2 SMC data

2.1 Observations

Observations were performed using telescopes at ESO (1.52 m) and at Laboratório Nacional de Astrofísica (LNA), Brazil (1.60 m). At ESO a Boller & Chivens Cassegrain spectrograph was used with a Loral/Lesser CCD and grating allowing a reciprocal dispersion of about 2.5 Å/pixel. At LNA, a Boller & Chivens Cassegrain spectrograph was also used with a SITe CCD but with a smaller dispersion, namely, 4.4 Å/pixel. A long east-west slit was used in all observations. The log of the observations is given in Table 1.

   

Table 1: Log of the observations

Object	$\alpha_{2000}$	$\delta_{2000}$	Date	Observatory
SMP 01	00 24 05	-73 37 22	1999 Aug. 19	ESO
SMP 02	00 32 39	-71 41 59	1999 Aug. 19	ESO
SMP 03	00 34 21	-73 13 24	1999 Aug. 19	ESO
SMP 04	00 40 42	-75 17 00	1999 Aug. 20	ESO
SMP 05	00 41 22	-72 45 17	1999 Jul. 18	LNA
SMP 06	00 41 28	-73 47 07	1999 Jul. 18	LNA
SMP 07	00 42 28	-73 20 55	1999 Aug. 15	ESO
SMP 08	00 43 25	-72 38 18	1999 Aug. 20	ESO
SMP 09	00 45 21	-73 24 00	1000 Dec. 28	ESO
SMP 10	00 47 00	-72 49 16	1999 Jul. 20	LNA
SMP 11	00 48 36	-72 58 00	1999 Jul. 20	LNA
SMP 12	00 49 21	-73 52 59	1999 Aug. 16	ESO
SMP 13	00 49 52	-73 44 23	1999 Dec. 26	ESO
SMP 14	00 50 35	-73 43 00	1999 Dec. 27	ESO
SMP 16	00 51 27	-72 26 11	1999 Dec. 29	ESO
SMP 17	00 51 56	-71 24 45	1999 Aug. 17	ESO
SMP 18	00 51 58	-73 20 32	1999 Aug. 17	ESO
SMP 19	00 53 11	-72 45 07	1999 Aug. 18	ESO
SMP 21	00 56 31	-72 27 01	1999 Dec. 26	ESO
SMP 22	00 58 37	-71 35 49	1999 Aug. 17	ESO
SMP 23	00 58 42	-72 56 59	1999 Aug. 20	ESO
SMP 25	00 59 41	-71 38 16	1999 Dec. 28	ESO
N 9	00 43 37	-73 02 26	1999 Aug. 15	ESO

Image reduction and analysis were performed using the IRAF package, including the classical procedure to reduce long slit spectra: bias, dark and flat-field corrections, spectral profile extraction, wavelength and flux calibrations. Atmospheric extinction was corrected using mean coefficients for each observatory, and flux calibration was secured by the observation of standard stars (at least three) every night.

Emission line fluxes were calculated assuming Gaussian profiles, and a Gaussian de-blending routine when necessary. A table with dereddened line intensities is available in electronic form at CDS. Adopting a scale in which $I({\rm H}\beta$) = 100, typical errors in the intensities are of about 15% for lines stronger than 10 and of about 30% for weaker lines. Interstellar reddening was estimated using the Balmer ratio H$\alpha$/H$\beta$, assuming Case B (Osterbrock 1989) and adopting the extinction law by Cardelli et al. (1989). E(B-V) values derived for each nebula are given in Table 2.

2.2 Physical parameters

Electron densities were estimated from the [SII] ratio $\lambda$6716/$\lambda$6731 and from the [ArIV] ratio $\lambda$4711/$\lambda$4740, when the appropriate lines were available. Electron temperatures were derived from both [OIII] $\lambda$4363/$\lambda$5007 and [NII] $\lambda$5754/$\lambda$6584 line ratios. Whenever these temperatures were comparable, we have adopted a mean value, otherwise the [OIII] temperature was used to estimate abundances of higher ionization potential ions like O⁺², S⁺², Ar^+2,+3, Ne⁺² and the [NII] temperature of lower potential ions like O⁺, N⁺, S⁺. Temperatures and densities adopted for each nebula are also listed in Table 2. A compilation of physical parameters for Magellanic planetaries was prepared by Richer (1993). A comparison between common objects indicates that the average temperature difference between literature and our estimates amounts to 640 K, with a dispersion of 1680 K. Electron densities have a higher difference (of about 60% on the average). These differences are probably due to artifacts. They can be related to errors in the flux determination of the [SII] lines, frequently weak and with considerable uncertainties in their fluxes. Different origins of the atomic data used to derive densities from the [SII] ratio could also account for this dispersion, however, it should be emphasized that no systematic effects were observed between the different samples.

