000 M⊙ star (Cassinelli et al. 1981, Feitzinger et al. 1983, although Walborn (1973) had earlier predicted that R136a would have a structure analogous to the core of NGC 3603. Consequently, observations at high spatial resolution are required to resolve individual stars within the core of R136a, requiring either Hubble Space Telescope or large ground-based telescopes with adaptive optics (VLT/MAD: Campbell et al. 2010).
The star-formation history of NGC 2070 has recently been investigated by Cignoni et al. (2015) using Hubble Tarantula Treasury Project (HTTP) imaging, revealing an upturn in star formation ∼7 Myr ago, peaking 2-3 Myr ago, and a total stellar mass of 9×10 4 M⊙ assembled within 20 pc of R136. The mass function for intermediate-and lowmass stars in R136 have been obtained from HST imaging (Hunter et al. 1995;de Marchi et al. 2011). However, the colour/reddening degeneracy of hot massive stars at ultraviolet and optical wavelenths necessitates spectroscopy for robust physical properties, and in turn, masses, measuring present-day mass functions (though see Maíz Apellániz et al. 2014). To date, Koornneef & Mathis (1981) and Vacca et al. (1995) have presented UV scans across NGC 2070 with the International Ultraviolet Explorer (IUE), while UV spectroscopy of individual stars of R136a has been limited to the brightest members with HST/GHRS (de Koter et al. 1997(de Koter et al. , 1998. Meanwhile, Massey & Hunter (1998) have obtained optical spectroscopy for dozens of stars within NGC 2070 using HST/FOS, which revealed large numbers of very early O stars, indicating a young cluster age (<1-2 Myr). More recently Schnurr et al. (2009) used VLT/SINFONI to obtain spatially resolved K-band spectroscopy of the brightest sources, namely R136a1, a2, a3 and c, which are WN-type emission-line stars. Analysis of the VLT and HST spectroscopy plus VLT/MAD imaging led Crowther et al. (2010) to conclude that such stars possess very high (initial) stellar masses of 160-320 M⊙, results largely supported by Bestenlehner et al. (2011) and Hainich et al. (2014).
The existence of such high-mass stars in R136 challenges the previously accepted ∼ 150M⊙ upper stellar mass limit (Figer 2005), unless they are products of stellar mergers (Banerjee et al. 2012;Schneider et al. 2014a). If such stars are indeed merger remnants one would expect younger apparent ages for this with respect to lower mass stars, and differences in the present day mass function with respect to a Salpeter slope (Schneider et al. 2014a;Schneider et al. 2015). Massey & Hunter (1998) suggested that the mass function of R136 is normal, although their study only included a subset of the brightest stars in the vicinity of R136a.
A census of the hot, luminous stars in R136a with STIS is the focus of the present series of papers. We discuss the suitability of ultraviolet spectroscopy for spectral classification of O stars, based in part upon spectroscopic datasets in the usual blue visual wavelength range. Wind velocities for O stars can uniquely be obtained from P Cygni profiles in the far-UV (e.g. , Prinja et al. 1990). Comparisons of empirical (reduced) wind momenta 1 with theoretical predictions usually rely on approximate v∞ calibrations (e.g. Lamers et al. 1995). Indeed, high quality far-UV spectroscopy has been restricted to relatively few LMC O stars to date (Prinja & Crowther 1998;Walborn et al. 1995;Massey et al. 2005). One byproduct of the present study is a calibration of wind velocities for (mostly) early O dwarfs at the ∼ 0.5Z⊙ metallicity of the LMC.
In the first study in this series, we introduce our new STIS observations of R136a in Sect. 2. We present and analyse ultraviolet STIS/MAMA spectroscopy of individual stars in Sect. 3, including approximate spectral types and estimates of terminal wind velocities. An ultraviolet spectral atlas of LMC O stars is compiled using literature spectra, in a few instances complemented with new STIS/CCD spectroscopic datasets to provide templates of missing spectral types. The current dataset also allows the contributions from individual stars to its integrated ultraviolet spectrum to be quantified. This is discussed in Sect. 4, of relevance for other young massive star clusters in the local universe (e.g. Chandar et al. 2004;Leitherer et al. 2011) and the integrated rest-frame ultraviolet spectrum of high redshift star forming galaxies (e.g. Shapley et al. 2003). Finally, brief conclusions are drawn in Sect. 5.

OBSERVATIONS
Here we present our new long-slit HST Space Telescope Imaging Spectrograph (STIS) observations of R136, and use of archival WFC3/UVIS imaging discussed by de Marchi et al. (2011).  calstis slit loss corrections were applied. Therefore, absolute wavelength scales may be incorrect by up to ±2 pix (0.56Å/pix) while absolute far-UV fluxes will be underestimated by up to a factor of two in extreme cases.
Extractions were also independently made using multispec (Maíz Apellániz 2005; Knigge et al. 2008) which is particularly suited to the crowded region at the geometric centre of the cluster. Figure 2 shows two dimensional spectral datasets for particularly crowded slits NW1 (a1, a2 and a5) and NW3 (a4 and a8), including collapsed spectral images between λλ1475-1525. multispec extractions relied upon empirical point spread functions for G140L/MAMA spectroscopy employing the position of individual sources in slits from short exposure HST WFC3/UVIS F336W imaging of R136 from Oct 2009 (for details see de Marchi et al. 2011). The latter provided wavelength corrections if sources were not centred in individual slits, but a standard slit loss correction is again adopted.
For relatively isolated sources, calstis and multispec extractions agree extremely well, while the latter approach usually performed better in situations where sources were closely spaced within individual slits. For consistency, spectroscopic datasets presented here were extracted with multispec, except for R136 H70, 141, 149 for which multispec extractions were problematic.

Optical spectroscopy
We have obtained optical STIS spectroscopy using the Charged Couple Device (CCD) detector and identical 52 ′′ × Table 2. HST/WFC3 photometry (de Marchi et al. 2011) for sources (m F555W ≤ 16.0 mag). within a projected distance of 2.05 arcsec (0.5 parsec) from R136a1. In case WFC3 photometry is unavailable, WFPC2 photometry from Hunter et al. (1995) is used, albeit offset in F555W by -0.17 mag (indicated in italic font, see Sect. 2.3). Identifications are from Weigelt & Baier (1985, WB85) or Hunter et al. (1995, HSH95), while spectral types are obtained from optical spectroscopy. A F555W is either obtained via the observed m F336Wm F555W colour plus R 5495 = 4.1, or A F555W = 1.72 mag is adopted (values for WN5 stars are from Crowther et al. 2010).  0. ′′ 2 slit at identical positions and position angle (or PA + 180 • ) to the UV spectroscopy, although the coarser plate scale of 0.05 arcsec/pix provided complete spatial coverage of the slit length. In total, two exposures at four grating positions with the G430M grating provided complete blue coverage λλ3793-4849Å, with a pixel scale of 0.28Å/pixel, and point source resolving power of 7700 at λ=4400Å. A single grating position with the G750M grating provided λλ6482-7054Å spectroscopy, with a pixel scale of 0.56Å/pixel, and point source resolving power of 6000 at Hα. For all CCD observations, the second exposure was offset spatially along the slit by 7 pix (0.35 arcec). Three of the four G430M observations were taken over the course of a few days, with the fourth (G430M/4194) obtained approximately 6 months later to enable to search for low period binaries. All G750M exposures were taken at a single epoch. In total CCD exposures were split between 7 visits during Cycle 19 (GO 12465, 26 orbits), plus 2 visits during Cycle 20 (GO 13052, 7 orbits), also set out in Table 1. multispec and calstis extractions of CCD observations will be the focus of the next paper in this series (Caballero-Nieves et al. 2016, in prep).

