A unique Galactic planetary nebula with a [WN] central star

Mon. Not. R. Astron. Soc. 346, 719–730 (2003)

A unique Galactic planetary nebula with a [WN] central star D. H. Morgan,1  Q. A. Parker2,3 and Martin Cohen4 1 Institute

for Astronomy, University of Edinburgh, Royal Observatory, Blackford Hill, Edinburgh EH9 3HJ of Physics, Macquarie University, NSW 2109, Australia 3 Anglo-Australian Observatory, PO Box 296, Epping, NSW 1710, Australia 4 Radio Astronomy Laboratory, 601 Campbell Hall, University of California, Berkeley, CA 94720, USA 2 Department

Accepted 2003 August 12. Received 2003 August 9; in original form 2003 January 30

ABSTRACT

We report the discovery of the first probable Galactic [WN] central star of a planetary nebula (CSPN). The planetary nebula candidate was found during our systematic scans of the AAO/UKST Hα Survey of the Milky Way. Subsequent confirmatory spectroscopy of the nebula and central star reveals the remarkable nature of this object. The nebular spectrum shows emission lines with large expansion velocities exceeding 150 km s−1 , suggesting that perhaps the object is not a conventional planetary nebula. The central star itself is very red and is identified as being of the [WN] class, which makes it unique in the Galaxy. A large body of supplementary observational data supports the hypothesis that this object is indeed a planetary nebula and not a Population I Wolf–Rayet star with a ring nebula. Key words: stars: individual: PM5 – stars: Wolf–Rayet – planetary nebulae: individual: PHR 16196−4913.

1 INTRODUCTION The high-resolution AAO/UKST Hα Survey of the Milky Way (Parker, Phillipps & Morgan 1999) is providing a rich source for new Galactic planetary nebulae (PNe) (Parker et al. 2003b). A systematic visual search of ∼80 per cent of the survey material has so far yielded ∼1000 new PNe; when complete, it is expected to more than double the known population of Galactic PNe as recorded by Acker, Marcout & Ochsenbein (1996). This preliminary list is detailed in version 1.0 of the Macquarie/AAO/Strasbourg Catalogue of Galactic Planetary Nebulae, available on CD-ROM (Parker et al. 2003a). Obvious candidate central stars of planetary nebulae (CSPNs) can be seen in
10 per cent of these PNe. So far, we have observed and confirmed about 700 of these PNe spectroscopically and have discovered seven with central stars that show Wolf–Rayet (WR) emission features. This is the third paper in a series on these WR central stars. The first two papers (Morgan, Parker & Russeil 2001, hereafter Paper I; Parker & Morgan 2003, hereafter Paper II) described the spectra of six central stars and their associated nebulae. Four of these (PMR1, PM3, PM6 and PM7) are [WC4] or possibly [WO4] stars, while the other two, PMR2 and PM4, are [WC9–10] and [WC8] stars, respectively. The discovery and identification of these six stars raised the number of known WR CSPNs in the Galaxy from 56 (Jeffery et al. 1996) to 62. Almost all the stars are of the late [WC] or early [WC]/[WO] types, with no confirmed intermediates. Until now, there was just one known CSPN belonging to the [WN]  E-mail: [email protected]  C

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sequence (Pe˜na et al. 1995) – Henize N66 (Henize 1956) – but N66 resides in the Large Magellanic Cloud (LMC) and is peculiar in other respects. There are also 34 CSPN identified as having very weak ˚ (Tylenda, Acker emission lines including N III lines near 4640 A & Stenholm 1993). Unlike our seven new WR CSPNs, we have so far only identified one new weak emission-line (WEL) CSPN from our PN spectroscopic follow-up observations (Parker & Morgan, in preparation). Two of the early-type [WC] CSPNs identified in Papers I and II (PMR1 and PM3) were seen to have a remarkably strong and unusual ˚ as well as the normal carbon and oxygen emission feature at 4610 A lines. This feature, unknown in any of the previously known [WC] stars, was attributed to N V. Also, one or perhaps both the cooler [WC] CSPNs (PMR2 and possibly PM4) appear to exhibit weak ˚ lines. Until this series of papers, nitrogen lines have N III 4634 A been rarely seen in [WC] CSPNs exhibiting WR characteristics. The presence of unusually strong N V or weak N III in three or four out of six WR CSPNs in Papers I and II came as a surprise and may be due to nitrogen enrichment in their stellar winds and outer atmospheres indicative of unusual stellar evolution. At this stage, the only common factor linking our newly identified PNe and the surprising discovery of strong nitrogen lines in some [WC] CSPNs is the common survey material used and the low surface brightness of the surrounding PNe (which is why they have evaded discovery until the advent of the AAO/UKST Hα Survey). There are undoubtedly further, as yet unknown, selection effects which render this new WR CSPN sample sensitive to cases of nitrogen enrichment. The seventh WR central star (PM5) was also discovered during follow-up spectroscopy of newly identified PN candidates from the

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AAO/UKST Hα Survey. PM5 turns out to be a remarkable object and, because of its unique nature, is the subject of this separate paper. It appears to be a WR star of the nitrogen sequence surrounded by an almost perfectly defined circular nebula with an extremely high expansion velocity. Since Population I WN stars can have associated ring nebulae around them, it has sometimes proved difficult to distinguish between these and classical PNe (e.g. WR124 ≡ M167). Consequently, the possibility that PM5 is a young massive star is discussed in detail after the observations of both central star and associated nebula have been described. This discussion is based on a large body of supplementary observational data, most of which is available in the public data bases. We are able to present compelling evidence which suggests that PM5 is indeed a [WN] CSPN. 2 S P E C T R O S C O P I C O B S E RVAT I O N S Spectra of the star PM5 and its associated external nebula (PHR 16196−4913) were first observed with the SAAO 1.9-m telescope: two 600-s exposures were obtained on 2000 June 20 and one 1200-s exposure was taken on 2000 June 25. The spectra cover the range ˚ at an instrumental resolution of 7.4 A ˚ pixel−1 . Further 3800–7300 A observations were made on 2000 July 2 using the Double Beam Spectrograph (DBS) on the MSSSO 2.3-m telescope, consisting of ˚ at 1200-s exposures in the ranges 6085–7050 and 4050–5960 A ˚ pixel−1 respectively. The instrumental resolutions of ∼1 and 2.2 A setup was as described in Paper I. The IRAF package was used to reduce and analyse the spectra. Relative spectrophotometry was possible through having used a wide slit in good seeing conditions and was achieved through the observation of selected standard stars (Stone & Baldwin 1983). 3 O B S E RVAT I O N A L DATA 3.1 Identification of PM5 The planetary nebula around PM5 is PHR 16196−4913 (Parker et al. 2003a). Details are given in Table 1. The J2000 coordinates are those of the central star as measured from the SuperCOSMOS Sky Survey (SSS) R-band image (Hambly et al. 2001a) and are accurate to ∼1 arcsec. According to the Simbad data base and the latest PN Catalogue Supplement (Acker et al. 1996), there is no known optically identified PN or WR star at this location. 3.2 Hα images

Figure 1. Identification chart for PM5 showing the surrounding nebula. The image is from the original Hα Survey discovery exposure HA18885. The chart dimensions are 2 × 2 arcmin2 ; north is to the top and east is to the left.

than the surrounding sky. There is no sign of any outer filament or fainter outer ring beyond this ring. Fig. 2 shows the radial profiles of the PM5 image from two separate UKST Hα exposures taken 24 yr apart. It was constructed using SuperCOSMOS pixel data extracted from scans of both the current Hα Survey film and an earlier-epoch pre-discovery plate (see later). The counts are scaled to match the maximum counts in the nebular ring and are (uncalibrated) relative pixel intensities azimuthally summed in annuli about the central star. The points plotted in each curve are median values, which were derived after excluding counts attributed to other stars in the field and counts that deviated by more than two sigma from the mean (in each annulus). The smoothness of the plotted curves shows that the system is genuinely circular. The central star, bright ring and central zone are clearly seen. A good working diameter for the nebula is 32 arcsec; the maximum