   

Table 2: Electron temperatures and densities

Object	E(B-V)	$T_{\rm e}$[OIII](K)	$T_{\rm e}$[NII](K)	$N_{\rm e}$[SII](cm-3)	$N_{\rm e}$[ArIV](cm-3)
SMP 01	0.11	11900	11800	47600:	-
SMP 02	0.11	13700	11460	4090	4270
SMP 03	0.04	13270	-	-	21000:
SMP 04	0.00	14980	12700:	-	4860
SMP 05	0.42	12400	10800	4160	2300
SMP 06	0.23	12800	9700	>3300:	60000:
SMP 07	0.09	15660	9300:	1140	-
SMP 08	0.02	11800	10200:	700	3780
SMP 09	0.60	13470	10000:	770	880
SMP 10	0.57	9800	-	860:	-
SMP 11	0.37	14000	11500:	1800:	-
SMP 12	0.00	13800	9100:	880:	370:
SMP 13	0.03	10800	-	10000	-
SMP 14	0.03	11600	-	430	500
SMP 16	0.02	10700	-	3700	-
SMP 17	0.27	11230	10300	3270	400
SMP 18	0.17	12400:	12100:	4090	-
SMP 19	0.28	11570:	9780:	1580	2900
SMP 21	0.16	23500	23300:	8400	11220
SMP 22	0.24	22400	10620	1550	-
SMP 23	0.00	13700	11700:	-	1210
SMP 25	0.00	38700	-	3760	5390
N 09	0.15	12700	12000	170	-

   

Table 3: Chemical abundances

Object	He	$\varepsilon$(O)	$\varepsilon$(N)	$\varepsilon$(Ne)	$\varepsilon$(S)	$\varepsilon$(Ar)	type
SMP 01	0.086	8.01	7.05	7.33	5.95:	5.73	-
SMP 02	0.110	8.28	7.07	7.70	6.13:	5.66	-
SMP 03	0.103	8.18	7.04	7.41	6.62	5.48	-
SMP 04	0.132	7.92	-	-	-	5.46	I (?)
SMP 05	0.120	8.51:	7.17	7.51	6.41	5.80	I
SMP 06	0.12:	8.36	7.34	-	6.35	5.88	I
SMP 07	0.102	8.17	7.14	7.46	-	-	-
SMP 08	0.116	8.24	6.74	7.75	6.15	5.77	-
SMP 09	0.11:	8.48	7.08	-	6.49	6.27	I
SMP 10	0.14:	8.55	7.30	-	6.67	6.48	I
SMP 11	0.100	8.37	6.39	7.90	6.26	6.08	-
SMP 12	0.11:	8.08	7.62	7.27	6.37	5.71	I
SMP 13	0.128	8.41	6.80	-	-	5.73	I
SMP 14	0.116	8.41	7.13	-	6.49	5.73	I
SMP 16	0.086	8.41	6.90	-	5.90	5.86	-
SMP 17	0.14:	8.48	7.19	7.60	6.25	6.05	I
SMP 18	0.095	7.99	6.83	7.25	6.16	5.75	-
SMP 19	0.115	8.47	6.82	-	6.14	5.57	I
SMP 21	0.15:	7.57	7.41	-	6.31	5.53	I(?)
SMP 22	0.152	7.59	8.26	6.97	6.10	5.48	I
SMP 23	0.098	8.05	-	7.38	-	5.50	-
SMP 25	0.105	7.01	6.93	-	5.54	5.09	-
N 09	0.085	8.20	6.52	7.55	6.28	5.79	-