Stellar census of R136
In sources for which de Marchi et al. (2011) photometry is unavailable, we have substituted Hunter et al. (1995) results to which these offsets have been applied. Overall, agreement between early WFPC2 and new WFC3 photometry is generally very good, with a few notable exceptions -WFC3 photometry from de Marchi et al. (2011) is significantly brighter for R136a1, a2 and a6.
In Table 2 we provide a catalogue of all 55 sources within 2.05 arcsec (0.5 pc) of R136a1 brighter than mF555W = 16.0 mag, while a further 8 bright sources observed with STIS/MAMA are presented in Table 3. Additional sources brighter than mF555W = 17.0 mag within 4.1 arcsec (1 pc) are listed in Appendix D (Table D1). There is evidence for a non-standard extinction law in 30 Doradus, and significant variation thereof (Maíz Apellániz et al. 2014). Specifically for R136, Doran et al. (2013)  For stars without UV derived spectral types, we adopt O6 V (T eff = 39,900 K) for stars brighter than mF555W = 16.0 mag, and O9.5 V (T eff = 32,900 K) for fainter stars. We adopt AF555W = 1.72 mag for all sources without WFC3 photometry.
Severe crowding results in significant differences in photometry between Hunter et al. (1995)   HST/WFC3 photometry implies F336W/F555W/F814W flux ratios of 0.88/0.95/0.84 for a2/a1, while K-band VLT/SINFONI observations favour a significantly lower ratio of 0.76 . From our STIS/MAMA spectroscopy, the 1500Å flux ratio inferred from the multispec extraction is lower still (0.56), while Gaussian fits to spatial profiles in our 2D ultraviolet NW1 dataset (lower left panel in Fig. 2) favour 0.62. We adopt this flux ratio, so m a1 F555W − m a2 F555W = 0.52 under the reasonable assumption that their interstellar extinctions are equivalent. From comparison between the UV fluxes of a1 and a2 and relatively isolated sources in the NW1 slit, we support WFC3/F555W photometry of a1, resulting in a revised a2 apparent magnitude of mF555W = 12.80 mag in Table 2. This leads to a downward revision of 0.18 dex in stellar luminosity, and corresponding 20% lower mass estimate.
In addition, WFC3 photometry of a6 from de Marchi et al. (2011) is significantly brighter than Hunter et al. (1995). From our MAMA SE1 dataset, we obtain a 1500Å flux ratio of 0.70 for a6/a3, in perfect agreement with m a3 F555W − m a6 F555W = -0.38 from WFC3, so adopt photometry from de Marchi et al. Finally, de Marchi et al. (2011) provide a single photometric measurement for a4+a8, which agrees satisfactorily with the combination of these sources in Hunter et al. (1995). Again, we utilise UV datasets (lower right panel in Fig. 2) to verify their relative fluxes. From our multispec extraction we obtain a 1500Å flux ratio of 0.66 for a8/a4, in good agreement with 0.65 from Gaussian fits to their spatial profiles in the NW3 slit. We obtain m a4 F555W − m a8 F555W = 0.46, again assuming identical interstellar extinctions for these stars. From the combined WFC3 photometry of m a4+a8 F555W = 13.41±0.01, we obtain revised a4 and a8 apparent magnitudes of mF555W = 13.91 and 14.42 mag (Table 2).

SPATIALLY RESOLVED FAR-ULTRAVIOLET SPECTROSCOPY
We have classified individual R136 stars, based upon a comparison between their ultraviolet morphology and LMC template stars with reliable optical classifications. As noted above, we have incorporated 10 reference O2-O9.5 stars based on optical classifications provided in Caballero-Nieves et al. (2016, in prep) which are indicated in Table 2-3. Our approach therefore complements the UV clasification system of LMC OB stars by Smith Neubig & Bruhweiler (1999) using IUE SWP/LORES spectroscopy. We adopt existing (optical) subtypes for the three WN5 stars in R136a (Crowther & Dessart 1998).
In addition to the point sources, a significant fraction of the far-UV spectrum of the central parsec of R136 arises from the intra-cluster background, on the basis of the multispec fits to the 100 brightest cluster members. This is illustrated in Fig. 4, and represents a combination of unresolved cluster members plus extended emission from bright cluster stars originating in adjacent slits. In Fig. 5, we present the far-UV spectra from the sum of selected background regions of SE3 (1 ′′ E of R136a1) and SE4 (0.7 ′′ SE of R136a1), reminiscent of R136 H70 (O5 Vz), except for extended He ii emission originating from either R136a1 or a2. ∼20% of the integrated far-UV continuum arises from the diffuse back- Table 4. STIS/MAMA-derived properties of R136 stars within a projected distance of 2.05 arcsec (0.5 pc), brighter than m F555W = 16.0 mag), sorted by projected distance from R136a1. References to optical (literature) spectral types are provided in Table 2. F 1500 fluxes are averages for the range 1500±25Å (estimated values are provided in parentheses). Aside from the WN5 stars , temperatures and luminosities are from spectral type calibrations (Doran et al. 2013), while inferred current masses M and ages τ for O stars are from bonnsai (Schneider et al. 2014b) using the Bonn evolutionary models (Brott et al. 2011;Köhler et al. 2014). Velocities include v edge , v black and v∞ (colons indicate less secure values).

HSH95 Optical
Ultraviolet  Table 5. STIS/MAMA-derived properties of R136 stars beyond a projected distance of 2.05 arcsec (0.5 pc) from R136a1, brighter than m F555W = 16.0 mag). References to optical (literature) spectral types are provided in Table 3. F 1500 fluxes are averages for the range 1500±25Å. Aside from the WN5 stars , temperatures and luminosities are from spectral type calibrations (Doran et al. 2013), while inferred current masses M and ages τ for O stars are from bonnsai (Schneider et al. 2014b) using the Bonn evolutionary models (Brott et al. 2011;Köhler et al. 2014). Velocities include v edge , v black and v∞ (colons indicate less secure values).

