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Chemical Abundances and Ionizing Clusters of HII Regions in the LINER Galaxy NGC 4258

We present long-slit observations in the optical and near infrared of eight HII regions in the spiral galaxy NGC 4258. Six of the observed regions are located in the SE inner spiral arms and the other two are isolated in the northern outer arms. A detailed analysis of the physical conditions of the gas has been performed. For two of the regions, an electron temperature has been derived from the [SIII] 6312 A line. For the rest, an empirical calibration based on the red and near infrared sulphur lines has been used. The oxygen abundances derived by both methods are found to be significantly lower (by a factor of two) than previously derived by using empirical calibrations based on the optical oxygen lines. In the brightest region, 74C, the observation of a prominent feature due to Wolf-Rayet stars provides an excellent constraint over some properties of the ionizing clusters. In the light of the current evolutionary synthesis models, no consistent solution is found to explain at the same time both the WR feature characteristics and the emission line spectrum of this region. In principle, the presence of WR stars could lead to large temperature fluctuations and also to a hardening of the ionizing radiation. None of these effects are found in region 74C for which the electron temperatures found from the [SIII] 6312A and the Paschen discontinuity at 8200A are equal within the errors and the effective temperature of the ionizing radiation is estimated at around 35,300K. Both more observations of confirmed high metallicity regions and a finer metallicity grid for the evolutionary synthesis models are needed in order to understand the ionizing populations of HII regions.


INTRODUCTION
Theoretically, the evolution of a young stellar population depends on metallicity at least through two very important effects: the increasing opacity of the stellar material and the dependence of mass loss on metal content in high mass stars. As a consequence of the first, the effective temperature of ionizing stars should be lower in regions of higher metallicity (see for example McGaugh 1991) but, on the other hand, as a consequence of the second, that is if the strength of stellar winds increases with metallicity, the loss of the outer en-⋆ Visiting Fellow, IoA, Cambridge velopes of the most massive stars can increase their surface temperature to very high values. These highly evolved massive O stars are identified with the Wolf Rayet population.
Most theoretical evolutionary models for ionizing star clusters predict the appearance of WR stars -although at slightly different cluster ages depending on the specific details of the models -and in all of them the fraction of WR stars increases with metallicity since the limiting mass for a star to enter the WR phase decreases with increasing metallicity. However, the effect of these stars on their surrounding ionized gas depends on the characteristics of the assumed wind opacity and therefore while some models predict the existence of high metallicity HII regions of high excita-tion (e.g. Stasińska & Leitherer 1996), others do not (e.g. Cerviño & Mas-Hesse 1994). The finding of this extreme HII region population would constitute a firm evidence of the existence of WR stars with high effective temperatures, thus favoring the first kind of models but, even if these regions are not definitively identified, the study of high metallicity HII regions is of crucial importance to investigate the processes of star formation and evolution in high metallicity environments. This investigation gains even more relevance if we take into account the recent observations that point out to a firm connection between star formation and activity in active galactic nuclei (Heckman et al. 1997, González-Delgado et al. 1998) already predicted on theoretical grounds (Terlevich & Melnick 1985;Terlevich et al. 1992), since these nuclei seem to reach rather high metallicities (Phillips et al. 1984; as deduced from the study of their circumnuclear star forming regions. Unfortunately, the term "high metallicity HII region" is rather ill defined. Most authors apply the term to HII regions of solar or oversolar metal content, but the difficulty to derive directly the metallicity of these cool HII regions is well known and the empirical calibrations commonly employed provide estimations with an uncertainty larger than a factor of 2 (see e.g. Díaz et al. 1991). On the other hand, the characteristics of WR features are highly dependent on metallicity. Both the luminosity of the WR "bump", that is the combination of the NIII and the HeII lines at around λ 4660Å, and its equivalent width are predicted to be largest at the highest metallicity and almost an order of magnitude difference is found between the computed WR "bump" intensity relative to Hβ in clusters of half solar and solar metallicities (Schaerer & Vacca 1998). Therefore, the detection of WR features combined with a detailed analysis of the emission line spectra of suspected high metallicity HII regions, can constitute excellent tools for the simultaneous determination of both the metallicity and the age rather accurately in these regions.
