Chemical composition of the Orion nebula derived from echelle spectrophotometry

Transcription

1 Mon. Not. R. Astron. Soc. 295, (1998) Chemical composition of the Orion nebula derived from echelle spectrophotometry C. Esteban, 1 M. Peimbert, 2 S. Torres-Peimbert 2 and V. Escalante 2 1 Instituto de Astrofísica de Canarias, La Laguna, Tenerife, Spain 2 Instituto de Astronomía, UNAM, Apdo. Postal , México D.F., México Accepted 1997 November 3. Received 1997 June 13; in original form 1997 March 24 1 INTRODUCTION The Orion nebula is the most observed Galactic H II region. Much work has been devoted to studying the chemical abundances of this object (see, e.g., Peimbert & Torres- Peimbert 1977, Rubin et al. 1991a, Baldwin et al. 1991). Traditionally, the abundance studies for H II regions have been based on determinations from forbidden lines, which are strongly dependent on temperature variations over the observed volume. In contrast, recombination lines are almost independent of such variations and, in principle, they should be more precise indicators of the true chemical abundances. Recently, Peimbert et al. (1993) have derived O /H values from the O II line intensities measured by Osterbrock et al. (1992, hereafter OTV) finding that they ABSTRACT We present echelle spectroscopy in the to 7060-Å range for two positions of the Orion nebula. The data were obtained using the 2.1-m telescope at Observatorio Astronómico Nacional in San Pedro Mártir, Baja California. We have measured the intensities of about 220 emission lines, in particular 81 permitted lines of C, N, N, O 0, O, Ne 0, Si, Si and S, some of them produced by recombination only and others mainly by fluorescence. We have determined electron temperatures, electron densities and ionic abundances using different continuum and line intensity ratios. We derived the He, C and O abundances from recombination lines and find that the C/H and O/H values are very similar to those derived from B stars of the Orion association, and that these nebular values are independent of the temperature structure. We have also derived abundances from collisionally excited lines. These abundances depend on the temperature structure; accurate t 2 values have been derived comparing the O II recombination lines with the [O III] collisionally excited lines. The gaseous abundances of Mg, Si and Fe show significant depletions, implying that a substantial fraction of these atoms is tied up in dust grains. The derived depletions are similar to those found in warm clouds of the Galactic disc, but are not as large as those found in cold clouds. A comparison of the solar and Orion chemical abundances is made. Key words: line: identification ISM: abundances H II regions ISM: individual: Orion nebula. are a factor of 2 larger than the O /H values obtained using forbidden lines. This discrepancy can be interpreted in terms of temperature fluctuations with a t We have taken long-exposure CCD high-spectral-resolution echelle spectrograms to obtain accurate measurements of O II recombination lines and other permitted lines of heavy element ions in the spectrum of the Orion nebula. The main aim of this work is to determine physical conditions based on two sets of highly accurate line intensities. In particular, we want to determine the O abundance from individual recombination lines, avoiding the problem of line blending, and to compare this with the value derived using forbidden lines from the same spectrum. A similar comparison can be performed for the C abundance, but in this case we have to compare our determinations using recombi RAS

2 402 C. Esteban et al. nation lines with those obtained by other authors from the UV C III] λλ1906, 1909-Å emission lines (Walter et al. 1992). 2 OBSERVATIONS AND DATA REDUCTION The observations were carried out with the 2.1-m telescope, in its f/7.5 configuration, of the Observatorio Astronómico Nacional (OAN) at San Pedro Mártir, Baja California, Mexico, in 1994 October and 1995 March. High-resolution spectra were obtained using the OAN echelle spectrograph; its general characteristics have been reported by Levine & Chakrabarty (1994). The spectrograph gives a reciprocal linear dispersion of Å mm 1 at Hα when used with the University College London (UCL) camera. Two different pixel 2 CCD detectors were used: a Thompson CCD chip with a 19 μm 2 pixel size and a Tektronic CCD chip with substantially higher efficiency in the blue with a 24 μm 2 pixel size, for the 1994 and 1995 observations, respectively. At Hα, each pixel corresponds to Å and Å for the Thompson CCD and the Tektronix CCD