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Astron. Astrophys. 347, 556-564 (1999)

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2. Observations and results

Imaging observations of the region of the X-ray error box of RX J0045.4+4206 (see Fig. 1) were performed during three nights in October 19-21, 1996 with the 1.5 m telescope of the Observatorio Astronómico Nacional (OAN) de San Pedro Mártir (SPM), México, equipped with a 1024[FORMULA]1024 pixel CCD TEK. The total field of view was 4:03[FORMULA]4:03. The seeing was typically [FORMULA], thus we observed with the CCD chip binned to 2[FORMULA]2 and with a corresponding pixel scale of [FORMULA]/pixel. Images were obtained in various filter bands (see Table 1, listing central wavelength, FWHM and peak transmission) to allow a diagnostic of the excitation level. The data reduction used standard MIDAS procedures for bias and cosmic ray subtraction and flat-fielding.


[TABLE]

Table 1. Details on narrow-band filters used


Spectroscopy was performed on two occasions in August 1996 and August 1997, respectively, using the 2.1 m telescope of OAN SPM. The Boller & Chivens spectrograph with a 600 l/mm (1996) and 400 l/mm (1997) grating was employed to get spectra with overall FWHM resolution of 4.5 Å and 6-7 Å. The seeing was typically [FORMULA], and consequently a [FORMULA] slit has been used. The data reduction again used standard MIDAS procedures of the long-slit spectrum reduction package.

Finally, we also observed the nebula with the UNAM Fabry-Perot interferometer PUMA. The Fabry-Perot observations of the nebula BA 1-642 were carried out during the night of November 2, 1997 at the f/7.5 Cassegrain focus of the 2.1 m telescope of the OAN SPM using the UNAM Scanning Fabry-Perot Interferometer PUMA (Rosado et al. 1995).

A 1024[FORMULA]1024 thinned Tektronix CCD detector, with an image scale of 0.59 arcsec pixel-1, was used with a 2[FORMULA]2 on-chip binning in both dimensions. Thus, the resulting image format was 512[FORMULA]512 pixels, with a spatial resolution of 1.18 arcsec pixel -1.

An interference filter centered at [FORMULA] Å and having a bandpass of 20 Å was used in order to isolate the H[FORMULA] line. The scanning Fabry-Perot interferometer is an ET-50 of Queensgate Instruments with a servo-stabilization system. The main characteristics of this interferometer are: interference order of 330, free spectral range of 18.92 Å (equivalent to a velocity range of 908 km s-1) and sampling spectral resolution of 0.41 Å (equivalent to 18.9 km s-1) at the wavelength of the H[FORMULA] line, achieved by scanning the interferometer gap at 48 positions. Thus, the resulting data cubes have dimensions of 512[FORMULA]512[FORMULA]48.

Under these conditions, we have obtained one nebular data cube with an exposure time of 48 min. We also obtained two calibration data cubes spaced at the beginning and at the end of the observations in order to check for possible flexures of the equipment. For the calibration cubes we have used the Ne line at [FORMULA] 6598.95 Å.

The data reduction and analysis were performed using the specific reduction package CIGALE (Le Coarer et al. 1993). This software was used to remove the cosmic rays, to carry out the wavelength calibration of the data cubes, to obtain continuum subtracted [FORMULA]-cubes and to carry out the emission line profile analysis.

2.1. The central star MLA 1159

The spectrum of the central star, MLA 1159 (Fig. 2), shows strong and broad emission lines at 4650 Å and 5805 Å (which we interpret as CIII and CIV ), and clearly suggest a carbon Wolf-Rayet star thus confirming the Wolf-Rayet candidacy proposed by Meyssonier et al. (1993). The ratios of CIII 5696 / CIV 5805 [FORMULA] 1 and CIII 5696 / OV 5592 [FORMULA] 1 identify MLA 1159 without much doubt as WC6 subtype (according to Smith 1968, van der Hucht 1981). The recently proposed refinement of the WC classification scheme (Smith et al. 1990, further developed by Crowther et al. 1998) does not alter this result. With a FWHM of the CIV 5805 Å line of [FORMULA] Å and [FORMULA], we again arrive at a WC6 subtype (see Table 3 in Crowther et al. 1998).