HSH95 Optical
Ultraviolet  ground, with bright stars in adjacent slits contributing the remaining 15% of the total. Walborn et al. (1985) have previously presented an atlas of O stars based on IUE/HIRES spectroscopy. The overwhelming majority of these stars are located in the Milky Way. More recently, Walborn et al. (1995Walborn et al. ( , 2002aWalborn et al. ( , 2004 have presented HST and FUSE ultraviolet spectroscopy of Magellanic Cloud O stars, although subtypes and luminosity classes are poorly sampled. Since then, a number of HST programmes have obtained ultraviolet spectroscopy for a range of LMC O stars (e.g. Massey et al. 2004Massey et al. , 2005Massey et al. , 2009). In Appendix A we present a montage of ultraviolet spectroscopic sequences for dwarfs, giants and supergiants (Figs. A1-A4), with luminosity sequences shown in Figs. A5-A11. A log of observations of these datasets is presented in Table A1. . Intra-cluster far-UV spectra of selected inner regions from SE3 and SE4 slits, in close proximity (0.7-1 ′′ , solid) to R136a1 together with R136 H70 (O5 Vz, dotted) for reference. Extended He ii λ1640 emission from R136a1 is apparent in the extracted background spectra.

UV morphology of O stars in the LMC
There is a clear variation in UV morphology amongst O stars from early to late subtypes and from dwarf to supergiant, although we wish to highlight wind line profile variation at an individual subtype/luminosity class. By way of example, we present C iv λλ1548-51 P Cygni line profiles amongst six O4-5.5 dwarfs in the LMC in Fig. 6. Both LH81/W28-5 and R136 H78 possess saturated P Cygni (black) absorption troughs, whereas other stars possess shallow extended absorption troughs which may be flat (e.g. Sk-70 • 60, LH 58-496, Sk-70 • 69) or deeper at high velocities (R136 H70). Within this sample, P Cygni emission components are relatively weak, with the exception of Sk-70 • 69 for which C iv emission is very prominent.
This variation indicates (at least) a third parameter affecting the wind profiles beyond ionization and luminosity class. While further investigation is required, it may well be age, i.e. evolutionary distance from the ZAMS, since some of these stars are in very young regions. While LH 81/W28-5 does have a strong He ii λ4686 absorption line (Massey et al. 2004, their Fig 17), it is substantially weaker than He ii λ4541, indicating a higher luminosity and/or stronger wind than for typical class V. It is well known that mass-loss rates of O stars at a given temperature and luminosity span a wide range of values (Hillier et al. 2003). Binaries and/or binary evolution may also be relevant contributors to this diversity.

UV morphology of stars in R136
In Appendix B we present a series of montages of ultraviolet spectroscopy of stars in R136 (Figs. B1-B9), together with UV templates where appropriate.
In Figure B1 we present HST/STIS spectroscopy of the three luminous WN stars R136a1, a2 and a3 ). These stars have previously been observed in the ultraviolet with HST/GHRS (de Koter et al. 1997), although a1 and a2 were blended in the GHRS SSA 0.22 ′′ entrance aperture. Our long-slit STIS spectroscopy was obtained at a position angle of 64 • and extracted with multispec to disentangle a1 from a2. Nevertheless, the ultraviolet morphologies of all WN stars are extremely similar to one another, as is the case optically (HST/FOS, de Koter et al. 1997) and in the near-IR (VLT/SINFONI, Schnurr et al. 2008). We include an IUE/LORES spectrum of R146 (Brey 88, WN5) for comparison.
HST/STIS spectroscopy of O2-3 (super)giants in R136 is presented in Fig. B3, including the LMC O2-3 III and O2 If reference stars R136 H48 and H36, from our new STIS/CCD dataset. Massey & Hunter (1998)   . Another example of a O2-3 III star is observed beyond the central parsec, namely R136 H102 located 4.8 arcsec to the NE. Figure B4 presents STIS/MAMA spectroscopy of 8 O2-3 V stars in R136 together with the templates BI 237 (O2 V) and R136 H35 (O3 V) based on our new STIS/CCD datasets. All stars exhibit prominent O v λ1371 absorption, P Cygni C iv λλ1548-51, N v λλ1238-42 and negligible/weak absorption at He ii λ1640. Of these, six stars have previously been spectroscopically observed with HST/FOS by Massey & Hunter (1998), who assigned O3 V or O3 III classifications. R136a4, a8, H31, H50 and H62 are newly classified as O2-3 V in this study.
HST/STIS spectroscopy of 7 UV-classified O3-4 dwarfs in R136 is presented in Fig. B5 together with the reference stars R136 H35 (O3 V) and H78 (O4 V) from our STIS/CCD spectroscopy. All O3-4 dwarfs possess weak O v λ1371 absorption, plus prominent, albeit unsaturated C iv λλ1548-51 P Cygni profiles (Fig. A6). Of these, four have previously been spectroscopically observed by Massey & Hunter (1998), who assigned O3 V, except for R136a7 for which O3 III was proposed from HST/FOS spectroscopy. R136 H86 and H149 are newly classified here as O3-4 V from STIS/MAMA datasets, although the latter, centrally located in the SE2 slit, suffers from significant He ii λ1640 contamination by R136a3 (SE1). Figure B6 presents STIS/MAMA spectroscopy for dwarf O4-5 stars in R136 together with new templates R136 H78 (O4 V) and H70 (O5 V) from new STIS/CCD observations. We favour O4 V for R136 H65, H71, H75, H89 and H90 from their morphological similarity to H78, while O4-5 V is preferred for H45, H49, H94 and H119. Classifications are in good agreement with Massey & Hunter (1998) who proposed O3 V for H45, H49 and H94, O5 V for H90 and O3-6 V for H71. Figure B7 extends this sequence to O5 V and O6 V subtypes, exploiting a STIS/CCD subtype of H123 (O6 V). We propose O5 V subtypes for H64 and H69, O6 V subtypes for H30, H77 and H92, and O5-6 V for H141, again in reasonable agreement with Massey & Hunter (1998) who adopted O7 V for H30 and H64, O3-6 V for H69, O3 V for H92. Massey et al. (2002) obtained O5.5 V+O5.5 V for H77.  (Walborn et al. (1995); Massa et al. (2003); Massey et al. ( , 2005Massey et al. ( , 2009, Evans et al. (2010); Bestenlehner et al. (2014). N is the sample size for each category. The UV morphology of R136 H68, located 4.6 ′′ from R136a1, is also consistent with O6 V. Figure B8 presents STIS spectroscopy for O7-8 dwarfs in R136, exploiting new optical subtypes of O7 V and O8 V for R136 H134 and H80, respectively. All stars are newly classified except for H112, for which Massey & Hunter (1998) proposed O8.5 III. A very weak wind signature of C iv λλ1548-51 is observed, via either very shallow P Cygni absorption or a clear blueshifted absorption. We assign O9+ V subtypes for the final group of stars, presented in Figure B9, for which solely photospheric C iv λλ1548-51 absorption is observed, in common with H121 for which O9.5 V is obtained from our STIS/CCD spectroscopy. Low S/N becomes increasingly problematic for these stars, whose far-UV fluxes are typically F1500 ∼ 10 −14 erg s −1 cm −2Å−1 or lower.