respectively. The spectral resolution is of 0.5 Å full width at half maximum (FWHM) and the accuracy in the wavelength determination of emission lines is 0.1 Å. We observed two positions on the nebula: Position 1, centred at 45 arcsec N of θ 1 Ori C, and Position 2, centred at 25 arcsec S and 10 arcsec W of θ 1 Ori C. The selection of these positions was made to facilitate comparison with results from previous works (i.e. Peimbert & Torres-Peimbert 1977; OTV; Walter et al. 1992). We also observed θ 1 Ori C, which is the hottest and most luminous star of the Trapezium, and is also the most important ionization source of the nebula. In all the observations, the slit was oriented in the east west direction. A journal of observations is presented in Table 1. We obtained spectra covering two overlapping wavelength ranges for each slit position. The blue spectral range covers from 3550 to 5800 Å in 23 spectral orders, the red one covers from 4500 to 6850 Å or 4600 to 7060 Å, depending on the CCD detector used, both including 17 spectral orders. We had more than 1000 Å of overlap between both wavelength ranges in both slit positions. Three or four individual exposures of 20 min were added to obtain the definitive blue and red nebular spectra. Complementary short 2-min exposures were taken to obtain good intensities for the brightest emission lines, which were close to saturation in the long-exposure spectra. Slits covering arcsec 2 in the red exposures and arcsec 2 in the blue ones were used to avoid overlapping between orders in the spatial direction. We used a Th Ar lamp for wavelength calibration in all the spectral ranges and a tungsten bulb for internal flat-field images. The absolute flux calibration was achieved by taking echellograms of the standard stars indicated in Table 1. All the standards used are from the list of Hamuy et al. (1992), which includes bright stars with fluxes sampled at 16-Å steps. The correction for atmospheric extinction was performed using an average curve for the continuous atmospheric extinction at the San Pedro Mártir site (Schuster 1982). The spectra were reduced using the IRAF echelle reduction package following the standard procedure of bias subtraction, aperture extraction, flat-fielding, wavelength calibration and flux calibration. 3 LINE INTENSITIES AND REDDENING Line intensities were measured by integrating all the flux in the line between two given limits and over a local continuum estimated by eye. In the cases of line blending, a Gaussian profile fit was applied to obtain the line flux of each individual line. All these measurements were made with the SPLOT routine of the IRAF package. As noted in Section 2, the brightest lines Hβ, Hα and the [O III] doublet were close to saturation on the longexposure spectra. As usual, all the line intensities have been normalized to a particular bright recombination line. For the long-exposure spectra, the reference lines were Hγ for the blue range and He I λ5876 Å for the red one. To produce a final homogeneous set of line intensity ratios referred to the same line, the relative line intensities for the long-exposure spectra were rescaled to Hβ by means of the Hγ/Hβ or He I λ5876/hβ ratios obtained from the short-exposure spectra. To check the consistency of the fit, we have compared the intensities of the same set of lines measured in both longer and shorter exposures; in all cases we found negligible differences. The different spectral orders covered in the spectra have overlapping regions at the edges. In these regions the optical sensitivity drops and the line intensities are not as accurate and, therefore, they have not been considered. For the lines in common in two consecutive orders with good flux measurement, i.e. not detected at the edge of the order, the line intensity was the average of the values obtained in both orders. In addition, the blue and red range spectra have a common region of 1300 and 1200 Å for Positions 1 and 2 respectively. The final intensity of a given line in the overlapping region was the average of the values obtained in both spectra. The differences between the relative fluxes of a given line with respect to the same Balmer line obtained in the blue and red spectra do not present any systematic trend that could be related to anomalies in the relative flux calibration. In fact, we estimate that the accuracy of the relative flux calibration achieved among the Table 1. Journal of observations.

3 Chemical composition of the Orion nebula 403 Table 2. Observed and reddening-corrected line ratios [F (Hb)\100] and line identifications.

