A spectroscopic study of the bipolar planetary nebula Mz 3

Transcription

1 Mon. Not. R. Astron. Soc. 337, (2002) A spectroscopic study of the bipolar planetary nebula Mz 3 Y. Zhang 1 and X.-W. Liu 2,3 1 National Astronomical Observatories, Chinese Academy of Sciences, Beijing , China 2 Department of Physics and Astronomy, University College London, Gower Street, London WC1E 6BT 3 Department of Astronomy, Peking University, Beijing , China Accepted 2002 July 30. Received 2002 July 26; in original form 2002 May 20 ABSTRACT We have obtained a medium-resolution, deep optical long-slit spectrum of the bipolar planetary nebula Mz 3. The spectrum covers the wavelength range Å. Over 200 emission lines have been detected, many of them permitted and forbidden transitions from Fe +,Fe 2+ and Fe 3+ and other iron-group elements. The spectra have been used to determine nebular thermal and density structures and elemental abundances. The very rich and prominent [Fe III] emission lines observed in the optical spectrum of Mz 3 are found to originate exclusively from an unresolved emission region centred on the central star. The relative intensities of [Fe III] lines arising from the same upper level are in good agreement with the theoretical predictions. The [Fe III] lines detected in Mz 3, arising from levels of different excitation energies and critical densities, provide powerful diagnostic tools to probe the physical conditions in the central emitting region. We find that all the observed [Fe III] diagnostic line ratios consistently yield an electron temperature of T e = K and a density of log N e (cm 3 ) = 6.5. The latter value is close to the densities where the ratios of these [Fe III] diagnostic lines are most sensitive to density variations, suggesting that the density in the central emission core could be even higher. In contrast, all the other standard nebular density- and temperature-diagnostic line ratios, all of lower critical densities than the [Fe III] lines and therefore only useful at densities 10 6 cm 3, yield consistently lower electron densities, with the resultant values correlating with their critical densities. This indicates that the central dense emission core has a highly stratified density structure such that forbidden lines of relatively low critical densities are collisionally suppressed in high-density regions. Given the highly stratified thermal and density structures of the core emission region, no reliable elemental abundances can be determined for this region, except possibly for iron, for which we find an abundance of approximately half the solar value. In contrast to the dense central core, emission from the extended bipolar lobes of Mz 3 are well represented by a mean electron temperature of 6800 K and density of 5200 cm 3. The elemental abundances derived for this region, on a logarithmic scale where H = 12, are He = 10.93, C = 8.41, N = 8.37, O = 8.50, Ne = 7.64, S = 7.15, Cl = 5.28 and Ar = In all cases, the abundances are close to the average values deduced for Galactic planetary nebulae. In particular, there is no evidence of He enrichment, as claimed in previous studies. The newly derived heavy-element abundances are significantly higher than those published in the literature. In those earlier studies, as a result of the contamination of emission from the dense central core, the average electron temperature in Mz 3 was significantly overestimated, resulting in grossly underestimated heavy-element abundances. However, Mz 3 does seem to have a relatively high N/O abundance ratio, in line with the result derived from the ISO observations of the far-infrared fine-structure lines. Key words: line: identification ISM: abundances ISM: lines and bands planetary nebulae: individual: Mz 3. zhangy@vega.bac.pku.edu.cn C 2002 RAS

2 500 Y. Zhang and X.-W. Liu 1 INTRODUCTION Mz 3 (PN G ) is a young bipolar planetary nebula (PN) with lobes extending over 50 arcsec on the sky. The nebular emission region consists of a bright core, two approximately spherical bipolar lobes and two outer large filamentary nebulosities on either side of the central star (Cohen et al. 1978; Redman et al. 2000). The salient features of Mz 3 are more easily studied than those of other bipolar nebulae because of its large angular size. It shows strong emission in the far-infrared. Imaging of Mz 3 at 10 µm reveals that there is an extended shell of warm ( K) dust surrounding the central star (Quinn et al. 1996). V-band imaging polarimetry by Scarrott, Scarrott & Wolstencroft (1994) suggests that the bipolar lobes, bounded by shocked gas swept up by the fast stellar wind, are filled with hot ionized gas. It is believed that bipolar nebulae are produced by more massive progenitor stars than those of the other morphological classes (e.g. Gorny, Stasinska & Tylenda 1997). Bipolar nebulae also tend to have higher He and N abundances. Peimbert & Torres- Peimbert (1983) first noticed the association of bipolar nebulae with Type I planetary nebulae (PNe), i.e. PNe with abundance ratios N(He)/N(H) and/or log (N/O) 0.3. In their study of the possible links between the morphology and chemical composition of PNe, Calvet & Peimbert (1983) classified Mz 3 as a Type I PN based on the abundances determined by Cohen et al. (1978), who found that Mz 3 has a helium abundance of Cohen et al. found that that Mz 3 also has a high nitrogen abundance. On the other hand, Meaburn & Walsh (1985) presented new spectroscopic observations and found that Mz 3 is enriched in helium but not in nitrogen. Several dynamical studies of Mz 3 have been carried out. Lopez & Meaburn (1983) showed that there are three distinct velocity components in the core of Mz 3. They found that the inner, bright, approximately spherical lobes expand radially at 50 km s 1. They also detected wings in the Hα profiles that extend to ±1200 km s 1, caused by electron scattering in the hot ionized inner regions around the core. Lopez & Meaburn presented two models to explain the observed kinematic properties of Mz 3, i.e. an ejection model and a steady-state model. Assuming a stellar wind velocity of V = 1229 km s 1, the two models yield mass-lose rates of Ṁ = and M yr 1, respectively. The outer, bipolar, filamentary lobes consist of two components of different spatial extent and expansion velocity. Meaburn & Walsh (1985) presented an empirical model that can explain many of the obtained velocity structure features and derived that the two components of the filamentary lobes expand respectively at velocities of 90 and 180 km s 1. The observed profile of Na I absorption lines suggests the presence of a neutral torus with an ionized inner surface in the core (Meaburn & Walsh 1985). The ionized inner surface has an electron density of N e cm 3 and expands at 20 km s 1. They suggest that the central star of the bipolar nebula may have recently undergone a mass-loss event. A recent spatially resolved spectral