Of the remaining stars within 2.05 arcsec of R136a1 that are listed in Table 4 two lie beyond our HST/STIS slit spectroscopy, namely R136 H39 and H137, of which H39 has been observed previously by Massey et al. (2002) and classified O3 V+O5.5 V. Finally, H186 lies close to a much brighter source (H86), within its slit, so we are unable to isolate its UV spectrum.

Wind velocities for R136 stars
The presence of strong, saturated P Cygni profiles amongst the O stars of R136a allows us to measure their wind velocities. To date, wind velocities for large numbers of Galactic OB stars have been obtained from IUE SWP/HIRES spectroscopy (Prinja et al. 1990;Howarth et al. 1997), whereas wind velocities for O stars in the Magellanic Clouds have been relatively scarce (e.g. Prinja & Crowther 1998;Massey et al. 2004Massey et al. , 2005. Here we measure the maximum blueward extend of saturated P Cygni absorption lines (usually C iv λλ1548-51), v black , and use it as a probe of terminal wind velocity, i.e. v∞ = v black (Prinja et al. 1990). The velocity at which the violet absorption meets the continuum, the edge velocity v edge , is usually larger than v black as a result of turbulent motions in the outflow. Such stochastic motions cause some of Figure 7. Representative HST/STIS C iv λλ1548-51 profiles for stars with saturated P Cygni profiles (top), unsaturated P Cygni profiles (centre) and non-P Cygni profiles (bottom). Profiles are shown in velocity space in the LMC rest frame of interstellar C iv 1548.2Å in all cases. the material to move at speeds beyond the local flow speed, hence will cause absorption beyond the terminal speed if the line is fully saturated. If P Cygni profiles are not fully saturated v∞ can be approximated as a function of v edge (see below). In such cases the terminal velocity could be underestimated because the ion may no longer be present in the outer wind.
We measure v black and/or v edge with respect to (LMC) interstellar absorption lines C iv λλ1548-51 and/or Si ii λ1527 (e.g. Howarth & Phillips 1986). The systemic radial velocity of R136 is ∼268 km s −1 (Hénault-Brunet et al. 2012). Individual measurements should be reliable to ±100 km s −1 , so v∞ should be accurate to ±200 km s −1 .
For the present sample, C iv λλ1548-51 exhibits one of three spectral morphologies. A significant subset, typically O2-4 stars and supergiants, exhibit saturated P Cygni profiles, from which v black (and v edge ) can be measured, as shown for R136 H36 in the top panel of Figure 7. A further subset possess strong P Cygni profiles, albeit without saturated absorption features, from which v edge can be mea-sured, typical of O5-6 dwarfs, as shown for R136 H70 in the central panel of Fig. 7. O7-8 dwarfs tend not to exhibit prominent P Cygni profiles, but some show blue absorption that are sufficiently distinct from interstellar C iv, as shown in the lower panel of Fig. 7 for H80.
Individual results are presented in Table 4. The highest terminal wind velocities are v∞ ∼ 3500 km s −1 for the O2 (super)giants H36 and H46, while late O dwarfs often possess v edge ≤ 1500 km s −1 , or no wind measurements are possible. Figure 8 shows edge and black velocities for 28 R136 stars, from which we obtain v black /v edge = 0.8 ± 0.05, which is adopted to estimate wind velocities for stars in which solely edge velocities are measured. For reference, Prinja & Crowther (1998) obtained v black /v edge ∼ 0.7 for LMC stars and v black /v edge ∼ 0.9 for 6 R136 stars (in common with the present sample).
In Table 6 average values of wind velocities in R136 are provided for O2-3, O3-4, O4-5, O5-6, O7-8 dwarfs, plus early O giants and supergiants, including cases for which wind velocities are based upon v∞ = (0.8 ± 0.05)v edge . Amongst dwarfs, the trend towards lower wind velocities for later subtypes is apparent. The standard deviation at O5-6 is unusually large since individual estimates span a wide range from 1510 km s −1 (H77, O6 V from UV morphology) to 2580 km s −1 (H69, O5 V from UV morphology).  Table C1 in Appendix C.
The wind velocity of 1400 km s −1 for R136b (O4 If/WN8) is extremely low, even with respect to O2-3.5 If/WN stars in the LMC (Table C1). In contrast, O2-3 supergiants within R136 possess higher wind velocities than dwarfs, which is the reverse of O2-3 stars elsewhere in the LMC, albeit based on low number statistics. For completeness, we obtain an average wind terminal velocity of 2475 ± 110 km s −1 for the three WN5 stars within R136a. Crowther et al. (2010) obtained v∞ = 2200 − 2600 km s −1 for these stars from their UV to near-IR spectroscopic analysis.
Our new dataset effectively double the sample of LMC early O stars with measured wind velocities. It is important to note that a range of wind velocities have been obtained for individual stars. In part this arises from the methodology, either direct measurement as applied here or line fitting based on the Sobolev with Exact Integration method (SEI, . Still, different authors applying the same approach to the same dataset have obtained significantly different wind velocities. By way of example, de Koter et al. (1998) obtained v∞ = 3000 km s −1 from HST GHRS spectroscopy of R136a5, while  obtained 3400 km s −1 from the same dataset; we find 3045 ± 200 km s −1 from STIS spectroscopy. Mindful of concerns about systematic differences between alternative approaches, an SEI analysis of H36 was carried out (R.K. Prinja, priv. comm.) revealing v∞ = 3500 km s −1 , v turb /v∞ = 0.15 for a β=1 velocity law, in agreement with the v black result. This does not imply that all such measurements will be entirely consistent, although it does suggest that there is no inherent systematic offset. Despite the low number statistics, wind velocities of O2-4 dwarfs in R136 are somewhat lower than elsewhere in the LMC. By way of example, the average wind velocity of 8 R136 O3-4 dwarfs is 2525 km s −1 , versus 2900 km s −1 for 5 counterparts elsewhere in the LMC (and 3040 km s −1 for Galactic O3-4 dwarfs). The reason for this is not clear, if not simply due to small sample sizes. If physical, it is likely not the result of a lower metallicity in 30 Doradus (v∞ ∝ Z 0.13 : Leitherer et al. 1992) relative to the rest of the LMC, although youth may be a factor, recalling the discussion linked to Figure 6.
Wind velocities for individual R136 stars are presented as a function of stellar temperature in Fig. 9 To date, comparisons of empirical (reduced) wind momenta of O stars with theoretical predictions either adopt wind velocities or are based on calibrations. Pre-empting our stellar mass estimates from Section 3.4, we are able to compare terminal wind velocities with escape velocities. Lamers et al. (1995) obtained v∞/vesc = 2.65 ± 0.2 for 16 Galactic O stars earlier than O7. From 35 O2-6 stars in R136, we find v∞/vesc = 2.5 ± 0.4, in agreement with the Galactic sample. Therefore, although R136 early O stars possess ∼15% lower wind velocities than Galactic counterparts, the ratio v∞/vesc is unchanged. Still, considerable variation is observed for other O stars in metal-poor galaxies (Garcia et al. 2014).