Fig. 1 shows PM5 as it appears on the Hα Survey discovery exposure; it can be seen as a central star surrounded by a circular nebula comprising a bright circular ring and a fainter central zone brighter Table 1. Details of the planetary nebula around PM5. RA (J2000) Dec. (J2000) l, b Diameter (outer) Diameter (bright ring) Hα Survey film Hα Survey field IRAS source 4850 MHz source PN designation PHR PN designation

16h 19m 40.s 16 −49◦ 13 59. 7 333.93◦ , +0.69◦ 32 arcsec 26 arcsec HA18885 349 IRAS 16159−4906 PMN 1619−4913 PN G333.9+0.7 PHR 16196−4913

Notes. The PHR designation follows Parker et al. (2003a). The radio source is from Wright et al. (1994).

Figure 2. Radial intensity distribution of the PM5 system from SuperCOSMOS pixel data (see text). The dotted line is from the Hα image of Fig. 1 from UKST Hα Survey exposure HA18885 and the dashed line is from a much earlier UKST Hα exposure taken in 1976, plate HA2416.  C

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A unique Galactic [WN] CSPN

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Figure 3. Part of the spectral image of the PM5 and its surrounding nebula, obtained with the DBS on the MSSSO 2.3-m telescope. The spectral dispersion direction lies along the horizontal axis with wavelength increasing to the right as indicated; the spatial dispersion along the slit lies along the vertical axis. The high-resolution 1200R grating was used.

diameter is 36 arcsec and the diameter of the brightest part of the ring is 26 arcsec. Comparing Fig. 2 with profiles expected from spherical shells with differing values of r i /r o (inner to outer radii) does not give a perfect match. The gradient beyond the maximum is best fitted with r i /r o ∼ 0.75, but the inner zone is too bright for this and is better fitted with r i /r o ∼ 0.5. This difference could be explained by the existence of excited gas within the ring or by inhomogeneities throughout the shell. A discussion of the evidence of nebular expansion indicated in Fig. 2 is given in Section 4.5.

3.3 Optical spectrum of the nebula Part of the charge-coupled device (CCD) spectral image of PM5 obtained with grating 1200R through the red arm of the DBS on the MSSSO 2.3-m telescope is shown in Fig. 3. Wavelength increases from left to right and position along the slit changes in the vertical direction. The bright horizontal band is the broad He II 6560 emission line in the spectrum of the WR central star, and the three ellipses are due to [N II] 6548, Hα and [N II] 6584 emitted by the nebula. The ellipses are near-perfect representations of a uniformly expanding spherical gas shell. [N II] 6584 is much stronger than Hα, at a level usually associated with nebulae of the Type I class. The most significant feature of Fig. 3 is the very large expansion velocity evident for the three strong nebular emission lines. Table 2 gives the details of the velocity measurements and line ratios at several points on the spectral image. The three lines give the same expansion velocity, which has a mean of v exp = 167 km s−1 . This is very much larger than normal PN expansion velocities of 15– Table 2. Nebular emission lines observed in PM5. Line [N II] 6548 Hα [N II] 6584 He I 6678 [S II] 6717 [S II] 6731

v exp

Flux (1)

Flux (2)

Flux (3)

169.0 165.8 166.5 – – –

40.5 100.0 165.3 – – –

66.1 100.0 165.5 – – –

58.8 100.0 187.1 3.7 5.9 7.6

Notes. v exp is the expansion velocity in km s−1 . Fluxes (1–3) are fluxes relative to Hα = 100.0. Flux (1) is for the blueshifted line on the central star. Flux (2) is for the redshifted line on the central star. Flux (3) is for the combined extreme points on the slit.  C

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30 km s−1 and immediately marks PM5 out as an unusual object. This topic will be discussed in more detail in Section 4.5. These nebular emission lines are much wider than the neighbouring sky lines which can be used as a measure of the instrumental line profile. The full width at half-maximum (FWHM) of each of the six lines (three blueshifted and three redshifted) was measured in the spectra extracted from each of the four single-pixel rows immediately below the stellar spectrum after subtracting a mean sky spectrum. A suitably weighted estimate of the FWHM of the six lines is 4.0 ± 0.5 pixel. The FWHM of 10 nearby sky lines is 2.4 ± 0.3 pixel. Subtracting these in quadrature gives a velocity spread of 80 ± 13 km s−1 . When superimposing the [N II] and Hα images from Fig. 3, it is clear that the extents of the nebula in [N II] and Hα are identical along the slit, i.e. there is no obvious inner higher excitation zone. Also, when the counts in Hα are plotted against those in [N II] 6584 for identical positions on the ring (excluding the points on the star itself), the correlation is good and little poorer than that obtained when the two [N II] lines are plotted against each other. So there is no [N II]-rich or [N II]-poor region (relative to Hα zones) around the ring. The high expansion velocity allows a limit to be placed on the general level of dust inside the nebula. The ratio of the redshifted and blueshifted components of the [N II] 6548, Hα and [N II] 6584 emission lines is found to be 0.96 ± 0.28 using measurements just above, below and on the central star. Hence, there is no evidence of internal extinction in excess of a few tenths of a magnitude in R. The final point to note from the [N II] lines is that the intensity between the red- and blueshifted components falls to the level of the stellar continuum, indicating that there is no low-velocity ionized gas in the nebula. No other nebular emission lines can be seen against the spectrum of the central star, i.e. [S II] 6717,6731 is weaker than 6 per cent of Hα. No lines were detected in the blue spectrum obtained at MSSSO, and no other nebular lines were seen in the SAAO spectrum despite its more extended wavelength coverage. However, [S II] lines can be seen in two spectra extracted at the bright extremities of the nebula where there is no underlying stellar emission. Although these lines are weak and difficult to see as a result of diffuse [S II] (and Hα) across the whole CCD image, they are certainly real, being 50 and 100 per cent greater in strength in the nebula + sky spectrum than in the sky spectrum, in contrast with the pure sky lines, which are all within 20 per cent of each other. Also, the [S II] lines are 60 per

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Figure 4. Spectrum of the central star PM5 normalized to a continuum level of 1.0. The spectrum is from the SAAO 1.9-m telescope. The identification at ˚ is tentative on account of the nearby atmospheric absorption band. 7590 A

cent wider in the nebula + sky spectrum than in the sky spectrum; the pure sky lines differ by no more than 20 per cent. This larger FWHM reflects the velocity spread in the nebula. 3.4 Optical spectrum of the central star ˚ as recorded at The normalized spectrum of PM5 from 5500 to 7800 A the SAAO 1.9-m telescope is shown in Fig. 4 with the major emission lines marked. The narrow nebular emission lines seen superimposed ˚ in Fig. 3 have not been removed because they on He II 6560 A are difficult to resolve in the lower-resolution SAAO spectrum. The ˚ because PM5 SAAO spectrum records nothing shortward of 5500 A is very red; hence the blue arm of the MSSSO DBS failed to register a spectrum. The emission lines of He II and He I are strong and there is a ˚ This is almost certainly strong broad emission feature at 7100 A. the blended multiplet 4 of N IV 7103–7128 with He I 7065 and He II 7177 in its wings, which is commonly seen in WN stars (Vreux, Dennefeld & Andrillat 1983). One consequence of the large width of the He I 7065 line is the appearance of the maximum intensity at ˚ rather than 7115 A ˚ as is usual in most WN stars. If this He I 7090 A line is assumed to have the same profile as He I 6678 and is subtracted from the observed feature, the residual profile is seen to be centred ˚ There is no evidence of any of the many strong carbon near 7115 A. ˚ is not so weak that the lines seen in WC stars; the signal at 5800 A extremely strong C IV 5808 blend would be undetected. Hence, PM5 is of the [WN] sequence. The presence of the next strongest N IV ˚ is supportive of this. Details of the stellar emission line at 6219 A