analysis revealed a pair of hypersonic ( 500 km s 1 ) emission knots along the axis of the inner bipolar lobes of Mz 3 (Redman et al. 2000). Redman et al. suggest that the tips of the inner bipolar lobes have encountered a much lower density environment and thus are undergoing a breakup of the accelerating parts of the shell as a result of the development of a Rayleigh Taylor instability. These observations seem to suggest the presence of shocks in Mz 3, which may play a role in heating the nebula. Riera (1996) suggest that the presence of shocked regions in Type I PNe is one of the reasons causing the spectral differences between Type I and non-type I PNe. Assuming a distance of 1.8 kpc and an effective temperature of K, Cohen et al. (1978) obtained a lower limit of 5700 L for the bolometric luminosity of the central star of Mz 3. Lopez & Meaburn (1983) argued that the distance to Mz 3 is probably about 1 kpc on the basis of the broad wing in the Hα profile produced by electron scattering. From a statistical study of 42 IRAS point sources, van der Veen, Habing & Geballe (1989) derived a distance of 3.3 kpc and a total luminosity of L for Mz 3. A lower Shklovsky distance of 1.3 kpc was deduced by Cahn, Kaler & Stanghellini (1992). Using a revised Shklovsky method (Barlow 1987), Kingsburgh & English (1992) obtained a distance of 2.7 kpc, implying a moderate luminosity of L. In this paper, we present a medium-resolution, deep optical longslit spectrum of Mz 3, taken approximately along the nebular major axis. The spectrum is dominated by very rich and prominent permitted and forbidden transitions of iron ions as well as of other iron-group elements. In particular, the [Fe III] line spectrum is phenomenal, with dozens of lines detected. The strong [Fe III] emission is found to arise exclusively from a central unresolved emission core that coincides in position with the central star. A plasma diagnostic analysis of the [Fe III] lines yields an electron temperature of K and an extremely high electron density of log N e (cm 3 ) = 6.5. There is clear evidence that the dense emission core has a highly stratified density structure and that other commonly observed density-diagnostic lines, such as the [S II], [O II] and [Cl III] doublets and the [N II] and [O III] nebular lines, are collisionally suppressed in the high-density emission regions. The extended emission regions surrounding the bright core of Mz 3 are found to have a much lower temperature and density. The two bright bipolar lobes are well represented by a more or less uniform ionized region of T e = 6800 K and N e = 5200 cm 3. The elemental abundances derived are much higher than those previously published in the literature and are comparable to the average values of Galactic PNe. Owing to contamination of the emission from the dense central core, the electron temperatures deduced in earlier studies were grossly overestimated, leading to significantly underestimated elemental abundances. Our results show that Mz 3 is not enriched in helium but probably has an above-solar N/O abundance ratio. Section 2 describes the observations and data reduction and presents the observed line fluxes. Dust extinction towards Mz 3 is discussed in Section 3, along with a brief description of several absorption features detected in the spectrum, including one identified as the strongest diffuse interstellar band (DIB) at 6284 Å. In Section 4, we present plasma diagnostic results using the standard forbidden-line diagnostic ratios as well as the rich [Fe III] lines observed in the spectrum of Mz 3 and the spatial variations of electron temperature and density across the nebula. Ionic and elemental abundances are derived in Section 5. In Section 6, we discuss briefly the nature and evolutionary status of this extraordinary nebula in light of these new observational results. 2 OBSERVATIONS AND DATA REDUCTION The observations were carried out with the ESO 1.52-m telescope using the Boller & Chivens (B&C) spectrograph. A journal of observations is presented in Table 1. The detector was an ultravioletenhanced Loral , µm 2 chip. The B&C spectrograph has a useful slit length of about 3.5 arcmin. In order to reduce the charge-coupled device (CCD) readout noise, in 1996 the CCD

3 Planetary nebula Mz Figure 1. Optical spectrum of Mz 3 from (a) 4020 to 4950 Å and (b) 4850 to 7250 Å. In each panel, the upper spectrum is that of the central core (position 1) whereas the lower spectrum is that of the extended bright inner lobes (position 2). Note the very rich and prominent [Fe III] emission lines in the spectrum of the central core. Some of the most prominent iron lines, along with a few from other ionic species, are marked. Note that Hα and [N II] λ6584 were saturated in this deep exposure. Dust extinction has not been corrected for. Table 1. Observational journal. Date λ range (Å) FWHM (Å) Exp. time (s) 1996 July , 60, 180, July June June , 4 180, was binned by a factor of 2 along the slit direction. A slit width of 2 arcsec was used throughout the observations. The slit was oriented in the north south direction (PA = 0 ), approximately along the nebular major axis, and through the central star. Two wavelength ranges were observed in 1996: Å at a resolution of 1.5 Å full width at half-maximum (FWHM), and Å at a resolution of 4.5 Å. At identical slit position two additional wavelength ranges were observed in 2001: Å at a resolution of 1.5 Å FWHM, and Å at a resolution of 3.2 Å. Short exposures were taken in order to obtain intensities of the brightest emission lines, which were saturated in the spectra of long exposure time. The spectra were reduced using the LONG92 package in MIDAS 1 following the standard procedure. The spectra were bias-subtracted, flat-fielded and cosmic rays removed, and then wavelengthcalibrated using exposures of a helium argon calibration lamp. Absolute flux calibration was obtained by observing the standard stars Feige 110 and the nucleus of PN NGC 7293 (Walsh 1993). The blue and red parts of the spectrum, after integrating along the slit, are plotted in Figs 1(a) and (b), respectively. Dust extinction has not been corrected for. For a 2 arcsec wide slit, our observations yield an Hβ flux of logf(hβ)(erg cm 2 s 1 ) = 11.75, compared to the value of for the whole nebula (Cahn et al. 1992). Note the rich [Fe III] lines detected in the spectrum. As shown in Fig. 1, the very rich and prominent [Fe III] and [Fe II] emission detected in the spectrum of Mz 3 arises exclusively from the (unresolved) central core. In Fig. 2, we compare the observed intensity profiles of several emission lines along the slit. The [Fe III] lines are almost certainly collisionally excited. Although [Fe II] lines can also be fluorescence excited (Bautista & Pradhan 1998), it seems that the strong [Fe II] emission observed in Mz 3 is mainly produced 1 MIDAS is developed and distributed by the European Southern Observatory.