In the frame of a long term programme to study high metallicity HII regions we are presently carrying out a search for the high metallicity and high excitation HII region population by performing spectrophotometric observations at a resolution high enough to detect and measure WR features. For that, we have selected from the literature the giant HII regions with solar or oversolar abundance, as deduced from empirical calibrations based on the optical forbidden lines, showing the highest excitation. This corresponds to a ratio of [OIII]/Hβ ∼ 1.0, which is actually quite moderate but still high compared to that expected in high metallicity HII regions. Here we present the first of these investigations involving several HII regions in the galaxy NGC 4258.
NGC 4258 (M106) is classified as SAB(s)bc and Sb(s)II types by De Vaucouleurs (1976) and Sandage & Tammann (1981) respectively. Heckman (1980) and Stauffer (1982) classified its nucleus as belonging to a transition type between LINER and H II-region like emission-line galaxy. As it is addressed by Courtés et al. (1993)(hereinafter C93), the H II region distribution is very peculiar in the sense that three general structures are found: the normal spiral arms, with dense and very bright HII regions following pure spiral shapes as a consequence of the barred spiral type of this galaxy; the faint outer arms (H II region poor), and the anomalous spiral arms, discovered by Courtés & Cruvellier (1961), without star formation and with no evidence of very blue stars (Deharveng & Pellet 1970). It is assumed that the anomalous arms were formed by powerful jets ejected by the active nucleus of the galaxy. Stüwe, Schulz & Hühnermann (1992) showed that the active nucleus must be obscured along the line of sight as is also indicated by its invisibility in IUE spectra. Cecil, Morse & Veilleux (1995) found, from the emission-line flux ratios, LINER-like gaseous excitation along the jets. They also point out the possibility that the circumnuclear region has undergone recent formation of massive stars induced by the jets that are flowing into the disk. However, no evidence for this star formation has been found, maybe due to the obscuration by molecular gas along the line of sight. Finally, a subparsec diameter gaseous disk shows keplerian motion around the centre of the galaxy, which indicates the presence of a mass of 3.6 × 10 7 M⊙ within a radius of less than 0.13 pc (Miyoshi et al. 1995). This is probably the best existing evidence for the presence of a massive black hole at the centre of a galaxy.
The problem of the distance determination for this galaxy still remains unsolved. The quoted distance values range from 3.4 Mpc (Roy, Arsenault & Joncas 1986) to 10.4 Mpc (Sandage & Tammann 1981). The adopted value in the present work is 5.5 Mpc (Martin et al. 1989).
We have observed eight HII regions along two different spiral arms. Four of them have been identified as: 74C, 69C, 5N(A) and 5N(B) in C93. Regions 5N(A) and 5N(B) are located in the northern outer arm at a galactocentric distance of 20.3 Kpc. The other six regions are located in the SE inner arm at 8.05 Kpc from the galaxy centre. Region 74C is a discrete radiosource at 4.9 GHz (Hummel et al. 1989).
Regions 74C and 5N were previously observed by Oey & Kennicut (1993) who derived values of 12+log(O/H) of 8.87 and 8.74 for each of them respectively (the solar value is 8.92). With a ratio of [OIII] to Hβ of 0.98 and 1.44 both regions fit our selection criteria.

OBSERVATIONS AND REDUCTIONS
Our spectrophotometric observations were obtained with the 4.2m William Herschel Telescope at the Roque de los Muchachos Observatory, in 1994 March 13, using the ISIS double spectrograph, with the TEK1 and a EEV CCD detectors in the blue and red arm respectively. Two gratings were used, R300B in the blue and R316R in the red arm, covering from [O II] λ3727 to [S III] λ9532 in three different spectral ranges of ∼1700Å each. The dispersion of 1.4Å pixel −1 with a slit width of 1. ′′ 5 gives a spectral resolution of ∼4Å. The nominal spatial sampling is 0. ′′ 7 pixel −1 . The seeing was ∼1. ′′ 5. The slit was centered at the positions of regions 74C and 5N, as given by Oey & Kennicutt (1993). A third position was observed in order to include other HII regions outlining the SE inner arm. In all cases the position angle was chosen as to maximize the number of HII rgions in the slit. A journal of observations is given in Table 1.