Physical parameters for R136 stars
De Koter et al. (1998) provided UV-derived parameters of the brightest dozen or so members in R136, while Massey & Hunter (1998) provided a more complete Hertzsprung-Russell diagram from HST/FOS spectroscopy. Since then, Doran et al. (2013) have provided updated physical parameters for stars within R136 using contemporary calibrations for O stars in the LMC, based on the Galactic scale from Martins et al. (2005).
We defer more robust R136 O-type classifications until future papers in this series, but already our UV morphological classification allows us to present a preliminary, near complete Hertzsprung-Russell diagram for luminous stars within a radius of 2.5 arcsec (0.6 parsec) of R136a1. Temperatures and bolometric corrections of O stars are adopted from Tables 5-7 of Doran et al. (2013), while properties of Of/WN stars are adopted from their Appendix A. HST/WFC3 photometry (de Marchi et al. 2011), except in particularly crowded regions (recall a1/a2, a4/a8), and interstellar extinctions provide absolute magnitudes as set out in Sect. 2.3.
In Figure 10 we present the first spectroscopicallyderived Hertzsprung-Russell diagram for stars brighter than mF555W = 16.0 mag within 0.6 pc from R136a1. A dotted line corresponding to mF555W = 16.0 mag for an adopted AF555W = 1.72 mag is indicated, highlighting the region where our sample is complete. A total of 25 stars (26 including R136c), mostly early O dwarfs, are more massive than ∼50 M⊙. The remaining stars are mid to late O dwarfs with initial masses above ∼20 M⊙.

Ages and masses of individual stars in R136
LMC isochrones for rotating stars with veq = 100 km s −1 from Brott et al. (2011) and Köhler et al. (2015) are included, together with tracks for non-rotating 25, 40, 60, 100 and 200 M⊙. These tracks were selected as respresentative since they closely match the average rotation rate of O stars in 30 Doradus (Ramírez-Agudelo et al. 2013), though do not differ significantly from non-rotating models. It is apparent that, R136b aside, all stars with initial masses above 100 M⊙ lie between the 0-1 Myr isochrones, whereas stars with masses between 40-100 M⊙ span 0-2.5 Myr, and those below 40 M⊙ typically span 0-4 Myr, although uncertainties in physical parameters do not exclude ages below ∼2 Myr, with the exception of R136 H73 (O9+ V) for which an age of ∼5 Myr is favoured.
For quantitative mass and age determinations we have exploited bonnsai 2 (Schneider et al. 2014b) which implements a Bayesian method of quantifying ages and masses Table 7. Comparison of mass and age estimates for the WN5h stars in R136a from Geneva Yusof et al. 2013) and Bonn (Köhler et al. 2015) based on physical and chemical abundances obtained by Crowther et al. (2010), incorporating the adjustment to the absolute magnitude for R136a2, as discussed in the present study.  Table 4, and include results for the WN5h stars solely based upon their positions in the HR diagram, i.e. for consistency initially neglecting their He-abundances.
As suggested by Fig. 10, application of the Bonn evolutionary models (via bonnsai) to the three WN5h stars implies very young ages of ≤0.3 Myr in all cases. Low initial rotation rates -favoured by observations of O stars in 30 Dor (Ramírez-Agudelo et al. 2013) -will lead to a redward evolution on the main sequence in evolutionary models, so young ages are required to explain the blue location of the WN5h stars. The next most luminous stars are extreme O supergiants, R136a5 (O2 If/WN5), a6 (O2 If) and b (O4 If/WN8), each with current masses of 100-150 M⊙, plus O2-3 V stars R136a4 and a8. Once again, young ages of 0.6±0.3 Myr are inferred for these stars, with the exception of R136b whose lower temperature suggests an age of ∼1.5 Myr. Schnurr et al. (2009) have previously argued that R136b, offset by a projected distance of 0.5 pc is more evolved than other core members of R136 (though see Appendix of Hainich et al. 2014).
Consequently, the HR diagram positions of the WN5h stars and the most luminous O stars in R136a favours very young ages. However, it should be emphasised that envelope inflation is implemented in the Bonn models, which is most pronounced for the most massive models (≥ 50M⊙), since these lie closest to the Eddington limit (Sanyal et al. 2015). Inflation moves isochrones to lower temperatures, such that younger ages are inferred for stars with models accounting for inflation, compared with those that do not.
Turning to the lower mass stars, all O2-6 stars lie between the 0-3 Myr isochrones in Fig. 10, although all stars do not lie along a single isochrone. This is unsurprising given the approximate physical parameters of individual stars, initial rotation rate distributions, and uncertain binarity. For an extreme mass ratio system, the primary would completely dominate the spectrum of the system, so its location in the HR diagram would not be significantly displaced. In contrast, for a mass ratio of unity, individual stars would lie 0.3 dex lower in the HR diagram than the composite system, so the age would be overestimated. The effect would be modest for the highest mass systems at very young ages, but would become more significant for lower mass systems. Indeed, later O stars possess higher age estimates (3-6 Myr), although the inferred properties for such stars are least secure of the entire sample.
Individual age estimates follow solely from the position of stars in the HR diagram. We are in the process of analysing STIS/CCD datasets of R136a members, from which rotational velocities and helium abundances can be obtained. For the moment, He abundances for the WN5h stars are available . From their luminosities and He-rich atmospheres, we can infer minimum ages since the surface helium mass fraction sets a lower limit to the core helium mass fraction. The minimum ages of R136a1-3 are 1-1.3 Myr, in reasonable agreement with results for the O stars. In order to simultaneously reproduce the high surface temperatures of the WN5hs stars, Bonn evolutionary models favour rapid initial rotation (∼400 km s −1 ), and ages of 1.1-1.6 Myr, as shown in Table 7. These results compare favourably with Crowther et al. (2010) who applied Geneva evolutionary models for very massive stars (Yusof et al. 2013), also requiring rapid rotation (vrot/vcrit = 0.4 or vrot ∼ 400 km s −1 ) and ages of ∼1.5 Myr (see Table 7). Details of the Bonn and Geneva codes differ significantly regarding the consideration of mixing, core overshooting and the treatment of envelope inflation. Both employ the metallicity-dependent Vink et al. (2001) prescription for mass-loss on the main-sequence, so estimates of Minit from current luminosities ought to be consistent, even if there are uncertainties in their validity for very massive stars (see Crowther et al. 2010;Gräfener & Vink 2013). Figure 11. Probability density distribution of the apparent ages of UV-derived spectral types of individual O stars (excluding R136a1, a2, a3) based on bonnsai (Schneider et al. 2014b) and main-sequence evolutionary models (Brott et al. 2011;Köhler et al. 2015), favouring an age of 1.5 +0.3 −0.7 Myr. The shaded area is a bootstrapped ±1σ estimate of the age distribution.