4 502 Y. Zhang and X.-W. Liu Table 3. All fluxes were normalized such that Hβ = 100. We used the standard Galactic extinction law for a total-to-selective extinction ratio of R = 3.1 (Howarth 1983) and adopted a logarithmic extinction at Hβ of c(hβ) = 2.3 for position 1 and 2.0 for positions 2 and 3 (cf. Section 3). A few lines in Table 2 remain to be identified. Figure 2. The observed intensity profiles along the slit (PA = 0 ) for six emission lines, on the scale where the total intensity of Hβ equals 100 after integration along the slit. Note that the [Fe III] λ4658 and [Fe II] λ7155 lines are confined to the core emission region. Hβ and [O III] λ5007 lines are also strong in the core where the [S II] and [N II] emission is weak. by collisional excitation. Fig. 2 shows that the core also has relatively strong [O III] emission, indicating a higher ionization degree. In contrast, the spectrum of the bright inner lobes is dominated by strong forbidden lines from singly ionized ionic species typically observed in a low-excitation PN. As we will show in Section 4, the nebular emission from Mz 3 has highly stratified thermal and density structures. The electron density in the central emission core is several orders of magnitude higher than in the extended bright lobes. Electron temperature in the central core is also higher. We have therefore divided the nebula into three zones and analyse them separately: (a) a core emission region of nebular radii r < 4 arcsec (position 1); (b) an extended emission region consisting of the two bright inner lobes on either side of the central star between 4 < r < 15 arcsec (position 2); and (c) the faint, outer most filamentary emission region between 15 < r < 20 arcsec (position 3; cf. the HST WFPC2 Hα image of Mz 3 published by Redman et al. (2000). All line fluxes, except those of the strongest lines, were measured using Gaussian line profile fitting. For those strongest and isolated lines, fluxes obtained by direct integration over the observed line profiles were adopted. The observed fluxes and identifications of lines detected at positions 1 and 2 are presented in Table 2. In Table 2, the first column gives the observed wavelengths after correction for the Doppler shift as determined from the Balmer lines. The observed fluxes at positions 1 and 2 are given in columns 2 and 4, and the corresponding fluxes after correction for dust extinction are listed in columns 3 and 5, respectively. The remaining columns give the ionic identification, laboratory wavelength, multiplet number, the lower and upper spectral terms of the transition, and the statistical weights of the lower and upper levels. Similarly, the observed and reddening-corrected fluxes at position 3 are presented in 3 DUST EXTINCTION AND DIFFUSE INTERSTELLAR BANDS Using the Galactic reddening law of Howarth (1983), the logarithmic extinction at Hβ derived from the observed Balmer decrements, Hα/Hβ, Hγ /Hβ, Hδ/Hβ and H9/Hβ are listed in Table 4. The values deduced for position 1 are systematically higher than those of position 2. This could be caused by local dust surrounding the central dense core (Quinn et al. 1996). On the other hand, at position 1, the extinction yielded by the Hα/Hβ ratio is significantly higher than those derived from the higher-order Balmer lines, suggesting that the emission of the lower-order Balmer lines from the dense central region could be optically thick. The Balmer line intensity ratios Hα/Hβ and Hγ /Hβ as a function of the optical depths of Lyα and Hα are given by Cox & Mathews (1969). For Hα depth of 5, the observed Hα/Hβ and Hγ /Hβ ratios at position 1 yield very similar reddening constants of c(hβ) = 2.30 and 2.32, respectively. Here we have adopted c(hβ) = 2.3 to deredden the optical spectrum at position 1. As for position 2, the values of c(hβ) derived from the Hα/Hβ, Hγ /Hβ and Hδ/Hβ ratios were averaged with weights of 3 : 1 : 1, roughly proportional to the intrinsic values of the three diagnostic ratios. This yielded c(hβ) = 2.0, and was used to deredden the observed line fluxes measured at both positions 2 and 3. Several absorption features have been detected in the spectrum of Mz 3, including a strong broad feature near 6284 Å, which is attributed to the diffuse interstellar band (DIB) at the same wavelength. The identification is strengthened by the detections of three other DIBs, at 5780, 6177 and 6205 Å, respectively. Fig. 3 shows the profiles of the strongest DIB at 6284 Å. The λ6284 absorption feature has an equivalent width EW = 1.9 and 1.4 Å in the spectra of position 1 and 2, respectively. Given a dust extinction colour excess of E(B V ) = 1.59 and 1.38 for position 1 and 2, respectively, we obtain a DIB strength per unit E(B V ), EW(6284)/E(B V ) = 1.2 and 1.0 for position 1 and 2, respectively. The values are lower than the typical value of 2 observed in the diffuse interstellar medium (ISM) but higher than the value of 0.2 observed towards stars in the Orion nebula (Jenniskens, Ehrenfreund & Fing 1994). The DIBs detected in Mz 3 are probably caused by absorption of local material rather than by the ISM along the line of sight. 4 PLASMA DIAGNOSTICS In this section, the electron temperatures and densities in the emission regions of Mz 3 at position 1 and 2 are determined using various diagnostic line ratios. The very rich and prominent [Fe III] emission lines detected exclusively from the central emission region (position 1) are analysed in Section Other standard diagnostic lines observed at position 1 are analysed in Section Similar analyses for position 2, i.e. the extended bright inner lobes, are given in Section 4.2. Emission from position 3 is much fainter, and many weak diagnostic lines, including the [O III] λ4363 and [N II] λ5754 auroral lines, have not been detected. Position 3 will thus be ignored from our analyses from now on.