The data were reduced using the IRAF (Image Reduction and Analysis Facility) package following standard methods. The two-dimensional wavelength calibration was accurate to 1Å in all cases. The two-dimensional frames were flux calibrated using two spectroscopic standard stars observed during the same night with a 5. ′′ width slit. The agreement between the individual calibration curves was better than 8% in all cases. The spectra were previously corrected for atmospheric extinction using a mean extinction curve applicable to La Palma observing site. The removal of the atmospheric water-vapour absorption bands in the near infrared (Díaz, Pagel & Wilson 1985), was achieved dividing by the relatively featureless continuum of a subdwarf star observed in the same night as the galaxy. The most critical step in the reduction process was the background subtraction in the near infrared frames. There are numerous OH emission lines of the night-sky spectrum that contaminate the near infrared spectra. This contamination is particularly important in the range from 9376 to 9600Å (Osterbrock, Fulbright & Bida 1997). Hence, the [SIII] λ9532Å and Pa 8 lines intensities could be affected even after the background subtraction is made. Fortunately, the [SIII] λ9069Å and Pa 9 lines are unaffected by these night-sky lines. The observed λ9532/λ9069 ratio is above the theoretical value of 2.44 in four of the observed regions. In these cases, the [SIII] λ9532 line was scaled by means of the more accurate [SIII] λ9069 line intensity measurement. Figure 1 shows the spatial distribution of the Hα flux along the slit for the three different positions observed. Regions 5N A and B are clearly identified on the Hα profile corresponding to PA=255 • . Also regions 74C and 69C, showing a certain degree of substructure, are resolved at position angle PA=133 • . Four different regions are resolved at PA=150 • that we have named GA1, GA2, GA3 and GA4. A fifth region is clearly seen in the Hα profile that is not detected in the near IR frames.

RESULTS
The spectra corresponding to each of the identified regions is shown in Figs. 2 to 4. WR features around λ 4680 A are seen in the spectrum of region 74C and to a lesser extent in that of region 5N.  [OII] [OIII] [SIII] [SIII] [OII] [OIII] [SIII] [SIII] [OIII] [SIII] [SIII] 4000 6000 8000 0 Wavelength (Angstroms) NGC4258 (69C) [OII] [OIII] [SIII] [SIII] The observed line intensities relative to the Hβ line were corrected for interstellar reddening according to an average extinction curve (Osterbrock 1989) and assuming the Balmer line theoretical values for case B recombination (Brocklehurst 1971). The presence of an underlying stellar population is evident in the blue spectra of regions GA2, GA3 and GA4. The Hγ and Hδ Balmer lines are clearly affected by this stellar absorption. An iterative process was [OIII] [SIII] [SIII] [OIII] [SIII] [SIII] applied in order to fit observed and theoretical Balmer line intensities and to obtain the reddening constant c(Hβ). In all cases, Hα, Hγ and Hδ were fitted, except in regions GA3 and GA4 for which Hδ seems to be underestimated. Redden-ing corrected Paschen lines, when measured, are consistent with their theoretical values. The reddening corrected line intensities together with their corresponding errors are given in Table 2 for regions 74C and 69C and Table 3 for the rest of the regions. Also [OIII] [SIII] [SIII]  given in the tables are the extinction corrected Hα flux and the Hβ equivalent width. For all the regions, Hα luminosities are lower than 10 38 erg s −1 , with the exception of region 74C, in which a value of 10 39.3 is found. According to Kennicutt (1983), this region can be classified as a supergiant HII region.

Physical conditions of the gas
Electron densities for each observed region were derived from the ratio of the [SII] λλ 6717, 6731 lines following standard methods (see for example Osterbrock 1989). For regions GA2 and 74C densities of 160 cm −3 and 60 cm −3 respec- [OIII] [SIII] [SIII] For two of the observed regions, 74C and 69C, it has been possible to measure the weak auroral [SIII] λ 6312Å line, which, together with the observed intensities of the nebular lines at λλ 9069, 9532Å, can be used to obtain a value for the electron temperature. Using the expressions given by Osterbrock (1989) and the effective collision strengths recently calculated by Tayal (1997) we find for t(S 2+ ) values of 7500 ± 300 K and 7600 ± 500 K for regions 74C and 69C respectively.