[FIGURE] Fig. 2. Mean spectrum of the star MLA 1159. The broad, strong emission lines of CIII , CIV 4650 Å and CIV 5805 Å clearly identify this as a C-type Wolf-Rayet star. Weak CIV 4442 Å and CIII 5696 Å (see inset) emission is also visible.

With a measured [FORMULA] = 19.2 mag, using a distance of 650 kpc and assuming only galactic foreground absorption and no intrinsic M 31 extinction, i.e. [FORMULA]=0.6 mag, we derive [FORMULA] = -5.7[FORMULA]0.3 mag, at the bright end of the range of absolute magnitudes of WC5-6 stars of [FORMULA] = -3.9[FORMULA]0.5 mag (Lundström & Stenholm 1984).

We derive a position based on the digitized sky survey of RA. (2000) = [FORMULA], Decl. (2000) = [FORMULA] (error of [FORMULA]) which is consistent with the position of MLA 1159 (Meyssonier et al. 1993).

2.2. The nebula BA 1-642 [FORMULA] PAV78 915

2.2.1. F-P interferometry

Fig. 3 shows the PUMA [FORMULA]-maps corresponding to the velocity channels where the nebula BA 1-642 is detected. We see that it covers a velocity range from Vhelio = -157 km s-1 to -24 km s-1 having maximum intensity at Vhelio [FORMULA] -100 km s-1. It is interesting to note that other M31 nebulae are also detected in the 10 arcmin field of the PUMA at the same velocities but with a narrower range in velocity. This suggests that BA 1-642 has internal motions larger than the internal motions of classical HII regions.

[FIGURE] Fig. 3. 1[FORMULA]8[FORMULA]1[FORMULA]5 H[FORMULA] images ([FORMULA]-maps) of the nebula obtained in different velocity channels with the F-P interferometer PUMA. The other PUMA channels (0-39) do not show appreciable nebular emission. The channel numbers and heliocentric velocities are marked at the lower left corners of each image. The images are continuum subtracted.

Fig. 4 is an enlargement of one of the velocity channels (at Vhelio = -100 km s-1). The WR star MLA 1159 is located within a nearly circular nebula (see also Figs. 1, 5). The south-east rim of the nebula is very pronounced, particularly in the H[FORMULA] line.

[FIGURE] Fig. 4. The enlarged image of the nebular complex in the H[FORMULA] filter, in the velocity channel 43, corresponding to [FORMULA] km s-1. Note the large circular shell centered approximately on the northern edge of the compact bubble.

However, from this figure we also note, that there is some extended structure on the West and the total complex probably consists of two different sets of nebulosities: a nearly hemispherical internal bubble with an angular size of 34[FORMULA]26 arcsec corresponding to a linear size of 107[FORMULA]82 pc at the adopted distance of 650 kpc to M31 (this is BA 1-642 [FORMULA] PAV78 915) and a V-shaped nebulosity to the West of the internal bubble (note that this part has its own nomenclature, i.e. BA 1-641 [FORMULA] PAV78 913). A close inspection of Fig. 4 suggests that this second nebulosity seems to be a thick shell of 74 arcsec in diameter (equivalent to 233 pc) centered at the Northern edge of the internal bubble. The external shell is brighter near the MLA 1159 which is probably photoionizing this nebulosity. The fact that both nebulosities have their maximum intensities in the same velocity channel implies that they would be at the same galactocentric distance and consequently a physical link between them is favored against a chance superposition along the line of sight. The thick shell could be a trace of the photoionization or the interaction with the interstellar medium of the winds of the progenitor of the WR star or other massive stars. Deeper observations are required in order to confirm this hypothesis. In the following we will restrict our kinematical study to the internal bubble.