5 Planetary nebula Mz Table 2. Observed line fluxes at positions 1 and 2, on a scale where Hβ = 100. λ obs Position 1 Position 2 Ion λ lab Mult Lower term Upper term g 1 g 2 F(λ) I (λ) F(λ) I (λ) He I V6 2s 1S 5p 1P He I V28 2p 3P 8d 3D H H20 2p + 2P 20d + 2D H H19 2p + 2P 19d + 2D H H18 2p + 2P 18d + 2D H H17 2p + 2P 17d + 2D H H16 2p + 2P 16d + 2D 8 He I V25 2p 3P 7d 3D 9 15 O III V14 3p 3P 3d 3D H H15 2p + 2P 15d + 2D 8 O III V14 3p 3P 3d 3D H H14 2p + 2P 14d + 2D 8 [S III] F2 3p 2 3P 3p 2 1S [O II] F1 2p 3 4S 2p 3 2D [O II] F1 2p 3 4S 2p 3 2D H H13 2p + 2P 13d + 2D H H12 2p + 2P 12d + 2D H H11 2p + 2P 11d + 2D H H10 2p + 2P 10d + 2D 8 [Fe V] d 4 5D 3d 4 3P He I V22 2p 3P 6d 3D 9 15 [Fe V] d 4 5D 3d 4 3F H H9 2p + 2P 9d + 2D [Fe V] d 4 5D 3d 4 3F O II V12 3p 4D 3d 2F O II V12 3p 4D 3d 4D [Ne III] F1 2p 4 3P 2p 4 1D H H8 2p + 2P 8d + 2D 8 He I V2 2s 3S 3p 3P Fe II V3 3d 7 4P 4p 6D H H7 2p + 2P 7d + 2D 8 98 [Ne III] F1 2p 4 3P 2p 4 1D 3 5 Fe II s 4P 4p 4P [Ni II] d 9 2D 4s 2D N II V12 3s 1P 3p 1D [Fe V] d 4 5D 3d 4 3P [Fe III] V15 3d 6 5D 3d 6 3G He I V18 2p 3P 5d 3D 9 15 [Ni II] d 9 2D 4s 4P [Ni II] d 9 2D 4s 4P Fe II V13 a2p z6p [Fe III] d 6 5D 3d 6 3G [S II] F1 2p 3 4S 2p 3 2P [Fe V] d 4 5D sd 4 3P [S II] F1 2p 3 4S 2p 3 2P H H6 2p + 2P 6d + 2D [Fe II] d 7 4F 4s 2H He I V16 2p 3P 5s 3S He I V53 2p 1P 6d 1D [Ni II] s 4F 4s 2G [Fe IV] d 5 4G 4g 2H Fe II V27 4s 4P 4p 4D [Fe II] d 7 4F 4s 4G [Fe II] d 7 4F 4s 2H 8 10 [Fe V] d 4 5D 3d 4 3P [Ni II] d 9 2D 4s 2D Fe II V27 4s 4P 4p 4D [Fe II] d 7 4F 4s 4G [Fe II] d 7 4F 4s 4G 8 8 N II V48a 3d 3D 4f 1[3] 5 5

6 504 Y. Zhang and X.-W. Liu Table 2 continued λ obs Position 1 Position 2 Ion λ lab Mult Lower term Upper term g 1 g 2 F(λ) I (λ) F(λ) I (λ) O II V101 3d 2G 4f H [Fe III] d 6 3P 3d 6 1F C II V6 3d 2D 4f 2F [Fe II] d 7 4F 4s 4G [Fe II] s 2 6D 4p 6S Fe II V27 4s 4P 4p 4F Fe II V27 4s 4P 4p 4D [Fe II] d 7 4F 4s 4G [Fe II] d 7 4F 4s 4G [Ni III] d 8 3F 3d 8 1G 9 9 [Ni II] d 9 2D 4s 4P H H5 2p + 2P 5d + 2D [Fe II] d 7 4F 4s 4G [Fe II] s 4D 4s 4D 6 4 [Fe II] d 7 4F 4s 4G [Fe II] d 7 4F 4s 4G 4 6 [Fe II] s 6D 4s 2 6S [O III] F2 2p 2 1D 2p 2 1S O I V5 3s 3S 4p 3P 3 1 O I V5 3s 3S 4p 3P 3 5 O I V5 3s 3S 4p 3P [Fe II] d 7 4F 4s 4G Fe II V32 4s 4H 4p 4F Fe II V27 4s 4P 4p 4D He I V51 2p 1P 5d 1D [Fe II] s 6D 4s 2 6S [Fe II] s 6D 4s 4F N II V55a 3d 3P 4f 2[3] [Fe II] s 6D 4s 2 6S [Fe II] s 6D 4s 4F [Ni II] s 4F 4s 2G He I V14 2p 3P 4d 3D [Fe II] s 6D 4s 2 6S [Fe II] s 6D 4s 4F [Fe II] s 6D 4s 4F [Fe II] s 6D 4s 4F [Fe II] s 6D 4s 4F [Co III] d 7 4P 3d 7 2F Fe II V38 4s 4F 4p 4D [Fe II] s 6D 4p 4F Fe II s 4F 4p 4F Fe II s 4F 4p 4D Fe II d 7 2D 4p 2F 4 6 Fe II V38 4s 4F 4p 4D Fe II V37 4s 4F 4p 4F [Fe II] d 7 2G 4s 2D [Fe II] V38 4s 4F 4p 4D 10 8 Fe II V37 4s 4F 4p 4F [Ni III] d 8 3F 3d 8 1G N II V5 3s 3P 3p 3P [Fe III] d 6 5D 3d 6 3F [C I] V2F 2p 2 3P 2p 2 1S 3 1 N II V5 3s 3P 3p 3P [CoIII] d 7 4P 3d 7 2F 4 6 N II V5 3s 3P 3p 3P 5 5 [Ni II] d 9 2D 4s 4P [Fe II] s 6D 4s 4P N II s 3P 3p 3P O II V1 3s 4P 3p 4D [Fe III] F3 3d 6 5D 3d 6 3F [Fe III] d 6 5D 3d 6 3F 7 5

7 Planetary nebula Mz Table 2 continued λ obs Position 1 Position 2 Ion λ lab Mult Lower term Upper term g 1 g 2 F(λ) I (λ) F(λ) I (λ) [Fe II] s 6D 4s 4P [Fe III] F3 3d 6 5D 3d 6 3F He I V12 2p 3P 4s 3S [Fe II] s 6D 4s 4P [Fe III] d 6 5D 3d 6 3F [Fe III] d 6 5D 3d 6 3F [Fe III] d 6 5D 3d 6 3F [Fe II] d 7 4F 4s 4F [Fe III] d 6 5D 3d 6 3F [Fe II] s 6D 4s 4P N II V20 3p 3D 3d 3D [Fe II] d 7 4F 4s 4F [Fe III] d 6 5D 3d 6 3F H H4 2p + 2P 4d + 2D [Fe IV] d 5 4G 3d 5 4F 6 6 [Fe IV] d 5 4G 3d 5 4F 6 4 [Fe IV] d 5 4G 3d 5 4F [Fe II] d 7 4F 4s 4F [Fe III] F2 3d 6 5D 3d 6 3H [Fe II] s 6D 4s 4P [Fe II] s 2G 4s 2D 10 6 [Fe IV] d 5 4G 3d 5 4F 10 8 [Fe IV] d 5 4G 3d 5 4F [Fe II] d 7 4F 4s 4F 8 8 [Fe IV] d 5 4G 3d 5 4F [Fe IV] d 5 4G 3d 5 4F He I V48 2p 1P 4d 1D [Fe II] V42 4s 2 6S 4p 6P [Fe III] d 6 5D 3d 6 3P 3 1 [O III] F1 2p 2 3P 2p 2 1D [Fe II] d 7 4F 4s 4F [Fe II] d 7 4F 4s 4F [O III] F1 2p 2 3P 2p 2 1D [Fe II] d 7 4F 4s 4F [Fe III] d 5 5D 3d 5 3H N II s 5P 3p 5F 5 7 N II p 3S 3d 3P [O III] F1 2p 2 3P 2p 2 1D [Fe III] d 6 5D 3d 6 3P 5 3 N II s 3P 3p 3S 3 3 N II s 5P 3p 5P 5 3 N II s 5P 3p 5P N II p 3D 3d 3F [Fe IV] d 5 4G 3d 5 2F [Ni IV] d 7 4F 3d 7 2G [Fe II] d 7 2P 4s 2D Si II V5 4p 2P 4d 2D 4 6 Si II V5 4p 2P 4d 2D 4 4 [Fe II] d 7 2G 4s 2D 8 4 [Ni II] s 4F 4s 2P [Fe II] d 7 4F 4s 4H [Fe III] d 6 5D 3d 6 3P [Fe II] d 7 4F 4s 4P [Fe II] d 7 4F 4s 4H [Ni II] s 4F 4s 2P Fe II s 4F 4p 6F [Fe VI] d 3 4F 3d 3 2G [Fe II] d 7 4F 4s 4P 8 4 [Fe II] d 7 4F 4s 4H [Fe III] d 6 5D 3d 6 3P [Fe II] s 4D 4s 2F 8 8