Photoionization models indicate that, for abundances We have used that relation to obtain t(O 2+ ). Finally, Stasińska's (1980) relation between t(O 2+ ) and t(O + ) has been used to calculate the temperature of the t(O + ) zone. These different temperatures for regions 74C and 69C are listed in Table 4. Another independent way of temperature determina-tion is through the measure of the Paschen discontinuity at λ 8200Å. The ratio of the Paschen discontinuity to the Hβ intensity is, after correcting for the derived extinction value of c(Hβ) = 0.48, ∆P a I(Hβ) = (6.9 ± 0.8) × 10 −15 Hz −1 According to the models computed by González-Delgado et al. (1994), for the helium abundance derived for the region 74C (see next section), this ∆ Pa/I(Hβ) implies an electron temperature of 8000 ± 700 K, which agrees with that derived from the [SIII] lines within the errors. Therefore, no temperature fluctuations are apparent in the region which  (Pérez 1997).
For the rest of the regions it was not possible to obtain a direct measure of the electron temperature. An average temperature has been adopted from an empirical calibration by means of the sulphur abundance parameter S23 (S23 = ([SII]6717, 6731 + [SIII]9069, 9532)/Hβ) (Díaz & Pérez-Montero 1999). The values for these temperatures, together with similar ones derived for regions 74C and 69C, are also given in Table 4.

Chemical abundances
Ionic abundances of oxygen, nitrogen and sulphur have been derived following standard methods (Pagel et al. 1992) and using the temperatures found above. These abundances are listed in Table 4. We have assumed that most of the oxygen and sulphur are in the first and second ionization stages and therefore O/H = O + /H + + O ++ /H + and S/H = S + /H + + S ++ /H + . This assumption seems to be justified given the relatively low estimates of the electron temperature.
The helium abundance was determined from the HeI λ 4471 and 6678Å lines. Mean values of 0.070 ± 0.005 and 0.060 ± 0.005 respectively are obtained for regions 74C and 69C. We have estimated the contribution of neutral helium from the expression:

Wolf-Rayet features
A prominent Wolf-Rayet feature has been observed in region 74C (see Figure 6). Assuming a distance to NGC 4258 of 5.5 Mpc (Martin et al. 1989) and a constant extinction value through this region of c(Hβ) = 0.48, as derived from the Balmer and Paschen recombination lines, the total luminosity of the broad HeII feature, without the contribution of the [FeIII] λ 4658Å line, is (1.1 ± 0.1) × 10 38 erg s −1 . This value comprises the broad features of both NIII λλ 4634, 4640 and HeII λ 4686Å lines. The contribution of the NIII lines to the WR bump is metallicity-dependent according to Smith (1991). In our case this contribution represents 0.4 times the total emission; hence we obtain for the HeII line luminosity: L(HeIIλ4686) = (6.6 ± 0.6) × 10 37 erg s −1 The non-detection of NV λλ 4604,4620Å, suggests that the WR stars are of late N or intermediate type. Using the calibration of Vacca8& Conti (1992), 38 ± 5 WN stars are found in this region.
Another fainter Wolf-Rayet feature has been observed in region 5NA. The bump luminosity is 2.2 × 10 36 erg s −1 . The NIII lines contribution represent 0.45 times the total emission, hence the HeII line luminosity is 1.2 × 10 36 erg s −1 , which is compatible with the presence of 1 WN star.

FUNCTIONAL PARAMETERS OF THE OBSERVED HII REGIONS
Three are the fundamental parameters which control the emission line spectra of HII regions (Díaz et al. 1991): the ionization parameter, the shape of the ionizing continuum, and the metallicity. The ionization parameteri. e. the ratio of the ionizing photon density to the particle density -is a measure of the degree of ionization of the nebula and can be deduced from the ratio of two lines of the same element corresponding to two different ionization states, e.g.

[OII]/[OIII] or [SII]/[SIII].
Alternatively, it can also be determined from [OII/Hβ] or [SII/Hα] if the metallicity of the region is known (Díaz 1994). By using a large grid of single star photoionization models the following expressions are found (Díaz et al. 1991, Díaz 1999 We have derived U from the four expressions above. In all cases the value of U derived from the [OII]/[OIII] ratio is systematically lower than the rest thus implying low effective temperatures for the ionizing stars (see e.g. Díaz 1999). We have therefore discarded this value and computed U as the mean of the other three. These adopted ionization parameters -logU -are listed in Table 5 and their uncertainty is estimated to be around ± 0.2 dex.