Table 3. Stellar emission lines in PM5. Line

EW

FWHM

EW

FWHM

Notes

He I 5876 N IV 6219 He II 6560 He I 6678 N IV 7115 N IV 7583

16 11 182: 44 102 –

21 26 – 54 97 –

– 2 145 37 – –

– 8 59 54 – –

1 2 3 4,5

˚ Columns 2 and 3 are from SAAO, 4 Notes. The measurements are in A. and 5 from MSSSO. 1 – excludes nebular lines: the exclusion process was difficult for the SAAO data and no FWHM could be measured. 2 – includes He II 6683. 3 – includes He I 7065 and He II 7177. 4 – includes He II 7593. 5 – not measured because of strong atmospheric absorption.

lines in Fig. 4 are listed in Table 3 with equivalent width (EW) and FWHM measurements from both SAAO and MSSSO given when available. It should be noted that there are no significant differences between the three recorded SAAO CCD frames; in particular, the ˚ complex are effectively the same three components of the 7100 A ˚ in each. in each and there is a weak line at 6219 A The WR spectra plotted by Vreux et al. (1983) and Vreux et al. (1989) can be used to estimate a spectral class for PM5. The ratio of He I 7065 to He II 7177 greatly exceeds unity for types WN8 and later and is well below unity for types WN5 and earlier. This leads to a best estimate of the spectral type of PM5 as [WN6]; but [WN7] and [WN5] cannot be ruled out in view of the noise in the spectrum  C

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A unique Galactic [WN] CSPN and the star-to-star variation seen in the comparison stars. Also, the early types are ruled out by the weakness (not detected) of He II 6891 relative to He I 6678 + He II 6683. The presence of the atmospheric ˚ could introduce errors. However, the absorption band near 7200 A extracted spectrum was compared with that of a nearby star, and on dividing the first by the second, very little difference in the profile ˚ complex was found, indicating that the derived line of the 7100 A ratios can be used with reasonable confidence. Another spectral feature of PM5 is the great strength and width of the WR emission lines. Unfortunately, most work on the spectra of WN stars has involved lines in the blue and green parts of the spectrum that are too greatly extinguished in PM5. However, when the EW and FWHM of the He II 6560 line in PM5 are compared with measurements of those quantities made by D. H. Morgan on unpublished FLAIR spectra of LMC WN stars, it turns out that PM5 is typical of the WN4–6b stars as classified in the recent scheme of Smith, Shara & Moffat (1996). It should be noted that the strength ˚ line is also strong in WN6ha stars, but the width (EW) of the 6560 A (FWHM) is much smaller. Therefore, if PM5 were a Population I star, its class would definitely be WN6b rather than WN6ha. The widths of the He I 6678 + He II 6683 blend are consistent with those of the He II 6560 line.

3.5 Optical photometry of the central star Photometric measurements of PM5, compiled from several sources are listed in Table 4. The SAAO spectra were obtained on nights of good seeing (
1 arcsec) through a wide (2.4 arcsec) slit and are suitable for providing absolute fluxes. After accounting for atmospheric extinction, they were calibrated using observations taken on the same night of the spectrophotometric standard LTT 9239 (Stone & Baldwin 1983). The aim was to make measurements of the monochro˚ appromatic (continuum) magnitudes at 4270, 5160 and 6000 A priate for the b, v and r bands of the Smith (1968) system and adapted to the Hayes & Latham (1975) calibration of Vega. PM5 is so red that only the r magnitude could be measured properly. ˚ this gave an upHowever, some flux was detected at 5160 A; per (bright) limit of the v magnitude as 18.5, or an estimate of v ∼ 19.5 ± 1. Table 4. Photometry of PM5. Waveband BJ R I J H K r F(8 µm) F(12 µm) F(15 µm) F(21 µm) F(12 µm) F(25 µm) F(60 µm) F(100 µm) F(4850 MHz)

Mag./flux

Source

21.5 16.0 13.9 10.31 9.24 8.48 17.6

SSS SSS SSS 2MASS 2MASS 2MASS SAAO spectrum

0.74 4.25 4.89 10.95 4.21 19.58 33.13:
222.3 244.0

Note: the F(60 µm) IRAS flux is of low quality.  C

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MSX (Jy) MSX (Jy) MSX (Jy) MSX (Jy) IRAS (Jy) IRAS (Jy) IRAS (Jy) IRAS (Jy) PMN (mJy)

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The broad-band photometric data quoted in Table 4 were taken from the raw SuperCOSMOS Sky Survey (SSS) (Hambly et al. 2001a) measurements of the UK Schmidt Telescope Southern Sky Surveys. The SSS BJ and R measurements were calibrated with sequences from the HST Guide Star Photometric Catalogue GSPC 2.1 (Bucciarelli et al. 2001) and the I data were calibrated through a comparison of SSS and DENIS measurements of a nearby area of sky. It may be noted that the calibrated spectrum integrated over the R band gives R = 16.6, which is a little fainter than the broad-band measurement. The influence of the He II 6560 line and He II 6684 lines is to decrease the R magnitude of PM5 by 0.2 mag. Lines in the I band (mainly He II 8237) will have a lesser effect (Vreux et al. 1983). Note that it is not sensible to transform the R-band data from the natural photographic system to the Cousins system because the R band in the latter system has a long red tail which makes photometric corrections for stars as red as PM5 highly problematical. Corrections to adjust the effective wavelength of the photographic bands to those of the zero-points are very small. It is worth noting that PM5 appears in the HST GSC II Catalogue (McLean et al. 2000) with R = 15.90, but no other colours are available there. 3.6 Infrared photometry of the system Infrared flux measurements can be obtained from the IRAS and MSX catalogues available at the Infrared Science Archive.1 These too are quoted in Table 4. The IRAS photometry provides two good- and one moderate-quality measurements plus one upper limit. PM5 is not found in the IRAS catalogue of small extended sources and is therefore unresolved by IRAS. The MSX fluxes are reasonably consistent with the IRAS fluxes, suggesting that the bulk of the mid-infrared emission lies within the 20-arcsec FWHM point spread function of the MSX images. The position of PM5 is not covered by DENIS data released at the time of writing this paper. However, 2MASS All-Sky data are available for the central star. 2MASS detects this star at (J2000) RA = 16h 19m 40.s 19, Dec. = −49◦ 13 58. 96, in excellent agreement with the optical position. The JHK magnitudes, which are accurate to 0.03 mag, are given in Table 4. 3.7 Infrared images of the system The 8-µm MSX image is about 8 pixel or 48 arcsec across; point sources appear to be about 6 pixel across. Thus, PM5 is a little extended in the mid-infrared. As well as being extended, the MSX image of PM5 is more flat-topped than stellar images. Assuming that the source of emission is a ring like the Hα emission, profiles across the centre of the image are best-fitted by a ring of diameter ∼22 arcsec. A ring of size 26 arcsec (the Hα diameter) would produce a broader 8-µm image in the bright parts. Rings smaller than 18 arcsec across generate profiles that are too narrow. The emission that gives rise to the extended overall dimension is relatively faint. That is, PM5 appears to be slightly smaller at 8 µm than at Hα at bright levels but with faint emission outside. However, PM5 is some 50 per cent larger at 21 µm than at 8 µm (to the lowest credible contour levels). Close inspection of the very faint background at 8 µm shows PM5 to be located on some faint emission. It appears to be on the edge of a faint ring of diameter 3–4 arcmin; but since there is also extensive low-level emission to the south-east, it is not obvious that these features are actually associated with PM5. None are seen in 1

http://iras.ipac.caltech.edu/

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D. H. Morgan, Q. A. Parker and M. Cohen 4 DISCUSSION PM5 is a WN star that appears to be the central star of a planetary nebula. This would make it the only Galactic PN with a [WN] CSPN. Since there are now 62 other Galactic CSPNs of the [WC] type (including those discovered in our first two papers), and theories suggest that it is difficult to form [WN] stars, it is important to check and review all the available data to see whether they support this hypothesis. In particular, it is necessary to consider seriously the possibility that PM5 is a Population I star surrounded by a ring nebula. As noted earlier, it has sometimes proved difficult to distinguish between these and classical PNe. 4.1 Central star spectrum

