1 Introduction

1.1 Central stars of planetary nebulae showing the Wolf-Rayet phenomenon

It is well known that the Wolf-Rayet (WR) phenomenon is not restricted to bright, massive stars, but that it is also found among the central stars of some ($\approx$50) planetary nebulae (PN), the so-called [WC] stars (Acker et al. 1992; Tylenda et al. 1993; Peña et al. 1998). About twenty percent (van der Hucht 1996; van der Hucht 1999) of the known stars showing the WR phenomenon in our Galaxy are central stars of PN. All the PN nuclei exhibiting a WR spectrum belong to the WC sequence (Tylenda et al. 1993), and appear virtually hydrogen free. No WN-type central star is known now that M 1-67 (= Sh 2-80), surrounding WR 124 (Spectral type: WN 8), has been removed from the list of bona-fide PN (Crawford & Barlow 1991).

The similarities in line profiles suggests that the winds of [WC] central stars are scale models of the winds of the massive WC stars. However, the level of excitation conditions among WR central stars is quite different since it spans a large range, from [WC 2] to [WC 11] (Méndez & Niemela 1982; Hu & Bibo 1990; van der Hucht 1996), compared to WC 4-WC 9 (with extension to WO 1-WO 5 at the hot-end) for massive, Population I WR stars (van der Hucht 1996). Note that such an extended distribution of spectral types (although the [WC 5-7] subtypes are apparently less represented; Acker et al. 1996; van der Hucht 1996) may additionally provide a broad baseline for comparison and detection of overall trends that otherwise might be drowned out in data generally showing large intrinsic dispersions within a given spectral type.

Although the loss of the outer hydrogen-rich envelope appears to be a necessary condition to the onset of the WR phenomenon, it is clearly not a sufficient one: the large majority of PN central stars ($\approx$50 of $\approx$350 central stars for which the spectrum is known) does not have a WR-like spectrum (Acker et al. 1992). We still do not know exactly what determines some PN central stars to become [WC] stars. Moreover, the observational data, despite their incompleteness or low accuracy for many [WC] central stars, suggest that what distinguishes a [WC] star is not its present physical properties (Acker et al. 1996; Pottasch 1996), but rather more likely its initial properties and evolutionary history. This complicates the study of their precise origin. However, the status and evolutionary history of the PN central stars, as well as their ultimate fate as white dwarfs, is somewhat better known than that of their massive counterparts. The latter point combined with the broad range of excitation conditions of the nuclei suggests that [WC] central stars may facilitate understanding the WR phenomenon as a whole.

1.2 Fragmented, hot star winds; probing the turbulent structure of [WC] winds

Moving subpeaks are systematically seen on the tops of broad optical emission lines from massive WR stars (Robert 1992; Lépine & Moffat 1999, and references therein). These subpeaks suggest WR winds are inhomogeneous and non-stationary on a time-scale of hours. Eversberg et al. (1998) have also reported stochastic, variable substructures in the He II$\lambda$4686 emission line originating in one O supergiant star. The latter study points to the likely universal nature of wind clumping originating in massive, hot stars.

In Grosdidier et al. (2000; hereafter Paper I), wind fluctuations were described for four [WC 9-10] stars, including BD +30$^\circ$3639 ([WC 9]) observed intensively during 15 nights. In the latter study, the authors show that the wind clumping originating in BD +30$^\circ$3639 is remarkably similar to that reported for one of its massive, WC 9 counterparts, WR 103. Therefore, they interpreted this fact as strong evidence for understanding the WR phenomenon as a purely atmospheric phenomenon independent of the differences between massive and low-mass WR stars. The present paper will discuss the case of the hotter subtype [WC 8]. The case of the cooler subtypes, [WC 10-12], will be investigated in detail in future studies.

In order to resolve the narrow subpeaks present on the tops of the emission lines originating in BD +30$^\circ$3639 ([WC 9]), Grosdidier et al. (Paper I) and Acker et al. (1997) found it necessary to have a spectral resolution of about one Angström, or better. For NGC 40, the Balick et al. (1996) spectroscopic data and our first observations (January 96; see Sect. 2 for details) demonstrated to us that a 3 Å spectral resolution is sufficient. Since the subpeaks are very weak, securing time resolution along with sufficient S/N ratio imposes the use of large telescopes. As a compromise, one has to concentrate on relatively bright [WC] central stars observed intensively, especially when using 2 m class telescopes, as in the present study.