We have obtained radial velocity profiles integrated in several zones of the bubble, the brightest region of the shell and other HII regions of the field. We confirm that the shell has the same velocity as the borders of the bubble: Vhelio [FORMULA] -100 km s-1. Even if our velocity profiles have low S/N ratio (about 4) we are able to detect two heliocentric velocity components for the central region of the bubble: at -83 and -131 km s-1 with an uncertainty of about 4 km s-1. On the other hand, the velocity profiles at the edges of the bubble are simple and can be fitted by a single Gaussian function. This could be interpreted as an expansion motion with a velocity of 25 km s-1. Consequently, we can estimate the kinematic age of the nebula using tkin = 0.6 R/V where tkin is in units of 106 yr, R in pc and V in km s-1 (taken from Weaver et al. 1977 models for the evolution of wind blown bubbles, see also Rosado 1986). Taking R = 54 pc and V = 25 km s-1, we obtain a kinematic age of 2.1 106 yr. This value of the kinematic age is larger than the duration of the WR phase. Thus, this bubble could not be formed by the wind of MLA 1159 (i.e. W-type, Chu 1981, Rosado 1986). Instead, this bubble should be formed by the WR star progenitor and is only illuminated by the WR star (i.e. Rs-type, Chu 1981, Rosado 1986). This is typical for the nebulae associated with WC-type stars.

2.2.2. Narrow-band photometry and spectrophotometry

The shape and the size of the nebula is very similar at different wavelengths (see Fig. 5 with images of the nebula in the six bands HeII , H[FORMULA], [OIII] , H[FORMULA], [NII] and [SII] ). The nebula shows strong Balmer emission lines, [OII] and [SII] but weak [OIII] (see Fig. 6). We note that the emission lines are all blue-shifted by 3-4 Å, consistent with the above reported velocities.

[FIGURE] Fig. 5. The nebula BA 1-642 in 6 different filter bands: HeII (top left ; 1200 sec exposure time), H[FORMULA] (top right; 2400 sec), [OIII] (middle left ; 2400 sec), H[FORMULA] wide (middle right ; 1800 sec), [NII] (bottom left ; 1200 sec), [SII] (bottom right ; 1200 sec). All images are centered on the Wolf-Rayet star which is particularly bright in the HeII image due to the strong carbon emission (see Fig. 2). Images with exposure times larger than 1200 sec are sums of two individual exposures. The slit positions and extraction areas of the long-slit spectroscopy are drawn on top of the H[FORMULA] image.

[FIGURE] Fig. 6. Mean spectrum of the nebula BA 1-642. Note that H[FORMULA] is cut at the top because the ordinate has been expanded to show also the weak emission lines.

It is not completely clear whether the ring-like nebula is physically center filled or not, because it shows less emission in the inner parts. Nevertheless, we have measured the size for both the inner and the outer ring based on the pronounced, bright south-east rim visible in the H[FORMULA] image to [FORMULA] and [FORMULA], respectively. At the distance of M 31 this corresponds to radii of 34 and 50 pc, respectively.

Using the ratio of [OIII] 5007 Å/H[FORMULA] as criterion, the excitation class is estimated to be 0.5 according to the system of Feast (1968) and Webster (1975). This and the lack of nebular HeII emission lines indicate that BA 1-642 is a low-excitation nebula.

The integrated H[FORMULA] line flux in regions Ib and IIb is 3[FORMULA]10-15 and 2[FORMULA]10-15 erg cm-2 s-1. The analysis of the H[FORMULA] filter image shows that region Ib corresponds to the mean H[FORMULA] flux integrated over the full nebula. We therefore scaled the integrated H[FORMULA] line flux from region Ib by the fractional area of the extraction region to derive a total H[FORMULA] line luminosity of the nebula of 4[FORMULA]1036 erg s-1. This H[FORMULA] line luminosity is a few percent of the luminosity of the Wolf-Rayet star, and thus can be produced by the irradiation of the central Wolf-Rayet star.