Figure 5. The average, absolutely calibrated, IRAS LRS spectrum of PM5. Note the deep 10-µm absorption, probably due to interstellar silicates, and the [Ne II] line at 12.8 µm.

the other bands, but that is probably because the background noise is greater in these bands. 3.8 Infrared spectrum of the system There are four independent spectra of PM5 from the complete pointsource-extracted IRAS Low Resolution Spectrometer (LRS) data base. We have followed the paper by Cohen, Walker & Witteborn (1992) to: (i) assign individual uncertainties to the fluxes at each wavelength for each LRS spectrum; (ii) merge the higher-resolution blue and lower-resolution red segments; (iii) rectify the incorrect shapes of all archival LRS spectra; and (iv) recalibrate each spliced, complete, rectified spectral shape absolutely using the IRAS average 12-µm flux density. Fig. 5 presents the averaged LRS data. Although the spectrum is undoubtedly very noisy, several interesting aspects are still apparent. There is a deep absorption feature near 10–11 µm, most likely due to silicates, confirming that this object suffers high extinction. However, the LRS data are too noisy to afford a formal determination of the interstellar AV that produces this absorption. The spectrum peaks near 19 µm, suggestive of a cool dust emission component. Three of the four individual LRS spectra show the 12.8-µm [Ne II] line in emission (the fourth is very noisy there, and was not used in the LRS average). The absence of the 10.5-µm [S IV] and 15.5-µm [Ne III] lines characterizes the infrared spectrum as one of low excitation.

It is the spectrum of the central star that is the main reason for considering the possibility that PM5 is a Population I object. Not only is the spectral class typical of a Population I WN star rather than a CSPN, but the lines are as strong and broad as those seen in a massive star. However, the argument is not conclusive. Although most [WC] CSPNs have weaker lines than their early-type counterparts, this is not always the case; e.g. NGC 6751 and PMR1 (Paper I and references therein). N66, the only previously known [WN] CSPN in the LMC, is another case in point. It too has broad lines: the EW and ˚ reFWHM measurements for the He II 6560 line are 95 and 70 A ˚ recorded spectively (Pe˜na et al. 1995), compared with 160 and 60 A for PM5. Thus, the linewidth (FWHM) is about the same in the two stars and the linestrength (EW) is not much stronger in PM5 by the standards of EW variation seen in WR stars. Similarly, the FWHMs of He I 6678 + He II 6683 in PM5 and N66 are quite close (55 and ˚ respectively). (The EW of He I 6678 + He II 6683 is smaller in 40 A ˚ compared with 40 A) ˚ as expected because N66 is classified N66 (8 A as [WN4.5] whereas PM5 is [WN6].) It is interesting to note that Smith et al. (1996) suggest that the broad-line spectra belong to WR stars with low initial masses (for Population I), perhaps implying that [WN] stars would have broad lines. 4.2 Location In Paper II it was seen that several of the newly discovered [WC] stars had to be PNe rather than Population I objects because, if the latter, they would be too high above the Galactic plane. PM5, however, is so close to the Galactic plane that these arguments do not help. PM5 lies in a region that is relatively free from obvious diffuse Hα nebulosity as seen in the AAO/UKST Hα Survey (Fig. 6). The

Figure 6. A 900 arcmin2 region around PM5 as seen on the Hα Survey. PM5 and another nearby newly discovered PN are marked on the image. Note that the region is generally free from obvious Galactic Hα nebulosity.  C

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A unique Galactic [WN] CSPN spectral data do reveal a faint Hα ‘wash’ over the region, which is also evident in the low-resolution but sensitive SHASSA Hα Survey (Gaustad et al. 2001). [The sensitivity limit of the SHASSA Survey is ∼2 Rayleigh with 0.5 Rayleigh smoothing; preliminary investigations suggest that the AAO/UKST Hα Survey is sensitive to ∼5 Rayleigh (Parker et al., in preparation)]. PM5 lies mid-way between two large H II regions, the edges of which are more than one degree away. Curiously enough, there is another newly discovered PN just 7 arcmin north of PM5; it too is marked on Fig. 6 (Parker et al. 2003a). PHR 1619−4906, with a diameter of ∼45 arcsec, is a little larger than PM5, but is much fainter in Hα and has no discernible central star. Thus, the general immediate environment contains other Population II objects but shows no evidence of Population I stars in an H II region.

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WR ring nebulae: ∼0.22 (0.14–0.28) in the M33 nebulae (Esteban et al. 1994), and ∼0.1 in M1-67 (Esteban et al. 1991). Thirdly, there is the ratio of the [S II] lines – namely [S II] 6717/[S II] 6731 – which in PM5 is 0.78, giving a moderate electron density of n e = 6000 cm−3 for an assumed electron temperature of 104 K (Lynch & Kafatos 1991). The minimum density as derived from the nebula + sky spectrum is 3000 cm−3 ; the pure sky spectrum corresponds to the low-density limit as expected. In contrast, the [S II] 6717/[S II] 6731 line ratio is ∼1.3–1.8 in the M33 ring nebulae (Esteban et al. 1994) and is 1.2–1.3 in WR40 and NGC 3199 (Kwitter 1984), which corresponds to ne ∼ 200 cm−3 (Mathis et al. 1992). Yet again, M1-67 is an exception with n e = 1000 cm−3 (Esteban et al. 1991). So, the electron densities in PM5 are typical of PNe and about 30 times greater than in ring nebulae (but just six times that in M1-67).

4.3 Morphology PM5, with its almost perfect ring structure, looks like a PN. None of the Population I ring nebulae found in the surveys by Marston, Chu & Garcia-Segura (1994a), Marston et al. (1994b) and Marston (1997) look like PM5; they are either amorphous, complex or faint. Moreover, most are found in regions of complex large-scale nebulosity indicative of a youthful environment – PM5 is not (see above). PM5 could perhaps be compared to WR16 and WR40, which are described in more detail by Marston (1995), but again the differences are significant. WR16 is round but, rather than having a strong rim, it has a line of ejecta around it; WR40 appears as a clumpy ellipse of ejecta. Of the ‘small’ ring nebulae in the LMC studied by Chu, Weis & Garnett (1999), that around Br13 is clear and round, but is much fainter relative to its central star than is the nebula around PM5. Note also that the much discussed nebula M1-67 is not like PM5: though circular, it consists of many small bright knots seen right across its disc. What is observed in PM5 is more akin to the product of a single mass-loss event rather than the sweeping up of remnants of the local interstellar cloud. The key morphological point is that PM5 does not resemble Population I ring nebulae. The MSX images described earlier are also like those of PNe. For example, there is great similarity between PM5 and PM4 (Cohen & Parker 2003), which has a [WC8] central star (Paper II). The faint 8-µm emission around PM5 is seen in other MSX images of PNe and is unlike the bright 8-µm emission seen around compact H II regions (Cohen & Parker 2003).