4 404 C. Esteban et al. Table 2 continued Dubious identification (see text). different orders is about 2 per cent from the comparison of pairs of well-measured fluxes of the same line and their underlying continua. The final list of observed wavelengths and line intensities relative to Hb for both slit positions is presented in Table 2. The observational errors associated with the line flux intensities (including all the possible sources of uncertainties in line intensity measurement and flux calibration) are estimated to be 2 5 per cent ( dex) if the ratio F (l)/f (Hb) is of the order of or larger than 0.1, about 10 per cent (0.04 dex) when 0.01sF (l)/ F (Hb)s0.1 and about 20 per cent (0.08 dex) when 0.001sF (l)/f (Hb)s0.01. For the lines weaker than 0.001F (Hb), the uncertainty could be of the order of or larger than 30 per cent (0.12 dex). Colons indicate uncertainties of the order of or larger than 40 per cent (0.16 dex). For a given line, the observed wavelength is determined by the centre of the baseline chosen for the flux integration procedure or the centroid of the line when a Gaussian fit is used. For a line observed in different spectra and orders, we have adopted the average wavelength of the different measurements. The final values given in columns 4 and 7 of Table 2 are observed wavelengths relative to the local standard of rest. The identification and adopted laboratory wavelengths of the lines collected in Table 2 were obtained following previous identifications in the Orion nebula by Kaler et al. (1965) and OTV, the compilations of atomic data by Moore (1945, 1993) and Wiese et al. (1966) and the papers of Hyung, Aller & Feibelman (1994a,b) based on echelle spectra of bright planetary nebulae. Several lines identified as sky emission were not included in the list. About 10 emission lines, some measured in both slit positions, could not be identified in any of the available references. Some of a

5 Chemical composition of the Orion nebula 405 Table 3. Unidentified lines reported in this and previous works. a Position 2. b Average of Positions 1 and 2. these unidentified lines had been also reported previously. In Table 3 we list four unidentified lines that have been observed in the Orion nebula and/or other gaseous nebulae by different authors. The first column gives the observed wavelength in this work corrected for radial velocity, and the second column gives the wavelength reported by other observers. The third and fourth columns include the object and the reference where the line is reported in the literature. The last column contains the estimated rest wavelength of the unidentified lines. It is well known that the Orion nebula, and specially the Trapezium region, shows an anomalous extinction law. The reddening function, f (λ), normalized at Hβ for this nebula was derived by Costero & Peimbert (1970). The use of this reddening law instead of the standard one (Whitford 1958) would produce a different C (Hβ) but very similar line intensity ratios in the optical region. The reddening coefficient, C(Hβ), was determined by fitting iteratively the observed Balmer decrement to the theoretical one computed by Storey & Hummer (1995) for the physical conditions of each nebular position. The final C (Hβ) for each slit position was the average of the values obtained from the Hα/Hβ, Hγ/ Hβ and Hδ/Hβ ratios in each of the short-exposure spectra. The final values of C (Hβ) are and for Positions 1 and 2 respectively. These values are slightly lower than those (0.63 and 0.59) reported by Peimbert & Torres-Peimbert (1977) and OTV for positions almost coincident with our Position 1, and also lower than the value of 0.60 given by Peimbert & Torres-Peimbert (1977) for the place of our Position 2. Table 2 shows the reddening-corrected line intensity ratios, I (λ)/i (Hβ), for each line and slit position. The reddening-corrected Hβ line fluxes from the blue short-exposure spectra for the arcsec 2 slit are and erg cm 2 s 1 for Positions 1 and 2, respectively. In Fig. 1 we compare our line intensity ratios with those of Kaler et al. (1965) and OTV by means of least-squares fits. The comparison with the photographic data of Kaler et al. (1965) shows a slope of 0.919, indicating a systematic trend to overestimate the intensity of the weakest lines in their data. The overestimation of weak lines from photographic data has been known for a long time (e.g. Miller 1971). In contrast, the comparison with the data of OTV gives a slope Figure 1. Comparison of line intensity ratio without correction from reddening from this work with those of (a) Kaler et al. (1965) and (b) Osterbrock et al. (1992). of 0.976, much closer to 1, as would be expected for two different sets of CCD observations. Several line intensities