R136 cluster age and mass
Early estimates of the age of R136 were 3-4 Myr on the basis that the WN stars were classical He-burning stars. Subsequently, a lower age of <1-2 Myr was obtained once it was recognised that the WN stars are core H-burning stars, supported by a rich early O-type population (de Koter et al. 1998, Crowther & Dessart 1998. The most extensive study of massive stars to date has been by Massey & Hunter (1998) who inferred a very young age of ≤1 Myr using the Vacca et al. (1996) O star temperature scale, or ≤2 Myr from the Chlebowski & Garmany (1991) scale, although the majority of their sample was drawn from the R136 region, rather than R136a itself.
Here, we have combined the age estimates of 52 O stars in R136a (excluding WN5h stars) from bonnsai to produce a probability density distribution, presented in Fig. 11, from which an age of 1.5 +0.3 −0.7 Myr is obtained. The ±1σ estimate represents the standard deviation in each bin of the age distribution computed with a bootstrapping method. As discussed above, lower ages are obtained for the highest luminosity stars, although the treatment of envelope inflation for the most massive stars (≫ 50M⊙) remains uncertain.
Discrepant age estimates for massive stars in other young massive star clusters (e.g. Arches and Quintuplet in Galactic Centre) have been obtained from their Wolf-Rayet and O star contents, in the sense that the luminous WR stars possess apparent ages which are younger than O stars by 1-2 Myr (Martins et al. 2008;Liermann et al. 2012). Schneider et al. (2014a) have interpreted this as rejuvenation of the most massive stars through stellar mergers. Banerjee et al. (2012) have suggested the very luminous WN5h stars in R136 are post-merger systems. If this were the case, one might anticipate younger (apparent) age estimates for these rejuvenated stars with respect to the massive O star population.
R136 differs from the Quintuplet and Arches clusters in the sense that the luminous O stars also appear young, yet the He-enrichment of WN5h stars favours ages of ∼1.5 Myr and unusually rapid rotation, the latter predicted by stellar mergers (de Mink et al. 2013). According to Schneider et al. (2014a), the probability that the most massive star in a 5×10 4 M⊙ star cluster is a binary product rises from 50% at 1 Myr to 100% at 3 Myr. However, only the most massive system is anticipated to have merged after 1.5 Myr (4 are anticipated after 3 Myr) for a Salpeter-type mass function in the core. Therefore, the youth of R136a does not support a scenario in which all the very massive stars are merger products. Indeed, as we discuss below, the mass within the central parsec is significantly lower than 5×10 4 M⊙ further strengthening this conclusion. A more definitive answer to this question is deferred until the analysis of STIS/CCD datasets, but the apparent youth of the most luminous stars likely arises from the treatment of inflation in single star models than the physics of binary mergers.
A variety of mass estimates of R136 have been obtained. Hunter et al. (1995) infer a mass of 2.2 × 10 4 M⊙ for stars with > 2.8M⊙ out to a radius of 5 pc from HST/WFPC2, suggesting a total mass of ∼ 5 × 10 4 M⊙, while Andersen et al. (2009) derived a mass of 5 × 10 4 M⊙ down to 2.1 M⊙ out to a radius of 7 pc from HST/NICMOS. The actual mass of the R136 cluster may be somewhat lower, as it may suffer contamination from the surrounding halo of NGC 2070, for which Cignoni et al. (2015) infer a global mass of 9×10 4 M⊙. Indeed, the two components contribute equally to the surface brightness at a radius of ∼5 pc from R136a1 according to Mackey & Gilmore (2003).
For a total cluster mass of 5 × 10 4 M⊙, an Initial Mass Function following Kroupa (2001) with an upper mass limit of 300 M⊙, we would anticipate 10 4 /1250 = 40 stars with initial masses in excess of 50 M⊙. Since only 25 stars with such high masses are observed within the central 0.5 pc radius, the actual R136a cluster mass may be somewhat lower. Of course, some massive stars may have been ejected from the cluster (e.g. VFTS 16, Evans et al. 2010). Fujii & Portegies Zwart (2011) suggest a few percent of stars with ∼ 50M⊙ would be dynamically ejected from dense clusters of this mass. If we were to consider a radius of 5 pc, analogous to the region surveyed by Hunter et al. (1995), there are another ∼30 stars with masses in excess of 50 M⊙ (from the census of Doran et al. 2013), supporting their high enclosed mass for the R136 'region'. Sabbi et al. (2012) have proposed that R136 forms part of an ongoing merging binary cluster system within NGC 2070, on the basis of a clump of older stars located ∼5 pc to its NE (see also Cignoni et al. 2015). If this clump were to reflect stars in the process of dynamically accreting, one might expect older stars to have already merged with R136 itself (see Gieles 2013). Indeed, some of the apparently older O stars in R136 (e.g. R136 H73) may arise from this clump, although the determination of physical parameters for these late-type O stars awaits detailed spectroscopic analysis of our optical STIS/CCD datasets. The evolutionary status of some stars in the NE clump (e.g. Mk 33Sb, WC5) implies an apparent age of 3-4 Myr, providing these are coeval (Melnick 1985;Walborn & Blades 1997). The age of this region may be higher, should the highest luminosity stars be rejuvenated binary products.  Doran et al. (2013) have undertaken a study of the massivestar content of 30 Doradus, together with the contribution of R136 to the ionizing and mechanical feedback of the entire star forming region, albeit reliant upon an incomplete census of R136 itself. We are now able to reassess the total number of very luminous stars in 30 Doradus, updated to include tailored parameters of massive stars from Bestenlehner et al. (2014) based on VFTS observations . Table 8 lists 29 stars within 30 Doradus which are more luminous than log(L/L⊙) ≈ 6.2, corresponding to initial masses of ∼ 100M⊙ or more for zero age main-sequence stars. Only 9 (31%) of the most luminous stars in 30 Doradus lie within R136 itself! A similar number (10) of very luminous stars are located in the halo around R136, at projected distances of 1.5-5 pc 3 The other (11) high luminosity stars in 30 Doradus lie ≥10 pc away from R136a1, approximately half within NGC 2070 and the remainder elsewhere. The majority of these are lu- 3 We have excluded VFTS 512 (O2 V) from the Doran et al. (2013) statistics since the luminosity obtained by Bestenlehner et al. (2014) is log(L/L ⊙ ) = 6.0. Source 1222 from Parker (1993) has also been excluded owing to its O9 V spectral type from Walborn et al. (2002c), as has VFTS 1014 (Parker 863) despite Bestenhehner et al. obtaining log(L/L ⊙ ) = 6.2 since this ARGUS spectrum is likely composite (O3 V + mid/late O). Similarly, we favour the lower luminosities (log L/L ⊙ ) = 6.0 -6.1) inferred by Doran et al. (2013) for R140a2 (WN5+O), Mk 49 (WN6) and BAT99-96 (WN8) over those by Hainich et al. (2014), who obtained significantly higher values by assuming the WR dominated the continuum flux (R140a2) or on the basis of high interstellar extinctions (Mk 49, BAT99-96).
The overwhelming majority of the very massive stars in 30 Doradus are early O supergiants, transition Of/WN stars or luminous WN stars. Crowther et al. (2010) proposed that very massive stars may possess O-supergiant spectral appearances on the main sequence, owing to their proximity to the (classical) Eddington limit. However, R136 a4, a8 and Mk 25 indicate that very luminous O2-3 dwarf stars do exist despite having classical Eddington parameters of Γe ∼0.5, based upon bonnsai-derived masses.