2.2.3. Photoionization modelling

In the following, we derive some constraints from the emission line ratios measured at various locations of the nebula (see Table 2 and Fig. 5). We note that all measured Balmer decrements are consistent with the recombination value or only very small reddening, certainly lower than the global total galactic column density in this direction of [FORMULA] cm-2 (Dickey & Lockman 1990) which corresponds to [FORMULA] mag. We therefore have not applied a reddening correction.


[TABLE]

Table 2. Line intensity ratios relative to H[FORMULA]. The values for regions a and b are the mean of the sectors I and II, while the values for region III is the mean of IIIa and IIIb (see Fig. 5 for the relative locations of the different regions). Only the spectra taken of region III cover the sulphur lines.


First, we estimated some of the parameters using diagnostic emission line ratios:

Density via sulphur ratio : The observed ratio of the two sulphur lines can be used to estimate the density of the emission-line gas. The measured intensity ratio [SII] [FORMULA]6716/6731 = 1.4 is close to the low-density limit and implies a density of [FORMULA] cm-3 (Osterbrock 1989), a value typical for HII regions. In the photoionization calculations described below, the gas density was fixed to this value or varied within a small range of this value (within a factor [FORMULA]2).

Temperature via oxygen ratio : For the inferred gas density, the oxygen intensity ratio [OIII] [FORMULA]5007/4363 is a useful estimator of the gas temperature. Given that we only have an upper limit for [OIII] [FORMULA]4363, we derive a corresponding upper limit for the gas temperature of [FORMULA] K (Osterbrock 1989). We note that this is not very strict, but again consistent with what is known for HII regions.

Second, we have performed photoionization calculations with the code Cloudy (Ferland 1993). We assumed the gas clouds to be dust-free (except if stated otherwise) and illuminated by the continuum of the central Wolf-Rayet star. The continuum shape was approximated by a blackbody, its temperature left free to vary (since there is no one-to-one match in temperature of a blackbody and a Wolf-Rayet atmosphere). Solar gas abundances (Grevesse & Anders 1989) were adopted if not stated otherwise. The gas was assumed to be ionization bounded. The aim of these photoionization calculations is to reproduce the observed line-ratios, get clues on the ionizing continuum shape (within the limits of the blackbody assumption), decide whether we see a filled volume or shell, derive various properties of the nebula (like ionization parameter, temperature, thickness), constrain abundances, and check for the presence of dust.

We have calculated a large grid of photoionization models varying the ionizing continuum (i.e., [FORMULA]) and the radius (equivalent to the ionization parameter). Test calculations were also performed for different density, metallicity and the presence of dust. The ionization parameter is defined as

[EQUATION]

where Q is the number rate of photons above 13.6 eV, r is the distance between the Wolf-Rayet star and the emission-line gas, [FORMULA] is the hydrogen density and c is the speed of light.

A main observational feature is the relative weakness of [OIII]/H[FORMULA] (though well within what is observed in `normal HII regions'; cf. Fig. 7). Whereas close distances and high temperatures ([FORMULA] K) significantly overpredict the strength of [OIII] (cf. Fig. 8), it is of the order of the observed value for larger distances of the illuminated gas (i.e., lower ionization parameters). Further, we note that the quite substantial change in strengths of some emission lines, particularly [OIII] , over a distance [FORMULA]1 pc to [FORMULA]40pc serves as argument that we are only seeing a shell of gas, as is also suggested by the visual impression.

[FIGURE] Fig. 7. Location of our WR nebula (filled square) in diagnostic diagrams, as compared to a few others (open symbols). The region enclosed in the dashed lines marks typical stellar ionization as observed in HII regions and HII galaxies (note, that the left borderline just reflects current observation limits, since weak lines might have escaped detection). The thick line is the dividing line (Osterbrock 1989) as compared to other types of ionization (like a hard continuum in AGN). Our WR-nebula is located well within the region populated by HII regions. For comparison, we plot some PNs surrounding WRs using the sample of Pena et al. 1997 (note: biased to low values of [OIII] since in several objects of their sample the [OIII]-line was saturated, thus not plotted.) Results from photoionization models are plotted in these diagrams as well: The lines correspond to several model sequences. Along a line, [FORMULA] varies between 30.000 and 60.000 K in steps of 5000 K from bottom to top. Each line corresponds to a fixed ionization parameter (log U = -2.4 (blue), log U = -3.2 (red)). The dotted (green) line corresponds to a sequence were dust was mixed with the gas and the metal abundances were depleted (log U = -2.4 as for the `blue' model). The model with log U=-3.2 and [FORMULA]=55000 K best matches all observed line ratios (Table 3).