given by OTV appear to be slightly higher than ours, especially the observed line intensity of N II λ Å, which is almost four times higher in OTV observations. Our data provide (i) line ratios of higher accuracy than those of Kaler et al. (1965), because of the linearity of CCD detectors, and (ii) higher spectral resolution than that present in OTV. 4 PHYSICAL CONDITIONS The large number of emission lines identified and measured in the spectra allows the derivation of physical conditions using different line ratios. The values of N e and T e given in Table 4 have been obtained using the five-level program for the analysis of emission-line nebulae of Shaw & Dufour (1995), except in the cases of the N e derived from [Fe II] and [Fe III] lines. For λλ4244, 4416, 4814, 5159 and 5262 Å of [Fe II], we have constructed diagnostic diagrams of line intensity ratios versus N e based on the computations by Bautista, Peng & Pradhan (1996). For [Fe III] we have used the diagnostic diagrams published by Keenan et al. (1993). The N e value derived from Ar 3 for Position 1 is highly

6 406 C. Esteban et al. Table 4. Physical conditions. uncertain and close to the low density limit, so it is not presented in Table 4. The values of N e obtained for the different ions available in each slit position are very similar, except for [Fe II]. An average value of 4000 cm 3 has been determined as representative for Position 1, coincident with the value obtained by OTV for this zone. For Position 2, the values obtained with the different ions except again those of Fe are remarkably consistent, and an average of 5700 cm 3 has been assumed as representative of that zone. A similar value of 5000 cm 3 was obtained by Peimbert & Torres-Peimbert (1977) for their position Ori 2b, coincident with our Position 2. The fairly similar values of N e obtained using forbidden line ratios of ions with different ionization potentials from S to Ar 3 suggest that the density does not vary considerably along a substantial part of the nebula. The densities from [Fe III] lines show a relatively large dispersion, between 10 3 and 10 4 cm 3, but are broadly consistent with the results obtained from the line ratios of the other ions. This consistency indicates that the [Fe III] lines are formed at the same places as the rest of the forbidden lines of the other ions. Fe is qualitatively different; the densities obtained with the different diagnostic diagrams give a very high dispersion. Bautista, Pradhan & Osterbrock (1994) find evidence of high-density regions in this nebula that could be responsible for these differences. These authors obtained N e cm 3 from diagnostic diagrams involving optical and near-infrared (near-ir) [Fe II] lines. As complementary evidence to support the existence of these high-density regions, Bautista et al. (1994) found that the Fe abundances derived from the near-ir and optical lines become consistent if large densities are considered. In the case of our results for this ion, we have to indicate that most of the diagnostic diagrams constructed become insensitive to N e at values around 10 6 cm 3. The observed ratios are usually near the high-density asymptotic regions in the diagnostic diagrams. The [N I] ratio is also close to the high-density limit and the N e derived from it is very uncertain. Electron temperatures from forbidden lines have been derived from [O III], [N II] and [S II] line ratios. The values obtained from different ions are systematically different in both slit positions, especially those of [S II]. T ([N II]) Figure 2. Section of the echelle spectrum of Position 1 at the Balmer limit region (observed fluxes). appears to be at least 1500 K higher than T([O III]). Baldwin et al. (1991) have observed this trend in several positions along the Orion nebula. Both temperature determinations are based on lines produced by ions with very different ionization potentials that occupy different zones of the nebula. On the other hand, the T([S II]) values are remarkably higher than the other T e values. The strong difference between the T ([S II]) derived in Positions 1 and 2 appears to be real, and comes from the diference of more than a factor of 2 in the intensity of the auroral [S II] λ4069, 4076-Å lines measured in both positions. The large value of T e in Position 2 could be caused by the presence of a shock or material with densities considerably higher than 10 5 cm 3 (the critical densities are above cm 3 ). For the S II diagnostics we have used the atomic parameters computed by Keenan et al. (1996). 