ACTA ASTRONOMICA Vol. 43 (1993) pp. 389±396

On the Evolutionary Status of WR-type Nuclei

by

R. Tylenda and S.K. G o r n y

Copernicus Astronomical Center, Chopina 12/18, PL-87100 ToruÂn, Poland

ABSTRACT

The planetary nebula nuclei showing WR-type spectra constitute a unique class of the central stars: they are He-burners. This work presents preliminary results of a study analyzing the observational characteristics of the WR-type nuclei and their nebulae. The existing He-burning models cannot account for the observations of the WR-type nuclei. In the discussion we consider two scenaria: (i) WR-type nuclei are single stars; (ii) WR phenomenon is an evolutionary phase of binary systems.

1. Introduction

The planetary nebula nuclei (PNN) showing Wolf-Rayet features in their spectra constitue a unique class of the central stars. This class includes, at present, about 50 objects and all of them are of WC-type (Tylenda et al. 1993). Their strong winds imply that these stars have active shell sources. This and the fact that their atmospheres are hydrogen-poor (Mendez 1991) imply that the WR-type PNN are burning helium. It is, thus, the only class of PNN for which we know the nuclear burning mode. We have undertaken an extensive study of the WR-type PNN in order to better understand their evolutionary status. This paper presents some preliminary results.

2. Observed Samples

From the list of WR-type PNN in Tylenda et al. (1993) we have selected a sample of 30 objects for which we have found necessary observational data (e.g.

PNN magnitudes, H ¯uxes, nebular diameters) in the Strasbourg-ESO catalogue (Acker et al. 1992). Similarly from the list of H-rich PNN of Mendez (1991) we have selected another sample of 38 objects. An important part of the reported study is devoted to a comparison between these two samples. The names of the objects included in the two samples are given in Table 1. 390 A. A.

Table1 Planetary nebulae included in the H-rich sample and in the WR-type sample.

H-rich WR-type PK name PK name 9±5.1 NGC 6629 0±1.6 M 2-20 25+40.1 IC 4593 2+5.1 NGC 6369 34+11.1 NGC 6572 2±9.1 Cn 1-5 37±34.1 NGC 7009 3+2.1 Hb 4 43+37.1 NGC 6210 4+4.1 M 1-25 45±4.1 NGC 6804 11+4.1 M 1-32 54±12.1 NGC 6891 17±4.1 M 1-30 80±3.1 NGC 6853 29±5.1 NGC 6751 63+13.1 NGC 6720 61±9.1 NGC 6905 64+48.1 NGC 6058 64+5.1 BD+30  3639 66±28.1 NGC 7094 89+0.1 NGC 7026 83+12.1 NGC 6826 120+9.1 NGC 40 93+5.2 NGC 7008 130+1.1 IC 1747 96+29.1 NGC 6543 144+6.1 NGC 1501 123+34.1 IC 3568 146+7.1 M 4-18 166+10.1 IC 2149 161±14.1 IC 2003 197+17.1 NGC 2392 189+19.1 NGC 2371-7 206±40.1 NGC 1535 243±1.1 NGC 2452 215±24.1 IC 418 278+5.1 PB 6 220±53.1 NGC 1360 278±5.1 NGC 2867 261+32.1 NGC 3242 286+2.1 He 2-55 264±12.1 He 2-5 307±3.1 NGC 5189 285±14.1 IC 2448 309±4.1 He 2-99 292+4.1 PB 8 309±4.2 NGC 5315 294+43.1 NGC 4361 321+3.1 He 2-113 307±4.1 MyCn 18 327±2.1 He 2-142 315±13.1 He 2-131 332±9.1 He 3-1333 316+8.1 He 2-108 336±6.1 PC 14 320±9.1 He 2-138 337+1.1 Pe 1-7 325±12.1 He 2-182 358±21.1 IC 1297 326±6.1 He 2-151 327+10.1 NGC 5882 329+2.1 Sp 1 331±3.1 He 2-162 345+0.1 IC 4637 345±8.1 Tc 1 350+4.1 H 2-1 358±0.2 M 1-26

3. Results

From the observational data we have calculated the blackbody Zanstra tem-

T ( ) F ( ) T ( ) Z peratures: Z HI from the H to PNN continuum ¯ux ratio, and HeII Vol. 43 391 if the HeII 4686 AÊ line is observed in the nebular spectrum. Then we have

Ê T = T ( ) T = T ( )

Z Z Z adopted Z HI if HeII 4686 A is not seen, and HeII if

Ê T ( ) T ( ) Z

Z HeII HI (this latter conditionis satis®edinall caseswhenHeII 4686 A T

is observable). Fig. 1 plots Z against the WC subclass for the WR sample. As expected, the correlation between the two parameters is clear.