4.5 Nebular expansion Although PM5 was identified on the AAO/UKST Hα Survey film, it can also be seen on an older UKST Hα plate taken in 1976 some 24 yr before the discovery film. This older plate was taken on the coarse-grained 098 emulsion using the AAO 10-inch Hα filter and so produces a much poorer image than the one shown in Fig. 1, which was from the fine-grained Tech-Pan emulsion (e.g. Parker & Malin 1999). It also suffers a little from having been taken in poorer seeing conditions than those for the Hα Survey film. Nevertheless, the older exposure of PM5 was scanned by SuperCOSMOS and treated in the same way as the Hα Survey image, though with some smoothing to reduce the grain noise; its radial profile is shown in Fig. 2 alongside the radial profile from the Tech-Pan survey image. The former was scaled to give similar count levels in the background and in the bright ring as found in the recent image. The central star appears much brighter relative to the nebula in the old data because the small-format Hα filter used for the older image has a considerably ˚ than the Hα Survey filter (70 A). ˚ Despite wider passband (100 A) the poorer quality of the older image, the plots show the ring to be about 1 arcsec smaller in radius in 1976 than in 2000. This is consistent with the expansion of 0.9 arcsec expected in 24 yr from a velocity v exp = 165 km s−1 in a nebula at a distance of 1 kpc. An alternative way of looking at the data is to enlarge the radial profile from the 1976 plate according to the measured v exp for a set of assumed distances and to compare these with the profile from the

4.4 Nebular spectrum Three useful line ratios have been measured for PM5. The first is [N II] 6584/Hα, which, at 1.87, is characteristic of Type I PNe, which are high-excitation usually bipolar objects. In contrast, this ratio is much lower in most ring nebulae: it is ∼0.20 (0.10–0.22) in seven WR nebulae in M33 (Esteban, Vilchez & Smith 1994), and is 0.8, 0.5 and 0.22 in WR40, S308 and NGC 3199 respectively (Kwitter 1984). But M1-67 is again an exception with [N II] 6584/Hα ∼ 1.1 (Esteban et al. 1991). Nevertheless, even in M1-67, [N II] 6584/Hα does not approach the level seen in PM5. Although most, if not all, PNe with strong [N II] 6584 are indeed high-excitation objects (Morgan & Parker 1998), it seems that this may not be the case for PM5; PM5 is not bipolar, does not exhibit high-excitation lines in its LRS spectrum, and does not have a particularly hot exciting star. It is not unreasonable to suppose that PM5 has a high nitrogen nebular abundance as a result of the relatively high nitrogen content of its stellar wind. The second ratio is [S II] 6717+6731/Hα, which is 0.14. This too is typical of PNe (Morgan & Parker 1998); but it is also common in  C

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Figure 7. Parts of the radial profiles of the nebula from Fig. 2. The profile for 2000 is the solid line and that for 1976 is the dashed line. The series of dotted lines are the 1976 profile shifted according to the measured v exp for distances 0.5, 1.0, 1.5, 2.0, 2.5 and 15.0 kpc, with the shortest distance being the farthest to the right.

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2000 film. This is shown in Fig. 7. The best fit is for a distance of 1 kpc. A good fit to a nebula at a distance of 0.5 or 1.5 kpc would require v exp = 80 or 250 km s−1 respectively. These are well beyond the error limits of the velocities measured from the spectra, but it not yet clear how errors in the surface photometry will distort the result. In simplistic terms, a distance less than 0.5 kpc or greater than 1.5 kpc seems to be ruled out. Expansion velocities for PNe with [WC] central stars can be as much as double those of typical PNe (Medina & Pe˜na 2002) and are typically in the range 20–50 km s−1 (Pe˜na, Stasinska & Medina 2001). Ring nebulae have expansion velocities in the wide range 15– 110 km s−1 (Chu et al. 1999). All these are well below the expansion velocity seen in PM5. However, very high expansion velocities are seen in some PNe, including N66, which shows a double-peaked structure in [O III] 5007 with a velocity separation of 82 km s−1 . With the nebula considered to be bipolar oriented at an angle of 60◦ to the line of sight (Dopita et al. 1993), the expansion velocity becomes 136 km s−1 , which is not so far below the 165 km s−1 recorded for PM5. The main difference between the two nebulae is the morphology: N66 is bipolar, PM5 is apparently spherical. Other PNe with very high expansion velocities are (Dopita, Ford & Webster 1985): He2-111, which is a bipolar Type I PN with velocities ∼400 km s−1 but in a very faint outer bipolar shell; NGC 6302, which has velocities of several hundred km s−1 and is a bipolar PN too; and NGC 2392, which has an expansion velocity of 75 km s−1 in a more spherical system. The expansion velocity, if uniform, leads to a dynamic age for the nebula of 373(d/kpc) yr. If PM5 is a PN at 1 kpc, then an age of 373 yr places PM5 amongst the youngest PNe (Dopita et al. 1987). Ages for WR ring nebulae are very much larger, 2.5 × 104 to 4.3 × 105 yr (Chu et al. 1999). Even if PM5 were at the very large distance of 10 kpc, its dynamic age of 0.4 × 104 yr would still be much smaller than normal for ring nebulae. Luminous blue variable (LBV) nebulae are younger by a factor of 10 (Chu et al. 1999). 4.6 Extinction It is clear from its colours that PM5 suffers a high level of extinction. The extent of this can be estimated from the broad-band colours, from the slope of the calibrated spectrum and from the ratio of the radio flux to the Hα flux. It is also useful to look at the extinction seen in stars in the vicinity of PM5. The broad-band colours of PM5 are not easy to use because the star is very faint in the blue and is located in a crowded part of the sky. Moreover, broad-band colours are not really suitable for such a red spectrum. The best photometric colour from Table 4 is (R − I). Although the (R − I)0 colours of WR stars are not known, continuum colours can be calculated on the basis of a power-law spectrum (F λ ∼ λ−α ), which has been shown to be valid for WR stars by Morris et al. (1993). In particular, α lies between 2.5 and 3.5 for WN6 stars. Using the zero-points of Bessell (1979), this gives (R − I)0 in the range −0.25 to 0.0. Applying the mean Galactic extinction law (Nandy et al. 1975) leads to a visual colour excess, E(B − V), in the range 3.1 to 3.8. The second approach is to use the calibrated SAAO spectrum. This too is not straightforward because the spectrum is so seriously truncated by the reddening. Nevertheless, the spectrum redward of ˚ was dereddened until it gave a satisfactory fit to an intrin5500 A sic power-law spectrum with α between 2.5 and 3.5. Acceptable solutions were found to lie in the range E(B − V) ∼ 3.0–3.5. Finally, the extinction was calculated from the ratio of the radio continuum and Hα fluxes using the theoretical Case B ratio

