3 Properties of the HII region sample

The properties of our galaxy sample as given by the NED database and the adopted distances, are summarised in Table 1. For NGC 3351 and the Virgo cluster member NGC 4321 we adopt the Cepheid distances from Freedman et al. (2001). The other two Virgo galaxies (NGC 4254, NGC 4303) are member of the same subgroup as NGC 4321 (Boselli, private communication). We therefore adopt an identical, approximate distance of 16 Mpc. The distance of NGC 3521 is taken from Tully's (1998) Nearby Galaxy Catalog.

A total of 121 spectra were extracted from the 95 slitlets. Nebular emission lines were detected in 88 spectra; 85 correspond to extra-nuclear regions.

Figure 1: Comparison of metallicities O/H of our H  II regions derived from various empirical calibrations. The Kobulnicky et al. (1999) calibration is taken as a reference (x-axis). Different symbols show O/H derived from the Pilyugin P method (filled triangles), and the R23 methods of Zaritsky et al. (1994, squares), and Edmunds & Pagel (1984, stars). See comments in text.

3.1 Metallicities

The metallicity O/H of the H  II regions has been estimated using the following empirical calibrations: the calibrations of Kobulnicky et al. (1999, hereafter KKP) using [O  III]  $\lambda\lambda$4959, 5007/[O  II] $\lambda$3727 and ([O  III]  $\lambda\lambda$4959, 5007+ [O  II] $\lambda$3727)/ ${\rm H}\beta$ (=R₂₃) based on the photoionisation model grid of McGaugh (1994), the similar P-method of Pilyugin (1991), and the older R₂₃ calibrations of Edmunds & Pagel (1984) and Zaritsky et al. (1994).

The O/H abundances obtained from these methods are compared in Fig. 1. Unsurprisingly rather large differences are obtained. As well known, at abundances $12 + \log({\rm O/H})\la 8.5$-8.6 the various R₂₃ methods yield similar results, while the differences increase towards higher metallicities (see e.g. comparison in Pilyugin 2001). Systematically lower values are found from the P-method of Pilyugin (2001). Although calibrated only for regions with $12 + \log({\rm O/H})\la 8.6$, this could indicate a systematic overestimate of the absolute metallicities using the other methods. To ease comparisons with the recent study of BK02 of metal-rich H  II regions we subsequently adopt the KKP calibration by default except otherwise stated.

The metallicity distribution of our entire sample is shown in Fig. 2. The mean metallicity is $<\!12 + \log({\rm O/H})>~= 8.88 \pm 0.22$ ( $8.57 \pm 0.24$) using the KKP (Pilyugin's P) calibration. The vertical dashed line indicates the solar value ( $12 + \log({\rm O/H})=8.92$) adopted in McGaugh's calculations used for the calibration of KKP.

Figure 2: Metallicity distribution of our H  II region sample (solid line) based on the empirical R23 calibration of Kobulnicky et al. (1999). The distribution of O/H for the WR regions is shown by the dashed line. The vertical line indicates the solar value adopted in the models used by these authors.

Figure 3: Comparison of metallicities O/H derived from the empirical calibration of Kobulnicky et al. (1999, KKP) with the determination in regions with measured electron temperature $T_{\rm e}([{\rm OII}])$. The errorbars include only the uncertainty on the [O  II] $\lambda$7325/ ${\rm H}\beta$ ratio. The diagonal shows the one-to-one relation. Discussion in text.

3.3.1 Regions with direct ${T}_{\rm e}$ determinations

The transauroral [O  II] $\lambda$7325 line has been detected in 11 H  II regions allowing thus a direct determination of the electron temperature from [O  II] $\lambda$7325/[O  II] $\lambda$3727. Other potential electron temperature indicators, e.g. [S  III] $\lambda$6312, [N  II] $\lambda$5755, are too weak or could not be measured due to the limited spectral reolution. Electron densities are determined from [S  II]  $\lambda\lambda$6717, 6731. $T_{\rm e}($O  II) and the resulting ionic abundance ratios of O⁺⁺/H⁺ and O⁺/H⁺were derived using this temperature for both ions (the atomic data are those listed in Stasinska & Leitherer 1996).

As shown in Fig. 3 the resulting O/H abundances (assuming O/H = O⁺⁺/H⁺+O⁺/H⁺) are on average found to be lower than those derived from the KKP calibration, the largest metallicity being closer to solar. However, the O/H derived here are lower limits, due to the strong temperature gradients expected at high metallicities (see Stasinska 2002). A deeper discussion of the abundances in our objects taking into account the observational constraints from the entire emission line spectrum is deferred to a forthcoming publication. For the purpose of the present paper, it is sufficient to note that the bulk of our H  II region sample with low values of R₂₃ have metallicities close to and above solar.

Figure 4: Upper panel: Histogram of ${\rm H}\alpha$ luminosities of the entire H  II region sample (solid) and the regions with WR detections (dashed). Lower panel: Number of WNL stars (derived from blue WR bump assuming WN7 type) as a function of the number of equivalent O7V stars derived from ${\rm H}\alpha$ luminosity. Note that the x-axis of both plots correspond directly through the definition of N07V (see footnote). Filled (open) symbols indicate regions with certain (candidate) WR detections (see text). The lower panel shows that secure WR detections are only found for regions with $N_{\rm WNL} \protect\ga$ 2-3, as physically expected. This indicates that the WR sample in our sample of regions with $\log({\rm H}\alpha) \protect\ga 38.2$ is fairly complete (cf. text).

 

 

Table 3: Statistics of WR regions.