Fig. 1. The Zanstra temperature against the WC subclass for the WR sample. T

In Fig. 2 we display the histograms of Z for both samples. As can be seen the T

H-rich sample tends to have somewhat lower Z than the WR sample. This is most

probably due to observational selection. It is easier to identify a H-rich PNN at low T

e when the Balmer lines are expected to be well seen in the spectrum. Note that

the histogram for the WR sample does not show any signi®cant de®ciency of PNN T

having middle Z as could have been expected from the observed lack of the WC 5 ± 8 subclasses (Tylenda et al. 1993).

The ®rst difference between the samples which we have noted from the obser- V

vational data concerns the nebular expansion velocities, exp . For the WR sample V

the mean exp is 31.0 km/s with a standard deviation of 9.0 km/s. For the H-rich

sample the ®gures are 18.7 km/s and 11.0 km/s, respectively.

V T Z Fig. 3 shows the relation between exp and . Full circles represent the

WR-type PNN while the open circles shows the objects from the H-rich sample. V

Both samples follow a similar trend in the sense that exp tends to increase with

T V exp the increasing Z . The obtained difference in the mean value of is partly due

to selection effects in the samples. As can be seen from Fig. 3 objects having low

T V exp

Z and measured are underpopulated in the WR sample. Nevertheless, a part V

of the difference in mean exp seems to be real. It can be seen that independently

T V exp of Z the WR-type PNN tend to situate at higher than the H-rich PNN.

392 A. A. T

Fig. 2. Top: The histogram of T for the H-rich sample. Bottom: The histogram of for the

Z Z WR-type sample. Vol. 43 393

Fig. 3. The relation between the PN expansion velocity and the PNN Zanstra temperature. Full

circles ± WR-type PNN; open circles ± H-rich PNN. f Fig. 4. TZ against the parameter . Full circles ± WR-type PNN. Open circles ± H-rich PNN.

Dashed curves ± theoretical H-burning tracks.

T f

Fig. 4 plots Z against the parameter de®ned as

2

F V

V exp

= ( )

f 2 1

F V

where V is the PNN ¯ux in the band and is the observed nebular radius. 394 A. A.

This sort of diagrams, for the ®rst time used by Tylenda and Stasinska (1989), can be used for a comparison between observed PNN and theoretical models. Its great advantage is that it is independent of distances. Full symbols in Fig. 4 represent

the WR-type PNN whereas the H-rich objects are shown as open symbols. Double V

arrows mark objects for which exp is not known and we have adopted the mean

V f

exp for the WR sample while calculating . Dashed curves show the theoretical tracks of SchonbernerÈ (1981, 1983) and BlockerÈ and SchonbernerÈ (1991). These tracks are shown for orientation rather than for a comparison between the models and the observations. They represent H-burning PNN whereas the WR-type PNN are certainly He-burners. As can be seen from Fig. 4 the WR-type PNN tend to situate towards upper- left in the diagram in comparison with the H-rich sample. This would normally imply higher PNN masses. It is, however, premature to conclude like that since the nuclear burning mode may be different in the two samples.