INTEGRATED FAR-ULTRAVIOLET SPECTROSCOPY
Ultraviolet spectroscopy of the central region of NGC 2070 was first obtained with IUE (e.g. Koornneef & Mathis 1981;Vacca et al. 1995), while the advent of HST/GHRS permitted high resolution spectroscopy (Walborn et al. 1992). The integrated R136 spectrum revealed prominent He ii λ1640 emission, together with strong N v λλ1238-42 and C iv λλ1548-51 P Cygni features, with stellar Si iv λλ1393-1402 absent. Figure 12 presents the integrated far-UV spectrum of R136 considering all sources up to 0.5 pc from R136a1. Its appearance closely reflects earlier descriptions, He ii λ1640 emission is relatively strong (W λ ∼ 4.5Å) and broad (FWHM∼10Å or 1800 km s −1 ), so the integrated spectral morphology is closer to an Of/WN supergiant than a genuine WN star -recall Figure 4 of Walborn et al. (1992) which compares R136a to Melnick 42 (O2 If). A comparison between the 2×2 arcsec aperture GHRS/G140L observation of R136a from April 1991 (GO 1188, PI D.C. Ebbetts) and Table 8. Census of the most luminous (log(L/L ⊙ ) ≥ 6.2) objects in 30 Doradus, sorted by projected distance from R136a1, drawn from the present study, Crowther et al. (2010), Doran et al. (2013) or Bestenlehner et al. (2014). Initial mass estimates from bonnsai (based upon estimated temperatures and luminosities) are limited to apparently main-sequence stars.  an equivalent region of our STIS dataset reveals excellent agreement, and is morphologically identical to the integrated spectrum shown in Fig. 12 aside from a 30% lower continuum flux. Fig. 12 includes the co-added spectrum of all 70 stars for which F1500 ≥ 5 × 10 −15 erg s −1 cm 2Å−1 , which collectively contribute ∼77% of the integrated continuum flux, with the difference resulting from UV faint late O-type stars, plus the unresolved early B population. The cumulative spectrum of the 7 most luminous members of R136a -each with inferred initial masses in excess of 100 M⊙ -comprising three WN stars (a1-a3), one Of/WN star (a5), an O2 supergiant (a6) plus two O2-3 dwarfs (a4, a8) -is also included in Fig. 12. Collectively, these most closely resemble a weak lined mid-WN integrated star. It is apparent that the He ii λ1640 emission line flux, and one third of the continuum flux originates from these few stars, the former dominated by R136a1-a3. The remaining 63 far-UV bright stars cumulatively resemble an early non-supergiant O star, with prominent O v λ1371, and P Cygni N v λλ1238-42, C iv λλ1548-51.
It is apparent that prominent He ii λ1640 emission in R136 arises from the presence of VMS (initial masses ≥ 100M⊙). Usually, the upper mass limit assumed in population synthesis models is 100 M⊙ so a synthetic model for R136 would fail to reproduce its He ii λ1640 emission. Indeed, if this feature is observed in the integrated light of sufficiently young (unresolved) high mass star clusters, one might legitimately conclude that the stellar mass function extends well beyond 100 M⊙.
At very young ages, early O stars will dominate the UV appearance of massive clusters, so strong O v λ1371 absorption will be present. The absence of mid to late O supergiants will ensure that P Cygni Si iv λλ1393-1402 is absent. At later ages (3-5 Myr) O v will weaken while Si iv will strengthen, plus He ii λ1640 from classic WR stars will appear. Therefore, the simultaneous presence of O v absorption, He ii emission together with the absence of P Cygni Si iv would suggest both a very young age and the presence of very massive stars. To date, the most promising example of such a star cluster is A1 within NGC 3125 (Chandar et al. 2004;Hadfield & Crowther 2006;Wofford et al. 2014).
In addition, since our datasets enable spatially resolved spectroscopy of the constituents of R136, we are able to quantify the contribution of individual stars to the integrated UV continuum. 5 stars brighter than mF555W = 17.0 mag located within 2.05 arcsec of R136a1 lie beyond our STIS slit spectroscopy, so we have estimated their fluxes by adopting an average F1500/FF555W ratio from cluster members; they are presented in parentheses in Tables 4 and B1. We estimate ∼2% of the integrated UV flux originates from these stars, primarily R136 H39 (O3 V+O5.5 V, Massey et al. 2002). Figure 13 presents the cumulative 1500Å flux of R136a to a radius of 0.5 pc, versus a ranked list of the 70 far-UV brightest stars in this region. The three brightest sources alone (R136a1-a3) contribute 20% of the far-UV continuum flux, while the 18 brightest far-UV sources provide 50% of the integrated far-UV flux. Walborn et al. (1992) cite an higher (pre-COSTAR) estimate of 40% for R136a1-a3 from GHRS acquisition images, according to D. Ebbets and E. Malumuth.

SUMMARY
We have introduced a census of the visually brightest members of the R136 star cluster based on new long-slit HST/STIS spectroscopy. We have obtained ultraviolet, blue visual and Hα spectroscopic datasets using 17 contiguous 52×0.2 ′′ slits sampling the core of R136. In the present study we reveal the following: (i) We classify 51 of the 57 stars brighter than mF555W = 16.0 mag within a radius of 0.5 pc of R136a1 based on their ultraviolet spectroscopic morphology. We confirm previous results of Massey & Hunter (1998) that the bulk of the visually brightest members of R136 are O2-3 stars. For the first time we classify all Weigelt & Baier (1985) stars within R136a, which comprise three WN5 stars (a1-a3), two O supergiants (a5-a6) and three early O dwarfs (a4, a7-a8). Eight additional sources are spectrally classified within the central ∼pc, including R136b (O4 If/WN8).
(ii) We obtain wind terminal velocities from measurements of v black (or v edge ) from P Cygni C iv λλ1548-51 profiles, from which we obtain average O2-3 wind velocities of 2750 and 3120 km s −1 for dwarfs and (super)giants, respectively. Comparisons with early O dwarfs elsewhere in the LMC and Milky Way indicate somewhat lower wind velocities in R136.
(iii) We estimate physical parameters for the visually brightest members of R136 based on the temperture calibration of LMC O-type stars from Doran et al. (2013) and exploit bonnsai (Schneider et al. 2014b) to estimate masses and ages based on rotating evolutionary models appropriate for the LMC (Brott et al. 2011, Köhler et al. 2015. Three WN stars (R136a1-a3), two Of/WN stars (R136a5 and R136b), newly identified O2 If R136a6 and O2-3 V stars R136a4 and a8 possess initial masses ≥100 M⊙.