(Osterbrock 1974, – equations 4.22 and 4.30 and table 4.4). Electron temperature is an unknown parameter but electron density is unimportant. A helium fraction of 10 per cent was assumed. The 4850-MHz flux from PM5 is 244 mJy (Wright et al. 1994). The beam size of the radio survey was 4.2 arcmin, so PM5 appears as a point source. The resolution of the MSSSO DBS data is good enough to allow a straightforward separation of the nebular and stellar spectra. The observed Hα flux was measured from the MSSSO data by making a spectral extraction over the whole nebula (as seen perpendicular to the dispersion direction). The stellar contribution, which here is the top of the broad stellar Hα emission line, was subtracted from this spectrum and the Hα flux determined by comparing the MSSSO data with the calibrated SAAO spectrum. The flux from the entire nebula was then calculated knowing that the detected flux came from a slit 2.0 arcsec wide and from a nebula of radius 15 arcsec. It was done for three models which assumed that all the flux came from a shell with r i /r o = 0.0, 5.0 and 10.0 arcsec. The results for the three models give E(B − V) in the ranges 3.1–3.3, 2.9–3.1 and 2.7–2.9 for log(T e ) = 3.7, 4.0 and 4.3 respectively. Extinction measurements for stars in this part of the Galactic plane at distances in the ∼1–2 kpc range have E(B − V) ∼ 1 (Neckel & Klare 1980), so it is expected that the extinction to PM5 is at least at this level. The Population I WN stars nearest in direction to PM5 are Nos 74 and 75 in the VIIth Catalogue of Galactic Wolf–Rayet Stars (van der Hucht 2001), which are ∼2–3◦ away from PM5. They have high extinction levels equivalent to E(B − V ) = 2.2 and 1.1 mag at distances of 2.6 and 2.2 kpc respectively. Data for stars within a 1 deg field around PM5 were extracted from the SSS and calibrated as described in Section 3.5. Most of the stars in the (R − I) versus (BJ − R) diagram form a clump displaced redwards by the equivalent of E(B − V ) = 1.0 from its normal position (Hambly, Irwin & MacGillivray 2001b). Thus, there is uniform extinction of this level right across the field. In addition, there is a tail of stars extending along the reddening line with PM5 at its extremity. Most of these would be reproduced if a similar (normal Galactic plane) population were reddened by an additional layer (or layers) of dust beyond the stars of the clump with E(B − V) ∼ 1.0–1.5. This would give a total E(B − V) ∼ 2.0–2.5 for these redder stars. Few stars are as red as PM5. No group of stars in any photometric diagram are found in positional groups across this field of view. Taken together, all these considerations lead to the conclusion that E(B − V) for PM5 most likely falls in the range 3.0 to 3.5, i.e. AV = 9.3–11. 4.7 Photometric distance Distances can be calculated once the extinction is known and an absolute magnitude is adopted. The latter is not obvious for PM5. According to the VIIth Catalogue of Galactic Wolf–Rayet Stars (van der Hucht 2001), a normal Population I WN6 star has Mv = −4.1 and a WN6ha star has Mv = −7.0 (though of course PM5 is not WN6ha – see Section 3.4). However, Mv = −4.1 is somewhat fainter than values derived for LMC stars, which do have accurate distances and relatively small levels of extinction (Vacca & TorresDodgen 1990). This discrepancy could have arisen because most of the stars used to derive the Mv = −4.1 level are binaries for which the contribution from the WR component had to be estimated. Consequently, a value of Mv = −4.7 has been adopted; this was based on the single stars available in the VIIth Catalogue, is more consistent with the LMC data and falls mid-way between the Mv = −4.1 and −5.4 values quoted for WN5 and WN7 stars respectively.  C

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−4.7 and v ∼ 19, its photometric distance would be 2.5–5 kpc as E(B − V) changes from 3.5 to 3.0 mag. The IRAS fluxes can be used to estimate the distance if the object is a PN. The data of Tajitsu & Shin’ichi (1998) show a relationship between the distance to a PN and the sum of the IRAS fluxes at 25 and 60 µm for specified ranges of the ratio of the two bands (temperature). This was used in Paper II. Superposing the data for PM5 on this relationship gives a distance of 1.4 ± 0.2 kpc, which is consistent with the distance calculated from the expansion. 4.8 Nebular size

Figure 8. Absolute magnitude of PM5 as a function of extinction E(B − V) for four distances marked (in pc) on the plots. The apparent magnitude of PM5 was taken to be v = 19.5 (solid lines); the dashed lines show the effect of changing this by ±1.0 mag. The four arrows mark the brightest absolute magnitude possible for each distance, corresponding to the maximum extinction [E(B − V ) = 3.5] and brightest v magnitude (v = 18.5).

Using 30 arcsec as the measured angular diameter of the nebula, the distance leads immediately to the nebular radius. At a distance of 1.0 kpc, the radius is 0.07 pc, which is well within the 0.01–0.3 pc range associated with PNe. Population I ring nebulae are much larger than this, mostly ∼5 pc (Chu et al. 1999), though one in the LMC is as small as 2 pc and some are smaller still: WR43 – 0.7 pc, WR42d – 1.0 pc, WR77 – 1.3 pc, WR93 – 1.9 pc (Marston 1997). Even if PM5 were at a distance of 5 kpc, its radius would be just 0.4 pc, which is still more appropriate for a PN than a ring nebula. Although most ring nebulae are large, any caused by the ejection of material from the star would have been smaller in the past, so size does not rule out the possibility that PM5 is a very young ring nebula. 4.9 Nebular mass

It should be remembered that there is a significant spread in Mv from star to star within any subtype. With PM5 having a spectral subclass error of one subclass, its absolute magnitude could be Mv = −4.7 ± 0.7. There are, of course, no precedents for providing the absolute magnitude of a [WN6] central star. WR central stars are usually much fainter than their Population I equivalents; but N66, which is of class [WN4.5], has M V = −2.23 (Pe˜na et al. 1995) (equivalent to Mv = −2.33, see van der Hucht 2001) and is not much fainter than a Population I WN4.5 star at Mv = −3.8. PM5 could be intrinsically brighter than N66 because it is 1–2 subclasses later than N66. Given the uncertainties in both the absolute magnitude and the extinction, it is useful to calculate Mv as a function of E(B − V) for several adopted distances. Although the v magnitude cannot be ˚ is too accurately measured because the SAAO spectrum at 5160 A noisy, the signal is consistent with v = 19.5 ± 1. The calculated values of Mv are shown in Fig. 8 for v = 19.5 and for distances of 0.5, 1.0, 1.5 and 2.0 kpc; the first and third of these distances are the extreme values allowed from the measurement of the expansion (Section 4.5), and the fourth is included for comparison purposes. The plots also show the effect of changing v by ±1.0 mag. Four arrows are shown in the figure; these mark the brightest absolute magnitude possible for each distance, corresponding to the maximum extinction [E(B − V ) = 3.5] and brightest v magnitude (V = 18.5). The faintest absolute magnitudes possible, corresponding to the minimum extinction [E(B − V ) = 3.0] and faintest v magnitude (V = 20.5) are Mv > 0 for the two upper quadrants and Mv > −1 for the lower pair. It is clear that the available photometry for PM5 is consistent with E(B − V ) = 3.0–3.5 and d = 0.5–1.5 kpc as proposed above, for stars with Mv between 0 and −4, which covers what is expected if PM5 is a PN. The Population I option is a much poorer fit but cannot be completely ruled out if Mv , v, E(B − V) and d are all near their limits, i.e. v ∼ 18, Mv = −4.1, E(B − V ) ∼ 3.0 and d = 1.5 kpc. If the distance were not constrained by the expansion measurement and PM5 were a Population I star with Mv =  C