Galaxy	# blue bump	# C  IV $\lambda$5808	# C  III $\lambda$5696	cand. blue bump	cand. C  IV $\lambda$5808	cand. C  III $\lambda$5696
NGC 3351	2			4		2
NGC 3521	4	2	1	6	1	1
NGC 4254	9	8	1		1	1
NGC 4303	9	4	3		3	2
NGC 4321	3		3	5	1	1
total	27	14	8	15	6	10

3.2 ${\rm H}\alpha$ luminosities, WR and O star populations

The histogram of the ${\rm H}\alpha$ luminosity of the H  II regions, as measured from our spectra, is shown in the upper panel of Fig. 4. As seen from this figure, $\sim$75% of the H  II regions correspond to giant extra-galactic H  II regions characterised by $L({\rm H}\alpha) \ga 10^{38}$ erg s^-1(Kennicutt 1984, 1991), while the remainder are less luminous objects similar to normal Galactic H  II regions. The corresponding number of equivalent O7V stars, $N_{\rm O7V}$ plotted in the lower panel of Fig. 4, ranges from $\sim$0.15 O7V stars (i.e. presumably corresponding to $\sim$1 late O or early B stars) to $\sim$400 O7V stars for the brightest region.

Our search for WR features in metal-rich H  II regions proved quite successful yielding with 27 WR detections a sample of unprecedented size (cf. Castellanos 2001; Bresolin & Kennicutt 2002). The number of regions where different WR features were detected (hereafter called "WR regions'') at various levels of confidence are listed in Table 3. The certain WR detections (defined as $\ge$$2 \sigma$ detections) are listed in Cols. 2-4; "candidate'' WR regions with emission line detections $1.1 \le \sigma < 2$ are given in Cols. 5-7. Visual inspection of the spectra yield essentially the same detection of the "certain'' WR regions. To illustrate the quality of our data sample spectra of a secure WR region and a candidate region are shown in Fig. 5.

Figure 5: Left panels: FORS1/VLT spectra of a WR region in NGC 4254 (top) and a candidate region in NGC 3351 (bottom) showing also the main line identifications. Right panels: Zoom on the spectral region of the blue WR bump (top) and the red WR bump (bottom) of the two spectra.

As also clear from Table 3, a large fraction of the WR regions shows signatures of WR stars of both WN and WC types as anticipated from theoretical expectations (Meynet 1995, SV98) and earlier studies of WR galaxies (Schaerer et al. 1997, 1999ab; Guseva et al. 2000). In our sample $\sim$50% of WR regions show WC signatures; predictions from the Meynet (1995) and SV98 models yield $\sim$30-77% at metallicities $Z \sim 0.008$-0.040. At least 1/3 of the WR regions harbour WC stars of late subtypes (WCL), characterised by their strong C  III $\lambda$5696 emission. The C  III $\lambda$5696/C  IV $\lambda$5808 ratio indicates subtypes WC7 or WC8 assuming that the contribution of WN stars to C  IV $\lambda$5808 is negligible; if this were not the case the mean spectral type could be of later subtype. So far relatively few WR "galaxies'' (true starbursts or extra-galactic giant H  II regions) with WCL stars are known (cf. Schaerer et al. 1999b). However, as late WC types are expected to occur preferentially in metal-rich environments (Smith & Maeder 1991; Maeder 1991; Philipps & Conti 1992) the high detection rate of C  III $\lambda$5696 is not surprising.

The ${\rm H}\alpha$ luminosity distribution of the WR regions is shown in Fig. 4 (upper panel, dashed line). Clearly, WR stars are only detected in the brightest regions. This is not due to the flux limit of our observations as can easily be seen by comparison of the smallest WR bump fluxes ( $F({\rm WR})_{2 \sigma} \sim 4\times 10^{-16}$ erg s^-1 cm^-2) with the detection limit of the faintest emission lines ( $F_{\rm lim}({\rm H}\beta) \sim 10^{-16}$ erg s^-1 cm^-2). In fact our observations are essentially deep enough to allow in all galaxies the detection of the blue WR bump of just $\sim$2-3 WNL stars, assuming the average 4650-4686 Å bump luminosity of a WN7 stars of 10^36.5 erg s^-1 (cf. Smith 1991; Schaerer & Vacca 1998). The number of WNL stars derived in this way is plotted in the lower panel of Fig. 4 for regions with certain WR detections (filled squares) and "candidate'' WR regions (open circles).

As our detection limit allows for the detection of few ($\sim$2-3) average WNL stars, the subsample of the brightest regions with $F({\rm H}\beta) \ga 5\times 10^{-15}$ erg s^-1 cm^-2could therefore represent a fairly complete sample of H  II regions "massive''/bright enough to allow a meaningful comparison between WR detections and non-detections. However, a possible bias against regions with small $W({\rm H}\beta )$ may exist (Sect. 2).

In this subsample containing a total of 47 regions we find 20 objects without WR signatures, or a fraction of $\sim$57% regions with WR signatures. Such a high fraction of WR detections compares fairly well with the predictions of 60-80% by Meynet (1995) and Schaerer & Vacca (1998) using the high mass loss stellar evolution tracks at metallicities $1/2.5 \la Z/Z_{\odot}\la 2$ for bursts with a standard Salpeter IMF and an upper mass cut-off $M_{\rm up}=120$ $M_{\odot }$. Given the fact that very young regions (ages 0 to $\sim$1.5-2 Myr) with large expected ${\rm H}\beta$ equivalent widths are notoriously absent (in the present sample and other samples of H  II regions and galaxies) it is, however, not clear how significant this finding is.