8 506 Y. Zhang and X.-W. Liu Table 2 continued λ obs Position 1 Position 2 Ion λ lab Mult Lower term Upper term g 1 g 2 F(λ) I (λ) F(λ) I (λ) Fe II s 2 6S 4p 6P [Fe II] d 7 4F 4s 4P [Ar III] F3 2p 4 1D 2p 4 1S N I F1 2p 3 4S 2p 3 2D 4 4 [N I] F1 2p 3 4S 2p 3 2D 4 6 Fe II V49 v2g [Fe II] d 7 4F 4s 4H Fe II V49 4s 4G 4p 4F [Fe II] d 7 4F 4s 4H [Fe III] d 6 5D 3d 6 3P 7 5 [Fe II] d 7 4F 4s 4P 10 6 [Fe II] s 4F 4s 4P Fe II s 2 6S 4p 6F [Ni IV] d 7 4F 3d 7 2G [Fe II] d 7 4F 4s 4H Fe II V49 4s 4G 4p 4F Fe II V49 4s 4G 4p 4F [Fe II] d 7 4F 4s 4H 6 10 [Co III] d 7 4F 3d 7 2P [Fe II] d 7 4F 4s 4P [Ni IV] d 7 4F 3d 7 2G [Fe II] d 7 4F 4s 4H He II f + 2F 7g + 2G [Fe III] d 6 5D 3d 6 3P Fe II V49 4s 4G 4p 4F [Fe II] d 7 4F 4s 4P N II p 3P 3d 3P [Co III] d 7 4F 3d 7 2P [Fe II] s 4D 4s 2P [Fe VI] d 3 4F 3d 3 4P N II p 3P 3d 3P Fe II d 4G 4f [Cl III] F1 3p 3 4S 3p 3 2D [Fe II] s 4D 4s 2P [Cl III] F1 3p 3 4S 3p 3 2D [Fe II] d 7 4P 4s 4D Si II p 2P 8s 2S [O I] F3 2p 4 1D 2p 4 1S O II s 2S 4p 2P [Fe II] d 7 4P 4s 4D [Fe III] d 6 3G 3d 6 1F [Fe II] d 7 4P 4s 4D N II V3 3s 3P 3p 3D N II V3 3s 3P 3p 3D 5 7 [Fe II] d 7 2G 4s 2G N II V3 3s 3P 3p 3D N II V3 3s 3P 3p 3D Si III d 3D 9f 3F ? [Fe II] s 4D 4s 2P [N II] F3 3s 1D 3p 1S [Fe II] d 7 2P 4s 2D [Fe II] V58 3d 7 2G 4s 2G He I V11 2p 3P 3d 3D [Mn V] d 3 4F 3d 3 2G [Ni IV] d 7 4F 3d 7 3P [Fe II] d 7 2D 4s 2D 4 4 [Fe IV] d 5 4P 3d 5 2F N II V28 3p 3P 3d 3D 3 5 N II V28 3p 3P 3d 3D N II V28 3p 3P 3d 3D 5 7

9 Planetary nebula Mz Table 2 continued λ obs Position 1 Position 2 Ion λ lab Mult Lower term Upper term g 1 g 2 F(λ) I (λ) F(λ) I (λ) Si II V4 4p 2P 4s 2S 2 2 N II V28 3p 3P 3d 3D 5 5 N II V28 3p 3P 3d 3D Si II V4 4p 2P 4s 2S 4 2 [Fe II] d 7 2D 5d 2S [Fe II] s 4G 4p 6F [Ni III] d 8 3F 3d 8 3P C I V p 3D 5d 3F O I V22 3p 3P 6s 3S 3 3 O I V22 3p 3P 6s 3S 5 3 O I V22 3p 3P 6s 3S [Fe III] d 6 3P 3d 6 1D [Ni IV] d 7 4F 3d 7 4P 6 4 Fe II s 4G 4p 6F Fe II V74 4s 4D 4p 4P [Mn V] d 3 4F 3d 3 4P N II d 3F 4p 3D [Fe II] d 7 2G 3d 7 2F [Co III] d 7 4F 3d 7 2G [Mn V] d 3 2D2 3d 3 2F Fe II V74 4s 4D 4p 4P Fe II V74 4s 4D 4p 4P ? [O I] F1 2p 4 3P 2p 4 1D [S III] F3 2p 2 1D 2p 2 1S N II d 3D 4p 3P Si II V2 4s 2S 4p 2P N II d 3D 4p 3P [O I] F1 2p 4 3P 2p 4 1D 3 5 [Ni II] s 4F 4s 2D Si II V2 4s 2S 4p 2P Fe II p 4D 4s 2 4D [Ni III] d 8 3F 3d 8 3P ? Fe II V74 4s 4D 4p 4P Fe II s 2 6S 4p 6D [Ni II] s 2F 4s 2P Fe II V74 4s 4D 4p 4P [Fe II] d 7 2G 3d 7 2F N II V8 3s 1P 3p 1P ? Fe II D 5p 6 F 3F 4d 4 H ? [Ni III] d 8 3F 3d 8 3P [N II] V1F 2p 2 3D 2p 2 1D [N II] F1 2p 2 3P 2p 2 1D H H3 2p + 2P 3d + 2D [N II] F1 2p 2 3P 2p 2 1D N II V31 3p 1D 3d 1F O II s 2P 3p 2S [Ni II] d 9 2D 4s 2F He I V46 2p 1P 3d 1D [S II] F2 2p 3 4S 2p 3 2D [S II] F2 2p 3 4S 2p 3 2D [Cr IV] d 3 4F 3d 3 2G [K IV] F1 3p 4 3P 3p 4 1D 3 5 [Ni III] d 8 3F 3d 8 3P [Fe II] s 4D 4s 4F 8 10 [Ni II] s 4F 4s 4P Fe II s 4D 4s 4F [Fe II] d 7 2G 4s 2G 10 10