The shape of the ionizing continuum is directly related to the effective temperature of the stars that dominate the radiation field responsible for the ionization of the nebula. A recent version of the photoionization code CLOUDY (Ferland 1996) has been used to estimate the mean effective temperature of these stars. We have used Mihalas NLTE single-star stellar atmosphere models, with a closed geometry and a constant particle density through the nebula. For all the regions, consistency is found for effective temperatures around 35,000 K -36,000 K. The results from these photoionization models compared to the observations are presented in Table 5.

DISCUSSION
The Hα fluxes for regions 74C and 69C, uncorrected for reddening, are 2692 and 287 × 10 −16 erg s −1 cm −2 respectively. These fluxes are to be compared with 3492 and 332 × 10 −16 erg s −1 cm −2 measured by C93. Considering that our observations have been obtained through a narrow slit and region 74C is rather extended, the agreement can be considered satisfactory. The combined Hα flux of our regions 5NA and 5NB is 187 × 10 −16 erg s −1 cm −2 close to the value of 249 × 10 −16 erg s −1 cm −2 given by C93 for region 5N. Regarding the rest of the regions, we tentatively identify regions GA2, GA3 and GA4 with regions 59C, 58C, 54C of C93. Region GA1 probably corresponds to the outer parts of region 69C.
For each of our observed regions we have calculated the Hα luminosity and the number of hydrogen ionizing photons from the observed Hα flux corrected for their derived reddening (see Tables 2 and 3  where F(Hα) is expressed in erg s −1 cm −2 and D is given in Mpc.
These values are given in Table 6 for the adopted distance of 5.5 Kpc. Only region 74C has an Hα luminosity greater than 10 39 erg s −1 and can be classified as a supergiant HII region as defined by Kennicutt (1983). The rest of the regions have Hα luminosities typical of HII regions in early spiral galaxies, although all of them are greater than 10 37 erg s −1 (Q(H) > 10 49 photons s −1 ) requiring more than a single star for their ionization (Panagia 1973).
Filling factors for each observed region can be determined from the reddening corrected Hα flux, F(Hα), and the derived ionization parameter, U, according to the expression: where αB(H o , T ) is the recombination coefficient for hydrogen and ne is the electron density of the emitting gas. The fillin0 factors computed in this way are listed in Table 6. We have used a value of αB(H o , T ) = 3.76 × 10 −13 cm 3 s −1 , corresponding to T= 7000 K and ne= 100 cm −3 (Osterbrock 1989). For regions 74C and GA2 we have used values of ne= 60 and 160 cm −3 respectively, as derived from the ratio between the red [SII] lines. For the rest of the regions for which only upper limits to the density could be derived, a value of ne=10 cm −3 has been assumed. According to C93 the angular sizes of regions 74C, 69C and 5N are 38.6, 16.8 and 20.5 arcsec respectively. The corresponding linear sizes, at the assumed distance to NGC 4258, are 1.0, 0.45 and 0.55 Kpc . These are comparable to the sizes of HII regions of moderate intensity found in other spiral galaxies (e.g. González-Delgado et al. 1997).
Hα angular effective diameters -that is the diameters containing half the Hα emission -of regions 74C, 69C and 5N are 4.3, 3.7 and 6.1 arcsec respectively (C93) which translates into 115, 99 and 163 pc. From the definition of U it is possible to obtain the angular size of the emitting region using the reddening corrected Hα flux and the derived electron density: where φ is the angular diameter in units of 10 arcsec . This angular diameter does not depend on the assumed distance to the galaxy. For region 74C the derived angular diameter is 3.3 ′′ close to the value of 4.3 ′′ given by C93. For region 69C their observed value of the effective diameter of 3.7 ′′ yields a value of the electron density of ne = 5 cm −3 close to the assumed one of 10 cm −3 . Using this value of ne for the rest of the observed regions, except GA2 for which ne = 160 cm −3 we obtain angular diameters between 1 and 2.6 arcsec.