(iv) We obtain a cluster age of 1.5 +0.3 −0.7 Myr from the O star content, which is consistent with that estimated for the luminous WN5h stars by Crowther et al. (2010), although apparent ages of these and other luminous O stars is 0.6±0.3 Myr, likely to arise from the treatment of envelope inflation in evolutionary models (Sanyal et al. 2015), rather than rejuvenation following binary mergers. The mass and youth of R136a indicate that no more than the most massive system is potentially a binary product (Schneider et al. 2014a). More definitive conclusions await analysis of STIS/CCD spectroscopy, permitting rotation rates and He abundances of O stars to be obtained. We identify 25 stars more massive than 50 M⊙ within 0.5 pc of R136a1 indicating a lower enclosed mass than previous estimates, although reasonable consistency is achieved with respect to Hunter et al. (1995) and Andersen et al. (2009) in the R136 'region' from the total number of high mass stars within a 5 pc radius.
(v) Incorporating results from our earlier VFTS-based census (Doran et al. 2013;Bestenlehner et al. 2014) we estimate a total of 29 stars more luminous than log(L/L⊙) = 6.2 in the entire 30 Doradus nebula, of which 26 possess O2-4 If, Of/WN or WN subtypes. Only 9 stars are located within R136 itself, with an additional 9 stars within the extended halo, 1.5-5 pc away from R136a1 (Melnick 1985) and another 11 stars at much greater distance (10-120 pc).
(vi) We have also considered the integrated ultraviolet spectrum of R136a in the context of its constituent members. The prominent He ii λ1640 emission feature arises primarily from the most massive stars. Specifically, close to 100% of the He ii λ1640 flux and 32% of the far-UV continuum originates in the 7 stars for which we estimate (initial) masses above 100 M⊙, the former dominated by from R136a1-a3. Usually, such stars are neglected in conventional population synthesis models. Prominent He ii λ1640 emission in the integrated spectrum of suitably young star clusters -prominent O v λ1371 absorption without strong Si iv λλ1393-1402 P Cygni emission -would favour a mass function that extends well above this value. Indeed, extremely strong λ1640 emission in cluster A1 within NGC 3125 favours extremely massive stars therein (Chandar et al. 2004;Hadfield & Crowther 2006;Wofford et al. 2014).
In future papers in this series we will present the corresponding blue, visual and Hα STIS/CCD spectroscopic datasets. These will be used to refine spectral classifications, obtain more robust physical and wind properties, and search for close binary systems. These will enable a refined Hertzsprung-Russell diagram for R136a. The multiepoch nature of the CCD datasets will enable an attempt at identifying close binaries (see Cottaar & Hénault-Brunet 2014), together new HST FGS observations of R136a (S.E. de Mink, P.I.).  Figure A1. Ultraviolet morphological progression of O2-5 dwarfs, showing strong, broad O v λ1371 absorption at O2-3, together with strong P Cygni N v λλ1238-42 and C iv λλ1548-51. O v weakens at O3.5 and disappears at O4, while C iv P Cygni weakens at O5. Figure A2. Ultraviolet morphological progression of O5-9.5 dwarfs, showing prominent broad, unsaturated P Cygni C iv λλ1548-51 at O5 V, which weakens in both emission strength and absorption width at O6-7 V and becomes photospheric at O9.5V. N v λλ1238-42 is apparent at O5 V but also weakens at O6-8 V and is undetected at O9.5 V. Figure A3. Ultraviolet morphological progression of O2-8 giants, showing strong, broad O v λ1371 absorption at O2-3, with prominent N v λλ1238-42 and C iv λλ1548-51 P Cygni profiles until O6 giants, and unsaturated thereafter. Figure A4. Ultraviolet morphological progression of O2-9.7 supergiants, showing prominent O v λ1371 absorption at O2, strong, saturated C iv λλ1548-51 P Cygni profiles until O8 supergiants, Si iv λλ1393-1402 P Cygni appearing weakly at O4 and saturating for O6 and later subtypes. He ii λ1640 emission is prominent between O2-6, but thereafter blended with iron forest features. Figure A5. Montage of O2 luminosity sequence. Strong/broad O v λ1371 absorption is common to all O2 subtypes. C iv λλ1548-51 and N v λλ1238-42 are prominent P Cygni profiles at all luminosity classes, although the former is unsaturated in dwarfs/giants. Note the reversal in the He ii λ1640 line from absorption to strong P Cygni emission from dwarf to Of/WN. Figure A6. Montage of O2-3 luminosity sequence. O v λ1371 absorption is common to all luminosity classes. Note the reversal in the He ii λ1640 line from absorption to strong P Cygni emission from dwarf to Of/WN. Figure A7. Montage of O3.5 luminosity sequence. O v λ1371 absorption weakly is present, with Si iv λλ1393-1402 P Cygni appearing in Of/WN stars together with prominent He ii λ1640 P Cygni emission. Figure A8. Montage of O4-5 luminosity class. O v λ1371 absorption is absent, although P Cygni N v λλ1238-42 and C iv λλ1548-51 remain prominent for all classes, albeit unsaturated in dwarfs/giants. P Cygni Si iv λλ1393-1402 and He ii λ1640 emission are seen in O4 supergiants. Figure A9. Montage of O6 luminosity sequence. C iv λλ1548-51 and N v λλ1238-42 exhibit a weak P Cygni profile for dwarfs, in contrast to strong P Cygni profiles at class II-I, with Si iv λλ1393-1402 also exhibiting saturated P Cygni profiles for supergiants. Narrow He ii 1640 emission is also observed in supergiants. Figure A10. Montage of O7-8 luminosity sequence. C iv λλ1548-51 exhibits a narrow, weak P Cygni profile for dwarfs, with N v λλ1238-42 barely present. Both lines display a prominent P Cygni profile at Class II (unsaturated for N v), while Si iv λλ1393-1402 exhibit saturated P Cygni profiles amongst supergiants. Figure A11. Montage of O9-B0.5 luminosity sequence. C iv λλ1548-51 is photospheric in dwarfs, N v λλ1238-42 absent, while C iv and Si iv λλ1393-1402 are saturated P Cygni profiles in luminous supergiants. C iii λ1247 is also prominent in luminous supergiants. Figure B1. HST/STIS ultraviolet spectroscopy of WN stars in R136, together with an LMC template WN5 star R146 (Brey 88=BAT99-117, IUE SWP/LORES) Figure B2. HST/STIS ultraviolet spectroscopy of Of/WN stars in R136, together with LMC templates Sk-67 • 22 and Melnick 37 (see Table A1) Figure B3. HST/STIS ultraviolet spectroscopy of O2-3 (super)giants in R136, together with LMC templates R136 H48 (O2-3 III) and H36 (O2 If), see Table A1. Figure B4. HST/STIS ultraviolet spectroscopy of O2-3 dwarfs in R136, together with templates BI 237 (O2 V) and R136 H35 (O3 V), see Table A1.   Weigelt & Baier (1985, WB85) or Hunter et al. (1995, HSH95). A F555W is calculated from observed colours for an adopted reddening law of R 5495 = 4.1, and DM = 18.49 (Pietrzynski et al. 2013). F555W-band photometry shown in italics is from Hunter et al. (1995) Massey et al. (2005)