2.3 Chemical abundances

Ionic abundances were calculated for each ion of interest by solving the statistical equilibrium equations for a three-level atom model, including radiative and collisional transitions. Elemental abundances were then derived through ionization correction factors (icf) adopted to account for unobserved ions of each element. We used the same icf's adopted in our precedent publications (Costa et al. 1996). Resulting abundances are given in Table 3. Typical errors are about 0.2 dex for O, N, Ne and 0.3 dex for S and Ar. Helium abundances deserve further attention, since collisional corrections by Clegg (1987) may be overestimated up to a factor of two (Peimbert & Torres-Peimbert 1987). Here we have adopted Clegg's formulae multiplied by an empirical factor equal to 0.6, since this gives a much better agreement between He⁺ abundances derived from the recombination lines HeI$\lambda$4471, 5876 and 6678. In the case of helium, when three digits are given, errors are of the order of 0.004 and, when discrepancies between the three HeI lines are greater than 0.02 the mean value is followed by ":''.

Previous chemical studies of PN in the SMC have shown that taking the whole sample, the average oxygen abundance tends to be higher than the mean values found for type I objects only (see, for instance, Leisy & Dennefeld 1996), a result not confirmed by the present study, as will be seen later. If some non-type I objects are misclassified, this could be a possible explanation. Galactic type I PN, besides the He and N excess, are characterized by their bipolar morphology, average distance to galactic plane and peculiar $<\Delta V>$velocity (Maciel & Dutra 1992). At the distance of the Clouds, morphologies are not presently available and one should wait for future high angular resolution observations, using large aperture southern telescopes (VLT, Gemini) now becoming operational. The Clouds and the Galaxy have different dynamical structures, so the kinematical and space distributions are not directly comparable. Thus, for the moment, He and N abundances are the only classification criteria. As de Freitas Pacheco et al. (1993a) have already emphasized, since the Clouds have lower metallicities, the self-enrichment condition to be applied is not necessarily the same as that first defined by Peimbert & Torres-Peimbert (1983). Here we adopted a more conservative position. Since our main goal is to investigate the reality of the anti-correlation N/O vs. O/H, we used only He abundance to classify the objects in our sample, assuming that objects with He/H > 0.11 are genuine type I PN as indicated in the last column of Table 3. This is a more drastic limit than that assumed either by de Freitas Pacheco et al. (1993a) or by Leisy & Dennefeld (1996), but it gives a higher confidence in our analysis. We will return to this point later.

Average abundances are shown in Table 4, either for type I PN as for all the sample. Objects with uncertain He abundances were excluded when average values were computed for type I planetaries. For comparison, average abundances for HII regions taken from Dennefeld (1989) are also given, excepting for oxygen, taken from the HII region sample by Russel & Dopita (1990).

   

Table 4: Average abundances

Object	$\varepsilon(\rm O)$	$\varepsilon(\rm Ne)$	$\varepsilon(\rm S)$	$\varepsilon(\rm Ar)$
Type I	8.33 $\pm$ 0.28	7.33 $\pm$ 0.24	6.35 $\pm$ 0.16	5.87 $\pm$ 0.30
All	8.22 $\pm$ 0.27	7.47 $\pm$ 0.23	6.28 $\pm$ 0.20	5.76 $\pm$ 0.25
HII reg.	8.13	7.22	6.32	5.78

Inspection of Table 4 indicates immediately that no significant differences exist between average abundances of type I PN and the ensemble average, although type I seem to have slightly higher values. However we should emphasize that we are still playing with small numbers and a large sample is necessary to draw a more firm conclusion. Nevertheless our mean values are quite consistent with interstellar medium abundances and, in particular oxygen. It is worth mentioning that Hill et al. (1997), from the study of six K-supergiants, obtained a mean oxygen abundance $\varepsilon(\rm O)$ = 8.14 $\pm$ 0.12, which agrees both with PN and HII region average values. A similar behavior is observed for galactic (de Freitas Pacheco 1993; Costa et al. 1996) and LMC planetaries (de Freitas Pacheco et al. 1993a). Taking at face value, these results may be interpreted in the sense that oxygen was preserved, not being affected by any dredge-up episode.