10 508 Y. Zhang and X.-W. Liu Table 2 continued λ obs Position 1 Position 2 Ion λ lab Mult Lower term Upper term g 1 g 2 F(λ) I (λ) F(λ) I (λ) ? [Fe II] d 7 4F 3d 7 2G [Cr IV] d 3 4F 3d 3 2G [Ni II] s 4F 4s 2D [Ni III] d 8 3F 3d 8 3P [Co III] d 7 4F 3d 7 4P N II d 3P 4p 3P 3 5 [Fe II] s 4D 4s 4F ? ? O I p 3P 4d 3D 5 3 O I p 3P 4d 3D 5 5 O I p 3P 4d 3D 5 7 O I p 3P 4d 3D He I V10 2p 3P 3s 3S [Fe III] d 6 3P 3d 6 1S 3 1 [Ni II] s 4F 4s 4P [Fe III] d 3 3F 3d 3 1D [Fe IV] d 5 4D 3d 5 4F [Ar III] F1 3p 4 3P 3p 4 1D [Fe II] d 7 4F 3d 7 2G [Fe IV] d 5 4D 3d 5 4F 2 4 [Fe III] d 3 3F 3d 3 1D [Fe IV] d 5 4D 3d 5 4F [Fe IV] d 5 4D 3d 5 4F [Fe IV] d 5 4D 3d 5 4F C II V3 3p 2P 3d 2D 4 6 C II V3 3p 2P 3d 2D O I V20 3p 3P 5s 3S 3 3 O I V20 3p 3P 5s 3S 5 3 O I V20 3p 3P 5s 3S 1 3 [Ni II] s 4F 4s 4P He I V45 2p 1P 3s 1S [O II] F2 2p 3 2D 2p 3 2P 6 2 [O II] F2 2p 3 2D 2p 3 2P 6 4 [Ni II] s 4F 4s 2D [O II] F2 2p 3 2D 2p 3 2P 4 2 [O II] F2 2p 3 2D 2p 3 2P [Ni II] d 9 2D 4s 2F [Fe II] d 7 4F 3d 7 2G Position 1: the central emission core The [Fe III] emission lines As in Perinotto et al. (1999) we adopt a 25-level atomic model for Fe 2+. The model includes energy levels from the ground 3d 6 configuration up to the 1 F term, which has an excitation energy of 5.3 ev (cf. Fig. 4). Two levels, 3d 5 ( 6 S)4s 7 S and 5 S, which have excitation energies of 3.73 and 5.08 ev, have been excluded. In Fig. 4, transitions detected in Mz 3 are indicated. The collision strengths ( ij ) are taken from Zhang (1996). The transition probabilities (A ij ) of [Fe III] lines were first given by Garstang (1957). His calculations did not include the configuration-interaction effects. Spectral observations of [Fe III] lines in the PN IC 4997, towards the Galactic Centre and the Orion nebula, yielded line intensities that are inconsistent with those predicted based on the transition probabilities of Garstang (Nahar & Pradhan 1996). More recent calculations of [Fe III] transition probabilities are given by Quinet (1996) and Nahar & Pradhan (1996). Both calculations include a larger set of forbidden transitions and the most important configuration-interaction and relativistic effects. Quinet used the relativistic Hartree Fock method developed by Cowan & Griffin (1976). Nahar & Pradhan used optimized configuration-interaction wavefunctions calculated with the SUPERSTRCTURE program in the Breit Pauli approximation. The large number of [Fe III] lines detected in Mz 3 provides an excellent opportunity to test the theoretical calculations of their transition probabilities. Such calculations are not easy because very accurate wavefunctions are needed for these weak relativistic electric quadrupole and magnetic dipole transitions. For optically thin emission, which is usually the case for forbidden transitions, the intensity ratio of two transitions arising from the same upper level depends only on their spontaneous transition probabilities, but not on the

11 Planetary nebula Mz Table 3. Observed line fluxes at position 3, on a scale where Hβ = 100. λ obs F(λ) I (λ) Ion λ lab Mult Lower term Upper term g 1 g [O II] F1 2p 3 4S 2p 3 2D [O II] F1 2p 3 4S 2p 3 2D He I V2 2s 3S 3p 3P H H7 2p + 2P 7d + 2D [S II] F1 2p 3 4S 2p 3 2P H H6 2p + 2P 6d + 2D H H5 2p + 2P 5d + 2D He I V14 2p 3P 4d 3D N II V5 3s 3P 3p 3P [Fe III] F3 3d 6 5D 3d 6 3F H H4 2p + 2P 4d + 2D He I V48 2p 1P 4d 1D [O III] F1 2p 2 3P 2p 2 1D [O III] F1 2p 2 3P 2p 2 1D N I F1 2p 3 4S 2p 3 2D 4 4 [N I] F1 2p 3 4S 2p 3 2D [N II] F3 3s 1D 3p 1S [O I] F1 2p 4 3P 2p 4 1D [N II] F1 2p 2 3P 2p 2 1D H H3 2p + 2P 3d + 2D [N II] F1 2p 2 3P 2p 2 1D He I V46 2p 1P 3d 1D [S II] F2 2p 3 4S 2p 3 2D [S II] F2 2p 3 4S 2p 3 2D He I V10 2p 3P 3s 3S [Ar III] F1 3p 4 3P 3p 4 1D [Ni II] d 9 2D 4s 2F 6 8 Table 4. Extinction derived from Balmer decrement. Ratio c(hβ) Position 1 Position 2 Hα/Hβ Hγ /Hβ Hδ/Hβ H9/Hβ Adopted nebular physical conditions. In Table 5, the intensity ratios of such pairs of lines observed in Mz 3 are compared to those predicted by the transition probabilities of Nahar & Pradhan (1996) and Quinet (1996). Esteban et al. (1998) measured a number of [Fe III] lines at two positions of the Orion nebula. Their position 1 was centred at 45 arcsec north of θ 1 Ori C and position 2 was centred at 25 arcsec south and 10 arcsec west of θ 1 Ori C. Line ratios derived from their observations are also included in Table 5. Table 5 shows that there is an excellent agreement between the intensity ratios observed in Mz 3 and the theoretical predictions. The agreement