We can also calculate the mass of ionized hydrogen following the expression: (Osterbrock 1989) where mp is the mass of the proton and The corresponding masses of ionized hydrogen in each observed region are given in Table 6. From evolutionary models of single ionizing clusters and radiation bounded HII regions, a relation between the number of ionizing Lyα photons per second per solar mass and the Hβ equivalent width (Díaz 1999) can be found, which allows the estimation of the ionizing cluster mass by taking into account the cluster evolution. log(Q(H)/M⊙) = 0.86log(EW (Hβ)) + 44.48 For our observed regions all values for Q(H) are around 10 49 photon s −1 except for region 74C for which a value of the order of 10 51 photon s −1 is found. Hence, in the absence of dust, a lower limit for the mass of the ionizing clusters can be estimated by means of the Hβ measured equivalent width and the Hα luminosity for each region. The estimated ionizing cluster masses range from 4 × 10 3 M⊙ for region GA4 to 1.12 × 10 5 M⊙ for region 74C. The high signal-to-noise spectrum of this latter region allows a more detailed modelling. The presence of WR features, in principle, provides a means to constrain the age of the ionizing population. The observed WR "bump" luminosity relative to Hβ (L(WR)/Hβ) is 0.15. According to the models by Schaerer & Vacca (1998), at metallicity Z= 0.008 (0.4 solar), these high values are only found at an age of 4.5 Myr. However, this ratio is strongly dependent on metallicity and changes by more than an order of magnitude between Z=0.008 (0.4 solar) and Z= 0.02 (solar). Therefore, the slightly higher metallicity of region 74C (Z ≃ 0.012) might allow earlier ages for the WR stellar popu- lation. Clearly, the available grid for stellar evolution models is too coarse to elucidate this matter. The same applies to the equivalent width of the "bump" (EW(WR)) whose observed value is 10.8Å. Since the NIII λ 4640Å line is strongly dependent on metallicity (Smith et al. 1996) it is probably more advisable to use only the flux and equivalent width of the HeII line. In that case, our observed values of L(HeII)/Hβ and EW(HeII) are 0.09 and 6.6Å respectively. Again, consistency is found for a model of Z=0.008 and an age of 4.5 Myr.
Since, according to the recent evolutionary models of Leitherer et al. (1999;STARBURST99), the observed value of the equivalent width of Hβ (81Å) points also to an age of around 4.5 Myr, we have computed the emission line spectrum corresponding to this single ionizing cluster. † This set of models makes use of the same evolutionary tracks as the models of Schaerer & Vacca (1998) and therefore provides a selfconsistent frame. The results of the computation can be seen in Table 7 together with the observed emission line ratios. It can be seen that the model ionizing spectrum results too hard and does not reproduce the observations. This is not surprising in view of the low effective temperature found from single star photoionisation models (Table  5). It can be better appreciated in Fig 8 where the spectral energy distribution of the cluster can be compared to that of the Mihalas stellar atmosphere of a star with T ef f = 35300 K which adequately reproduces the data. A high energy tail is clearly present in the ionising cluster as the result of the WR star contribution.
These models assume an enhanced mass loss during the evolution of massive stars, as prescribed by the stellar evolution models of Meynet et al. (1994). More moderate mass loss rates are assumed in the models of Bressan et al. (1993) and Fagotto et al. (1994) and with them a single population of 5.3 Myr can adequately reproduce the emission line spectrum (see Table 7). The corresponding ionizing spectrum is also shown in Fig 8. However, this ionizing cluster provides only a few WR stars (about 3× 10 −5 per solar mass of the ionizing cluster, for a metallicity Z=0.008; García-Vargas, Bressan & Díaz 1995). This is probably not enough to reproduce the observed values of L(WR)/Hβ.
The WR features observed in region 74C are similar to those found in other extragalactic HII regions. Table 8 shows the values of the WR feature intensities and equivalent widths for 74C and other three well studied extragalac- † All the models are computed assuming a standard Salpeter IMF for the stellar ionizing cluster and using the photoionization code CLOUDY (Ferland 1996). tic HII regions: NGC 604 in M33 (Díaz et al 1987;Terlevich et al. 1996), region A in NGC 3310 (Pastoriza et al. 1993) and region A in NGC 7714 (García-Vargas et al. 1997). Also given in the table are the metallicities of the regions in the form 12+log(O/H) and the effective temperature of the ionizing radiation estimated from single star photo-ionization models.