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The nebular mass follows from the radius, the filling factor and the mean density. One estimate of the density has been obtained from the nebular spectrum; a second can be derived from the free–free radio emission. Although there is only a single radio measurement, this affords an independent estimate of the nebular mean density, assuming that the 4.85-GHz flux density arises from optically thin free–free emission. For the nebula, we insert Z i = 2, n i = 0.5n e (Cohen, Barlow & Kuhi 1975) and T e = 104 K into equation (3-54) of Spitzer (1978) for radio free–free emissivity. The Gaunt factor was calculated as ∼5.5 at this frequency, whence ne ≈ 4800 cm−3 results, using PM5’s optical radius as the linear scale of the emission region So, at 1 kpc, M neb for PM5 is
0.17 M , which is the mass obtained for a density n e = 5000 cm−3 and r i /r o = 0.5 (filling factor 0.875) assuming that this density is uniform over the entire shell. This is entirely reasonable for a moderately dense PN (Meatheringham, Dopita & Morgan 1988). In comparison, the controversial ring nebula M1-67 has a filling factor of just 0.008 because it consists of small condensations located in a thin shell. Even so, M neb for M1-67 is 0.8 M (Solf & Carsenty 1982). If PM5 were a ring nebula with the same nebular mass as M1-67, then it would be at the reasonable distance of 1.7 kpc but would be extremely small, with a radius of 0.12 pc. It is worth noting that the measured (and extrapolated) Hα flux used in Section 4.6, when dereddened by AV = 9.3–11 and moved from a distance of 1.0–1.5 kpc to the 50 kpc of the LMC, is equivalent to a log(Fβ) flux of −12.25 to −13.14, which is typical of the LMC PNe observed by Meatheringham et al. (1988). 4.10 Infrared colours Walker et al. (1989) describe the distribution of a variety of stellar and non-stellar IRAS point sources in the [12] − [25] versus [25] − [60] colour–colour plane (see their table 1 and fig. 1). PM5 has the IRAS colours of [12] − [25] = 3.26 ± 0.10 and [25] − [60] =

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2.52 ± 0.16, which are matched solely by the ‘PNR’ population (red PNe; their redness derives from cool dust grains within the nebulae). The IRAS colours of WR ring nebulae (e.g. Mathis et al. 1992) are different, more like those of normal H II regions (Pottasch et al. 1984). The [25] − [60] colours do not distinguish PM5 from the colours of the samples of all Population I WRs (characterized by Cohen 1995), or all dusty WC stars. However, the [12] − [25] colours of Population I WRs are much bluer than those of PM5, whether one examines all WRs, or all dusty late [WC]. Thus, PM5 has the IRAS colours of neither a dusty Population I WR, nor a WR ring nebula, but it does match those of a PN. We have determined the MSX colour indices for the 87 categories of celestial object in ‘SKY’, the model of the point source sky that simulates observed source counts with high fidelity at wavelengths from the ultraviolet to 35 µm (Cohen 1993, 1994; Cohen, Sasseen & Bowyer 1994). SKY’s infrared wavelength flexibility derives from the empirical spectral library of 87 average spectra that underpins its application to new passbands. Using MSX’s letter designations (A, C, D, E) for the four long-wavelength bands (with isophotal wavelengths of 8.3, 12.1, 14.6 and 21.3 µm respectively), we find the observed colours of PM5 to be: [A]−[C] = 2.75 ± 0.08, [A]−[D] = 3.31 ± 0.08 and [A]−[E] = 4.98 ± 0.10. All magnitudes and uncertainties come from the MSX Point Source Catalog Version 1.2. The most reliable of these colours are those involving band A, the most sensitive MSX band. The closest matches to these colours within SKY’s representation of the MSX sky are: [A]−[C] PNR, [A]−[D] PNR, and [A]−[E] PNR. PM5 clearly has the IRAS and MSX colours of a red PN. 4.11 Infrared spectrum The occurrence of the peak of the LRS spectrum at 19 µm is as expected for a PNR source (see fig. 2 of Wainscoat et al. 1992). If we assign PM5’s central star to Population I WN status, then what would be the observable consequences for the mid-infrared (MIR) spectrum? The MIR survey of WR stars by Smith & Houck (2001) offers 8–13 µm spectra at a resolving power of ≈600 of 29 WC and WN stars, including several dusty WC and heavily reddened WN stars. With the exception of the WC9s with the most circumstellar emission, the dominant emission feature in all these WRs is the 9.7-µm He II (13 → 11) line, which is seen even in the somewhat dusty WC7 star, HD 192641. By types WN7 and 8, the blends of He I and He II lines near 11.3 and 12.4 µm dominate, even in WR145 and 147, which suffer interstellar extinctions of 8–11 magnitudes. Of particular relevance to PM5, Smith & Houck (2001) also present spectroscopy of WR124, a Population I WN8 star embedded in a stellar wind-blown (non-planetary: Crawford & Barlow 1991) nebula (M1-67). In WR124, these same lines are present but the 12.8-µm nebular [Ne II] line rivals the 12.4-µm blend. If PM5 contained a Population I WN6 star and dust made only a modest contribution to the MIR continuum, one might reasonably expect to detect at least one of the 9.7, 11.3 and 12.4 µm features; yet its LRS shows only [Ne II]. If dust continuum emission were to dominate the photosphere and wind emission in PM5, the absence of stellar emission lines would not preclude a Population I WN star, because the line-to-continuum contrast could have been significantly reduced. However, the relative brightness of the central star of PM5 might be difficult to explain in the presence of so much circumstellar obscuring material. Despite the noisy LRS spectrum, the absence of any discernible permitted lines from either photosphere or stellar wind is also completely consistent with a PN spectrum, like the LRS spectra of dusty red PNe such as NGC 40 and IC 418 (Volk & Cohen

1990), which interestingly are PNe with [WC] central stars, and the much fainter IC 2149 from the MIR spectral survey of dusty PNe by Casassus et al. (2001). It can also be noted that the observed JHK colours are consistent with those seen for other PNe (Garcia-Lario et al. 1997). 4.12 Infrared flux Infrared emission is also seen in ring nebulae around Population I WR stars. Mathis et al. (1992) have studied the three best WN ring nebulae for dust and have shown extended dust within the nebulae, preferentially located in a ring near the outer shells. For these stars, the log(λF λ ) versus log(λ) plot of the stellar data shows the expected steady linear decrease to the red, and, even after adding the full IRAS flux, still shows a large fall from R to 25 µm; but the IRAS flux for PM5 is the same as the dereddened R and I fluxes. That is, PM5 emits far more infrared flux relative to the stellar flux than do any of the three WN nebulae studied by Mathis et al. (1992). It was the high level of infrared flux that led van der Hucht et al. (1985) to revive the idea that the nebula M1-67 (PN G050.1+03.3) is a planetary nebula. The ring nebula theory, however, was supported by Esteban et al. (1991) on the basis of extinction levels in and around the nebula, and by Crawford & Barlow (1991) from measurements of the velocity extent of interstellar absorption lines and their consistency with a large distance for the star. Hence, there is at least one ring nebula with a significant infrared flux. Given that the alternative classifications of PM5 are a Population I WN star and a PNR, one might be able to distinguish between these possibilities using their IRAS absolute magnitudes. Cohen (1995) derives the average IRAS absolute magnitudes of a variety of Population I WR stars. If we treat PM5 as a single WN star, it should have −7.21, −9.08 and −12.18 at 12, 25 and 60 µm respectively. However, its obviously dusty nature suggests that the magnitudes of a dusty WC star would be more relevant; these would be −9.67, −10.51 and −13.95 at 12, 25 and 60 µm. (Wainscoat et al. 1992) give −8.35 and −12.06 at 12 and 25 µm for a PNR. Comparing these with PM5’s observed IRAS magnitudes gives an average distance for a single WN star of 0.6 kpc, for a dusty WC star of 1.6 kpc and for a PNR of 1.5 kpc. At only 0.6 kpc, PM5 would be grossly underluminous for a normal WN star, but a value around 1.6 kpc would be equally consistent with a dusty WC star as with our other estimates in this paper for a PN. Evidently, we are unable to distinguish between these two possibilities on the basis of the MIR absolute magnitudes. 4.13 Bolometric luminosity An estimate of the bolometric magnitude of the star alone can be derived by attaching a scaled version of the intrinsic expected spectrum for a Population I WN6 star (kindly provided by Dr P. Crowther: Crowther & Vacca, in preparation) to our dereddened R and I points, and using this to extrapolate the overall energy distribution into the ultraviolet. This method yields log(L/L ) ∼ 4.3–4.7 for E(B − V ) = 3.0–3.5 and a distance of 1 kpc. At 1.5 kpc, log(L/L ) ∼ 4.6–5.1. Strong-lined early Population I WN stars have log(L/L )  4.9 (Hamann, Koesterke & Wessolowski 1993). Although typical PN [WC] central stars have log(L/L )
4.0, some can be as bright as log(L/L ) = 4.5 (Leuenhagen, Hamann & Jeffery 1996). Importantly, the [WN] central star of N66 has log(L/L )
4.5 (Pe˜na et al. 1997). Thus, the model is consistent with PM5 being a PN if d = 1.0 kpc, but is ambiguous if d = 1.5 kpc. However, although the model used matches a power-law index (α) of 2.85, which is a  C