between the Orion nebula observations and theoretical values is slightly less satisfactory, probably caused by the larger observational errors in the Orion nebula data set. For all the line ratios tabulated in Table 5, the predictions based on the calculations of Nahar & Pradhan (1996) and Quinet (1996) differ only slightly and are therefore difficult to discriminate using the current observations. Here we will adopt the transition probabilities of Nahar & Pradhan (1996). The many [Fe III] lines detected in Mz 3 from a large number of energy levels of different excitation energies and critical densities pro- Figure 3. Spectra showing the DIB 6284 Å detected in Mz 3. The upper spectrum is for position 1 and the lower one for position 2. vide excellent diagnostic tools to determine the electron temperature and density in the region where the lines are emitted. The ratio of two lines emitted from two separated upper levels of very different excitation energies yields electron temperature, whereas the ratio of two lines arising from two levels of similar excitation but of very different transition probabilities gives electron density. Fig. 5 plots several density-sensitive ratios, λ4754/λ4881, λ4658/λ4667, λ4701/λ4733 and λ4769/λ4777, as functions of log N e for T e = 5000, and K. The λ4754/λ4881 ratio is only sensitive to density variations in the range of 6 < log N e (cm 3 ) < 8. For the latter three ratios, two densities are possible for a given observed value. The degeneracy can be removed based on the result from the λ4754/λ4881 ratio. In the case of Mz 3, the ratio yields a

12 510 Y. Zhang and X.-W. Liu as temperature increases. By combining the density-sensitive ratios plotted in Fig. 5 and the temperature-sensitive ratios in Fig. 6, the electron temperature and density can be determined simultaneously. The [Fe III] line ratios and the electron temperature and density derived from them are presented in Table 6. For the density-sensitive diagnostics, the densities were derived assuming an electron temperature of K, whereas for the temperature-sensitive ratios, the temperatures derived were for log N e (cm 3 ) = 6.5. At the density prevailing here, the densities derived from the various densitydiagnostic line ratios are fairly insensitive to the temperature assumed the results vary by less than 0.1 dex for T e in the range between 5000 and K. [Fe III] λ4667 is blended with the [Fe II] λ4664 line. The contribution of [Fe II] λ4664 to the blend was evaluated and corrected for using the observed intensity of the [Fe II] λ4640 line emitted from the same upper level as the λ4664 line. In Fig. 7, the constraints yielded by [Fe III] diagnostic line ratios measured in the bright central core of Mz 3 are plotted in the T e N e diagnostic plane. The various diagnostic curves converge remarkably well to a density of log N e (cm 3 ) = 6.5 and a temperature of T e = K, with a scatter of 0.1 dex in N e and 1000 K in T e. Figure 4. Energy diagram of Fe III. The solid lines correspond to transitions detected in our spectrum. unique value of log N e (cm 3 ) = 6.52, i.e. the high-density regime applies. Fig. 6 shows several temperature-sensitive [Fe III] diagnostic line ratios as a function of electron temperature, λ6096/λ4667, λ6096/ λ4658, λ7088/λ4607 and λ7088/λ4667, for log N e (cm 3 ) = 6, 6.5 and 7. The line ratios become increasingly sensitive to density Other standard plasma diagnostics The loci of several other standard plasma diagnostic line ratios also measured in the central core in the log N e T e plane are plotted in Fig. 8. The curves were constructed by solving the level populations for multilevel ( 5) atomic models. There is no intersection between the curves of density-sensitive ratios, such as the [S II], [O II] and [Cl III] doublet ratios, and those of temperature-sensitive ratios, such as the [N II], [O II] and [O III] nebular to auroral line Table 5. Comparison of observed and predicted intensity ratios of pairs of [Fe III] lines coming from the same upper level. Line ratio Observed Predicted Transitions Wavelength Mz 3 Orion nebula a Nahar & Pradhan Quinet Position 1 Position 2 (1996) (1996) a 3 F 2 5 b D a 3 F 2 5 D a 3 F 2 5 D a 3 F 2 5 D a 3 F 2 5 b D a 3 F 2 5 D a 3 F 3 5 D a 3 F 3 5 D a 3 F 3 5 D a 3 F 3 5 D a 3 F 3 5 D a 3 F 3 5 D a 3 F 4 5 D a 3 F 4 5 D P 1 5 c D P 1 5 D a 1 D 2 3 P a 1 D 2 3 F a Esteban et al. (1998). b [Fe III] λ4667 is corrected for the [Fe II] λ4664 line (8 per cent), using [Fe II] I (λ4664)/i (λ4640) = c [Fe III] λ5011 is blended with N II lines.