For most of the WR-type PNN in our sample, especially for those with low

T f

Z , we have obtained large values for the parameter (see Fig. 4). This indicates that their PN are rather compact and young. Therefore we can conclude that the H-poor layers in these stars have been exposed shortly after the PN formation at the tip of AGB. This excludes, from our considerations, scenaria like that of a ®nal helium shell ¯ash (Iben et al. 1983) which produce a H-poor, He-burning nucleus surrounded by an old, large nebula. He-burning PNN models have been calculated by a number of authors (e.g. Iben 1984, Wood and Faulkner 1986, Vassiliadis and Wood 1993). Unfortunately, most of them cannot be compared to the WR-type PNN since they conserve H-rich envelopes throughout the PN phase. The only set of He-burning PNN models in

which the H-poor layers are quickly exposed after having left the AGB is the one of = Wood and Faulkner (1986) which leaves the AGB at phase 0 and adopts the type B mass loss. However, this set cannot account for the WR-type PNN either.

The reason is that the effective temperature in these models very quickly jumps up

T =

to above log e 5 0 when the PNN atmosphere becomes H-poor. This is in

obvious contradiction with the fact that most of the WR-type PNN in our sample T

have log Z well below 5.0.

4. Discussion

The conclusion that arises from our analysis in Section 3 is that at present we have no theoretical evolutionary PNN model which could be compared to the observations of the WR-type PNN. We suggest that the reason why the He-burning PNN models quickly exposing H-poor layers (Wood and Faulkner1986) fail is that they adopt hydrostatic atmospheres. Strong WR-type spectral features prove that these PNN have very intense winds. If the wind is suf®ciently strong that it may happen that not only the emission-line features but also the continuum is formed in Vol. 43 395 the wind. In this case the observed effective temperature would be lower than that in the case of a hydrostatic photosphere. This idea is supported from the model atmosphere analysis of some WR-type PNN in Leuenhagen et al. (1993). In their model of NGC 6751 all the observed spectrum is formed in the expanding wind. Now a question arises: what is the mechanism of generation of such strong winds? We want to discuss two possible scenaria which, in our opinion, seem to be most plausible. The ®rst idea is based on the assumption that the WR-type PNN are single stars. This implies that the He-shell ¯ash has to be capable, at least in some cases, of generating a very intense mass loss. The mass loss should quickly remove H-rich matter and should remain strong also when the H-poor layers are exposed. In the model of NGC 6751 of Leuengahen et al. (1993) the wind carries two orders of magnitude more momentum than that available in the radiation. Thus the wind cannot be radiation-driven and the main problem is to ®ndawayin which the He-shell ¯ash can drive such an intense mass loss. In our second scenario the WR-type PNN appears as a phase in the evolution of a binary system. The evolution of binaries in the context of formation of PN has recently been thoroughly discussed by Yungelson et al. (1993). If the separation of binary components is not too large then the system passes by a phase of common envelope. This phase induces a very intense mass loss and in many cases gives rise to the PN phenomenon. If one wants to explore this idea for explaining the nature of the WR-type PNN then the abundances observed in the winds of these objects imply that the common-envelope phase should occur when the primary component is on the AGB. Evolutionary calculations of AGB stars (e.g. Vassiliadis

and Wood 1993) show that for initial masses below 3.5 M the stellar radius during a He-shell ¯ash reaches higher values than during the subsequent H-burning phase. This implies that if an AGB star enters into contact with its Roche lobe then in most cases it occurs during the He-burning phase. The Roche lobe over¯ow of an AGB star inevitably leads to the common-envelope phase (e.g. Yungelson et al. 1993). During this phase the binary spirals in and a very intense mass loss is generated (e.g. Livio andSoker1988). The time scale ofthe common-envelope phaseis much shorter than that of the He-shell ¯ash. Thus if the common-envelope ejection is suf®ciently effective, in the sense that all the H-rich matteris lost, ®nalphasesofthe common-envelope phase are expected to resemble to the observed characteristics of the WR-type PNN. The primary component is a He-burning post-AGB star with a strong H-poor wind provoked by the secondary rotating either within the remnant common-envelope or close to the primary surface. The consequence of this scenario is that some, especially more evolved (hotter) WR-type PNN, should be observed as binaries. As far as we know, no WR-type PNN has been detected to be binary. However, given the brightness of the primary, the presence of wide, intense emission lines in its spectrum and the contamination from the surrounding nebula an observational detection of a faint (e.g. low-mass main sequence star) secondary would be very dif®cult. 396 A. A.

Acknowledgements. This work was in part supported from the KBN grant No. 2-2114-92-03.

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