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A unique Galactic [WN] CSPN good working mean (Morris et al. 1993), α can vary by ±0.5 for any given spectral subclass (and does not correlate with subclass) and does not depend on absolute magnitude. On altering the calculations to account for a different continuum slope, it is found that for α = 3.35 a Population I WR is preferred at most distances >1.0 kpc (though a PN is valid at 1.0 kpc); but at α = 2.35, a PN is preferred at all distances
1.5 kpc. (A Population I star is ruled out for all models at d = 0.5 kpc.) It should be noted that the photometrically derived extinction values determined in Section 4.6 depend on α (the intrinsic colour), in the sense that a low α will give a low extinction and a lower luminosity (as well as a lower luminosity because it has less far-ultraviolet flux). Some reasons for preferring the lower values of α and extinction are: (1) the extinction is below the average when estimated using the methods that did not use an intrinsic colour; (2) the Hα flux is lower than is expected from the observed r magnitude; and (3) the IRAS spectrum implies a low ionization and hence a lower far-ultraviolet flux. 4.14 Variability The central star of N66 is known to have gained its WN characteristics through a dramatic mass-loss event around 1990 (Pe˜na et al. 1997). Its spectrum changed over a period of a few years from being featureless in 1990 into a [WN4.5] spectrum in 1994, along with an initial brightening of both ultraviolet and optical continua corresponding to a drop in T eff from ∼120 000 K to ∼50 000 K and an increase in brightness from M V = +1.24 (1987) to M V = −2.23 (1994). Other spectral variations in N66 have occurred on timescales of years and months (Pe˜na 1995), but the nebular lines have changed little. Until similar monitoring of PM5 is undertaken, information on the variability of the CSPN can only come from legacy exposures such as those in the UKST Plate Library. A preliminary examination of these has not revealed any significant (m  1 mag) changes in PM5 over the past 25 yr or so. 4.15 Radial velocity The radial velocities of the three lines [N II] 6548, Hα and [N II] 6584 in the MSSSO nebular spectrum (albeit broadened) are 8.7, 4.1 and 6.4 km s−1 respectively. The mean is 6.4 km s−1 , which is −3 km s−1 with respect to the local standard of rest. The radial velocities of PNe in this direction are in the range −20 to +40 km s−1 (Osterbrock 1974), so PM5 is consistent with this. It is helpful to note that the Galactic rotation model of Fich, Blitz & Stark (1989) with the Galactocentric distance to the Sun set at 8.5 kpc predicts radial velocities close to zero at distances out to ∼1 kpc, but becoming increasingly negative at greater distances. This means that the measured low radial velocity is consistent with the PN hypothesis (small distance) but less so with the alternative WR ring nebula proposition for which the photometry predicts distances of 2.5–5 kpc (unless the star has v = 18, which is unlikely). Even at 2.5 kpc the model predicts a radial velocity of −30 km s−1 . However, individual stars do not necessarily conform to the model; e.g. M1-67 has an abnormal radial velocity. 4.16 Summary PM5 appears to be a planetary nebula rather than the Population I WN star plus ring nebula that is expected from its WN-type spectrum. Its nebular spectrum (optical and infrared), density, infrared flux, radio flux, size, radial velocity and morphology all fit the PN hypothesis. Distances derived from its optical photometry, radio  C

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flux, MSX and IRAS fluxes are consistent with the 1 kpc derived from the observed radial expansion. The combined data presented here are not consistent with identification of PM5 as a Population I ring nebula unless the star is especially faint and unique in its other characteristics. 5 CONCLUSIONS One of the PNe newly discovered on the AAO/UKST Hα Survey (PM5) has a [WN6] central star. If confirmed, PM5 would be the only Galactic PN known to have a [WN] central star. Its unique status, coupled with the fact that the stellar emission lines are as strong as those in Population I WNb stars, raises the possibility that it might be a ring nebula surrounding a Population I WN6b star. All the available optical, infrared and radio photometric data support the hypothesis that PM5 is a PN at a distance of 1.0–1.5 kpc and is heavily reddened with AV ∼ 9.3–11, but they cannot rule out the possibility that PM5 is a particularly faint Population I WN star. However, the other data – location, morphology, nebular parameters (density, size and mass), MIR spectrum and colours – are all typical of a PN and not at all appropriate for a Population I star. PM5 is also unusual in that its circularly symmetric nebular shell has a very high expansion velocity, v exp ∼ 165 km s−1 . Best estimates of other nebular parameters are: ne ∼ 5000 cm−3 , radius ∼0.07–1.0 pc and M neb
0.17 M . The nebula appears to have a high nitrogen abundance, which is likely to be associated with the nitrogen content of the stellar wind. It is important to obtain more observations of PM5 to study its unique character in more detail. AC K N OW L E D G M E N T S The authors are grateful to the staff of the UK Schmidt Telescope Unit for taking the films of the AAO/UKST Hα Survey and to the Directors and staffs of the SAAO and MSSSO for the awards of telescope time and assistance at the telescopes. This publication makes use of data products from the Two-Micron All Sky Survey, which is a joint project of the University of Massachusetts and the Infrared Processing and Analysis Centre/California Institute of Technology, funded by the National Aeronautics and Space Administration and National Science Foundation. MC thanks NASA for supporting his participation in this work through LTSA grant NAG5-7936 with UC Berkeley. REFERENCES Acker A., Marcout J., Ochsenbein F., 1996, First Supplement to the Strasbourg–ESO Catalogue of Galactic Planetary Nebulae. Observatoire de Strasbourg Bessell M. W., 1979, PASP, 91, 589 Bucciarelli B. et al., 2001, A&A, 368, 335 Casassus S., Roche P. F., Aitken D. K., Smith C. H., 2001, MNRAS, 320, 424 Chu Y.-H., Weis K., Garnett D. R., 1999, AJ, 117, 1433 Cohen M., 1993, AJ, 105, 1860 Cohen M., 1994, AJ, 107, 582 Cohen M., 1995, ApJS, 100, 413 Cohen M., Parker Q. A., 2003, in Kwok S., Dopita M. A., Sutherland R., eds, Proc. IAU Symp. 209, Planetary Nebulae: Their Evolution and Role in the Universe. Astron. Soc. Pac., San Francisco, in press, p. 33 Cohen M., Barlow M. J., Kuhi L. V., 1975, A&A, 40, 291 Cohen M., Walker R. G., Witteborn F. C., 1992, AJ, 104, 2030 Cohen M., Sasseen T., Bowyer S., 1994, ApJ, 427, 848 Crawford I. A., Barlow M. J., 1991, A&A, 249, 518 Dopita M. A., Ford H. C., Webster B. L., 1985, ApJ, 297, 593

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2003 RAS, MNRAS 346, 719–730