13 Planetary nebula Mz Figure 5. Density-sensitive [Fe III] diagnostic line ratios plotted against log N e (cm 3 ) for T e = 5000 K (dotted line), K (solid line) and K (dashed line). ratios. There is clear evidence that the nebular lines from these ions, all of them having critical densities comparable to or lower than the [Fe III] lines, are suppressed in the high-density environment of the central emission core by collisional de-excitation. In fact, at the high densities prevailing in the dense core, as indicated by the [Fe III] lines analysed above, the [N II], [O II] and [O III] nebular to auroral line ratios have become more of a density diagnostic than a temperature diagnostic. If we assume a constant temperature of K as derived from the [Fe III] lines, then the observed ratios of [S II] λ6731/λ6716, [Cl III] λ5537/λ5517, [O II] λ3729/λ3727, [O II](λ λ7330)/λ3727, [N II](λ λ6584)/λ5754 and [O III] (λ λ5007)/λ4363 in the core yield electron densities of log N e (cm 3 ) = 3.7, 4.4, 4.9, 4.8, 5.4 and 6.5, respectively. There is a clear trend that, the higher the critical densities of the diagnostic lines, the higher the densities that they yield. The dense central emission core in Mz 3 clearly has a highly stratified density structure where nebular lines of lower critical densities are increasingly collisionally suppressed in the high-density regions. While emission of doubly ionized species, such as lines of [Fe III] and [O III], from the dense core of Mz 3 can be represented by a mean electron temperature of K and a mean electron density of log N e (cm 3 ) = 6.5, the average temperature and density under which lines from singly ionized species, such as [N II], [O II] and [S II], arise are not well constrained by the current observations. For example, the observed [N II] nebular to auroral line ratio gives any temperature between 5000 and K for density in the range 5 < log N e (cm 3 ) < 6.4 (Fig. 8). The ratio yields no physically realistic temperatures for the very low densities implied by the observed [S II] and [Cl III] doublet ratios. 4.2 Position 2: the extended inner lobes The plasma diagnostic diagram for the emission region of the extended inner lobes of Mz 3 is given in Fig. 9. In contrast to Fig. 8 for the central core, here all five standard nebular diagnostics, three for density and two for temperature, intersect at a point of approximately T e = 6800 K and N e = 5200 cm 3. The individual values derived from various line ratios are presented in Table 7. In Table 7, the temperatures derived from the temperature-sensitive ratios were calculated assuming a density of 5200 cm 3. In Table 7, we also include the electron temperatures deduced from the ratio of the nebular continuum Balmer discontinuity (Balmer jump, BJ) to H11, i.e. [I c (λ3643) I c (λ3681)]/i (H11), where I c (λ3643) and I c (λ3681) are the nebular continua at 3643 and 3681 Å respectively (Liu et al. 2001a). This yields T e = 6700 K, almost identical to values derived from the [O III] and [N II] forbidden line ratios. A similar analysis for the central core region yields T e (BJ) = K, again in good agreement with T e derived from the [Fe III] forbidden line ratios. 4.3 Spatial variations of N e and T e In Fig. 10 we present the variations of N e and T e across the nebular surface. The electron density was derived from the [S II] λ6731/λ6717 doublet ratio by assuming T e of 6700 and K