[FIGURE] Fig. 8. Run of the [OIII]/H[FORMULA] emission line ratio in dependence of distance r (in cm) from the WR star and blackbody temperature (= 3, 4, 4.5, 5, 5.5, 6 [FORMULA] 104 K from bottom to top ) for [FORMULA]. The observed range for the emission line ratio (corresponding to different slit positions) is bracketed by the dotted horizontal lines.

We further find that low-ionization lines are underpredicted for small distances (less than a few 1019 cm), i.e., large ionization parameters, consistent with general expectations (e.g. Stasinska 1982, Komossa & Schulz 1997).

In order to simultaneously match the observed strengths of low- and high-ionization lines we require a temperature of [FORMULA] K (see Table 3) and an ionization parameter of [FORMULA] corresponding to a distance of [FORMULA] cm (cf. model (i) of Table 3) for an ionizing (beyond the Lyman limit) luminosity of [FORMULA] erg s-1 and [FORMULA] K; or a slightly larger distance if n is lowered (cf. model (ii) of Table 3). We also carefully checked that no other, non-observed emission lines are overpredicted by the model.


[TABLE]

Table 3. Results of the photoionization calculations carried out with the code Cloudy . The first column gives observed values of emission lines (all relative to H[FORMULA]; Ib representatively chosen and [SII] added from region III), the next columns give the model results. Model (i) was calculated for [FORMULA] K, [FORMULA] and [FORMULA]; model (ii) for the same [FORMULA], [FORMULA] and [FORMULA] (see text for details). [FORMULA] is the gas temperature at the illuminated phase of the shell in K. We note that an even better match of individual lines could be obtained by further fine-tuning r or [FORMULA]. Given the uncertainties, particularly the approximation in representing the WR-star atmosphere by a blackbody, this was not attempted.
Notes:
(1) [FORMULA], [FORMULA]; [FORMULA].


This model (we used model (i) in the following; see Table 3) then predicts (for a shell of 4[FORMULA] covering) an H[FORMULA]-luminosity of [FORMULA] = 3[FORMULA]1037 erg s-1 which overpredicts the observed value, i.e. implies a filling factor of less than unity, of the order of 0.1.

Given that both, [NII] and [SII] are correctly predicted by this model, we do not find evidence for a deviation of the Nitrogen abundance from the solar value.

Re-calculating the best-fit model, now including Galactic-ISM-like dust (the species graphite and silicate) but non-depleted (i.e., still solar) abundances, has only weak influence on the resulting emission line spectrum (dust slightly contributes to the heating, but not much, in the present case). However, if additionally the abundances are set to the ISM value (a mean of Cowie & Songaila 1986, as included in Cloudy ) the effect is quite strong. Due to the depletion of important coolants, several line-ratios increase in intensity (cf. Fig. 7). The [SII] line becomes particularly strong, which can be partly traced back to the overabundance of sulphur, which happens to be twice the solar value in the employed set of ISM metal abundances. Under the assumptions made, i.e. dust properties and ISM metal abundances, the observed emission lines are better matched for a dust-free environment (or selectively depleted S abundance).

A mean nebular temperature of [FORMULA] K is derived for the best-fit (dust-free) model.

The total mass of the nebula amounts to [FORMULA] [FORMULA] assuming a covering factor of 0.1 (see above).

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© European Southern Observatory (ESO) 1999

Online publication: June 30, 1999
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