Some 17 [WC]-late ([WC 8-12]) central stars are known within the Galaxy (Górny & Stasinska 1995; Acker, private communication) and only two belong to the [WC 8] spectral type. The [WC 8] nucleus of M 2-43 (= PN G027.6+04.2) is certainly too faint ( $V\approx 15.7$) for our project. Therefore, the central star of NGC 40, HD 826, which is 4.1 mag brighter in the visible domain ( $V\approx 11.6$), appears obviously to be the best target for studying [WC 8] spectroscopic flickering. We concentrated our intensive spectroscopic program on the C III$\lambda$5696 and C IV $\lambda \lambda$5801/12+C III$\lambda$5826 emission lines originating in HD 826 observed intensively during 22 nights with 2 m class telescopes. Because they are relatively bright and have comparable intensities in [WC 8] stars, they are the best lines to study expanding stellar wind variability in the optical domain. In addition, the blend-free C III$\lambda$5696 emission line constitutes an excellent line to trace the movements of independent subpeaks, which are blurred by mixing in the adjacent blended C IV $\lambda \lambda$5801/12+C III$\lambda$5826 emission line.

1.3 The [WC 8] central star of the planetary nebula NGC 40: HD 826

The PN NGC 40 is a well known nebula, with a 48 $^{\prime\prime}$-diameter barrel-shaped core, surrounded by two haloes; a faint, diffuse, inner halo out to 90 $^{\prime\prime}$, and an outer halo with jet-like structures extending to 4$^\prime$; see Meaburn et al. (1996). These authors noted that turbulent motions exist in the nebula, an observation confirmed by the analysis of nebular line profiles (Neiner et al. 2000). NGC 40 is unusual, because the low excitation class (p=1) of the nebula suggests a stellar temperature of about 30000 K, whereas the UV spectrum of the nucleus is compatible with a temperature three times higher. This discrepancy can be explained by the existence of a "carbon curtain'' in the nebula (Bianchi & Grewing 1987). C II$\lambda$6578 emission was observed to be coincident with the 48 $^{\prime\prime}$ [N II] shell, implying that the expanding central envelope is relatively rich in carbon (see Meaburn et al. 1996). The PN NGC 40 probably originated from a relatively massive progenitor (6 $M_{\odot}$; Bianchi 1992).

 

 

Table 1: Log of spectroscopic observations of NGC 40's central star. The exposure times were 1000 s (OMM) and 1800 s (OHP)

Denomination	Central star		Journal	of	observations
PNG	Spectr. type	Telescope (diam.)	Spectr. range	S/N	Date	No. of
Usual name PN	Va	Spectrograph	Resol. power (RP)			spectra
Star name
(1)	(2)	(3)	(4)	(5)	(6)	(7)
120.0+09.8	[WC 8]	OMM (1.6 m)	5300-5960 Å	78	1996 Jan. 11	12
NGC 40	11.6	B&C	$\approx$2.8 Å (2000)	48	1996 Jan. 15	22
HD 826				44	1996 May 26	15
				42	1996 May 27	15
				40	1996 May 30	13
				80	1996 Jul. 28	18
				56	1996 Sep. 26	26
				65	1996 Sep. 30	19
				43	1996 Nov. 16	34
				63	1996 Nov. 17	32
		OHP (1.52 m)	5250-6000 Å	36	1997 Jan. 12	20
		AURELIE	$\approx$1.1 Å (5000)	39	1997 Jan. 13	19
				26	1997 Mar. 3	8
				21	1997 Mar. 5	13
				16	1997 Mar. 6	12
				29	1997 Mar. 7	14
				7	1998 Jan. 20	5
				17	1998 Jan. 21	18
				19	1998 Jan. 22	19
				17	1998 Jan. 23	17
				20	1998 Jan. 24	12
				18	1998 Jan. 25	9

^a From the Acker et al. (1992) catalogue.
^b Characteristic signal-to-noise ratio evaluated in the continuum adjacent to C III$\lambda$5696.

The effective temperature of HD 826 was estimated at 46000 K by Leuenhagen et al. (1996), although some authors report effective temperatures as low as 30000 K (Köppen & Tarafdar 1978), or as high as 90000 K (Schmutz et al. 1989; Bianchi & Grewing 1987). Using IUE data, Bianchi & Grewing (1987) reported a terminal velocity of 1800 km s^-1 and a luminosity of 26200 $L_{\odot}$. From these results they inferred a radius of 0.66 $R_{\odot}$ for the continuum-emitting region. Earlier UV spectroscopic data obtained by Benvenuti et al. (1982) led to an even larger terminal velocity, $v_{\infty}\approx 2370$ km s^-1. However, such a value is not reliable because of the poor spectral resolution of their instrumentation. Recent, reliable modelling of the expanding atmosphere (Leuenhagen et al. 1996) led to about 1000 km s^-1 for the terminal velocity and 0.33 $R_{\odot}$for the stellar core radius. PN nucleus spectroscopic flickering similar to that observed for massive WR stars was reported for the first times by Balick et al. (1996) and Grosdidier et al. (1997). They found that the [WC 8] central star of HD 826 shows fast moving subpeaks on the top of its flat-topped C III$\lambda$5696 emission line. Balick et al. (1996) also reported a nearly similar apparent acceleration for all features, the acceleration zone being at least 5 $R_{\odot}$ in extension.