The tabulated values are difficult to explain in the light of the recent models of Schaerer & Vacca (1998). The expected decrease of the intensities of WR features with decreasing metallicity is not clearly observed. However, a decrease of the temperature of the ionizing radiation with increasing metallicity is apparent. In fact, region A in NGC 3310 and region 74C in NGC 4258 show the same L(WR)/Hβ ratio despite their different metallicity (by a factor of about 3). This implies that both regions have similar ratio of WR to O stars, while, at the same time, their effective temperatures are considerably different.
The interpretation of the equivalent widths of WR features is even more difficult, since they may be affected by the continuum from any underlying stellar population. Composite populations have been postulated for region A in NGC 3310 (Pastoriza et al 1993) and region A in NGC 7714 (García-Vargas et al. 1997). Under this assumption, the equivalent widths, EW(WR) and EW(HeII), of the cluster with WR stars are increased to 6 and 4.8Å for region A in NGC 3310 and 3.6 and 2.5Å for region A in NGC 7714 which might still be compatible with Scharer & Vacca models. In the case of region 74C, combinations of two ionizing clusters can be found which adequately reproduce the emission line spectrum; however, this assumption would produce equivalent widths for the WR cluster higher than the maximum predicted for the models with Z=0.008, although probably compatible with solar metallicity models. Again, a finer metallicity grid for the WR cluster synthesis models would be needed in order to explore this possibility in further detail.

SUMMARY AND CONCLUSIONS
We have analyzed eight HII regions in the LINER galaxy NGC 4258 using spectrophotometric observations between 3700 and 9650Å. For two of the regions it has been possible to measure the electron temperature from the [SIII] λ 6312 A line, which allows the derivation of accurate abundances following standard methods. For the rest of the regions an empirical calibration based on the sulphur emission lines has been used to determine a mean oxygen content. The derived metallicities range from 0.3 to 0.6 Z⊙. In particular, the metallicities found for regions 74C and 5N, previously reported by Oey & Kennicutt (1993) to be close to solar, are found to be lower by a factor of two. For each observed region, we have also estimated the functional parameters: ionization parameter and the effective temperature of the ionizing cluster. Most regions show ionization parameters of the order of 10 −3 and effective temperatures of around 35300 K, except regions 74C and 69C for which higher ionization parameters are found. We have also derived the physical properties of the regions and their corresponding ionizing clusters: fillin0 factor, mass of ionized gas and mass of ionizing stars. Most of the regions, except 74C, have small ionizing clusters with masses in the range 4000 to 25000 solar masses. These values constitute in fact lower limits since the regions are assumed to be ionization bounded and the presence of dust has not been taken into account.
WR features have been detected in region 74C and, to a lower extent, in region 5NA. A detailed modelling has been carried out for region 74C using different sets of models: those of Schaerer & Vacca (1998) for WR populations and those of Leitherer et al. (1999) and García-Vargas, Bressan & Díaz (1995) for ionizing populations, to try to reproduce simultaneously both the WR features and the emission line spectrum. No consistent solution has been found: the models which better reproduce the WR features (a cluster 4.5 Myr old in the models of Schaerer & Vacc8 and Leitherer et al.) result too hard to reproduce the emission line spectrum. Conversely, a cluster 5.3 Myr old in the models of García-Vargas, Bressan & Díaz (1995) adequately reproduces the emission line spectrum, but produces too few WR stars to explain the observed WR features.
When comparing the WR "bump" and HeII intensities and equivalent widths with other well studied HII regions of different metallicities, the data on region 74C are found to fall into the observed ranges. However, the expected de-crease of WR feauture intensities with decreasing metallicity is not clearly observed. On the other hand a decrease in the effective temperature of the ionizing radiation with increasing metallicity is apparent. This seems to imply that, while the number ratio of WR to O stars is similar for regions of metallicities differing by a factor of 3, the WR stars are cooler for the regions with higher metallicity.
Both more observations of confirmed high metallicity regions and a finer metallicity grid for the evolutionary synthesis models are needed in order to understand the ionizing populations of HII regions.