5 RECOMBINATION SPECTRUM OF HI AND HeI 5.1 H I lines The H I Balmer spectrum can be detected up to H30 in our data. In Fig. 2, we show the spectrum of Position 1 at the Balmer limit region. The Balmer discontinuity and the high principal quantum number Balmer lines are very conspicuous in this figure. As is well known, the H I spectrum arises essentially from recombination, presenting very small deviations caused by the destruction of Lyman-line photons by dust, which affect only the Hα/Hβ ratio by a few per cent (Baldwin et al. 1991). The best calculations of the H I recombination spectrum are those by Hummer & Storey (1992) and Storey & Hummer (1995). In Table 5, we present the comparison between the observed intensities of the H I Balmer lines and the predicted ones using the machine-readable line-ratio tables by Storey & Hummer (1995). The theoretical line ratios have been evaluated for T e 9000 K and N e 4000 cm 3 for Position 1 and T e 9500 K and N e 5700 cm 3 for Position 2 (the H I line ratios are almost independent of the adopted temperatures and densities). The comparison cannot go further than H25 because this is the limit of the calculations of Storey & Hummer (1995). H8 and H14 have not been

7 Chemical composition of the Orion nebula 407 Table 5. Observed over predicted (Storey & Hummer 1995) H I Balmer emission line ratios. Table 6. Observed over predicted (Smits 1996) He I line ratios. included in Table 5 because they are blended with He I Å and [S III] Å, respectively. Columns 2 and 3 give the ratio of the observed to the predicted intensities. The differences between theoretical and observed line ratios are typically of the order of or below 10 per cent, which implies that the accuracy of the line ratios is of that order. 5.2 Balmer discontinuity temperature The echelle spectra obtained allow a satisfactory determination of the Balmer discontinuity, defined as I c (Bac) I c (λ3646 ) I c (λ3646 ), where I c (λ3646 ) and I c (λ3646 ) are the nebular continuum fluxes shortward and longward of the Balmer limit at λ3646 Å. In our case, we have obtained I c (λ3646 ) as the mean value of the continuum in the line-free region between λλ3615 and 3630 Å. The corresponding I c (λ3646 ) has been obtained from the mean value of the narrow regions of continuum between H18 and H22 (from 3676 to 3691 Å) in the spectral order shown in Fig. 2. Liu et al. (1995a) have studied the effect of stellar scattered light and sky background on their spectral observations. In our case, the higher spectral resolution of the observations and the closeness of the spectral regions used to obtain the continuum fluxes shortward and longward of the Balmer limit only 50 Å apart makes the possible effects quantified by these authors negligible. Moreover, this part of the continuum is not affected by the line blending that produces the raising of the apparent continuum just longward of the Balmer limit. We have obtained power-law fits to the relation between I c (Bac)/I(Hn) and T e for 3 n 21. The emissivities as a function of electron temperature for the nebular continuum and the H I Balmer lines are taken from Brown & Mathews (1970) and Storey & Hummer (1995) respectively. The T e (Bac) finally adopted for each slit position was the average value obtained using the different H I Balmer lines. These adopted temperatures are shown in Table 4. The error associated with T e (Bac) corresponds to the dispersion of the different temperatures obtained using the H I Balmer lines considered, which amounts to 500 K, plus an estimate of an additional error associated with the determination of I c (Bac). 5.3 He I lines There is a large number of He I emission lines identified in our spectra. These lines arise mainly by recombination, although they may have contributions from collisional and self-absorption effects. In Table 6, we present the comparison between the observed line intensity ratios and those predicted by Smits (1996). The theoretical line ratios have been evaluated for case B for singlets and triplets assuming T e 9000 K and N e 4000 cm 3 for Position 1 and T e 9500 K and N e 5700 cm 3 for Position 2. The ratios between observed and predicted line intensity ratios are presented in columns 3 and 4 of Table 6. As in the case of the H I Balmer spectrum, the observed ratios compare well with the predictions, taking into account that uncertainties of the order of 10 and 20 per cent are expected as a result of the relative faintness of most of the He I lines. An apparent trend is detected in the 2 3 P n 3 S lines: the observed intensities are systematically brighter than the predicted ones by more than 20 per cent. This effect can be explained by selfabsorption effects from the metastable 2 3 S level, which produce large optical depths in the lines for the 2 3 S n 3 P transitions. The optical He I lines most affected by selfabsorption are λλ3889, 7065 and 10830, but, unfortunately, we do not have measurements of the last two lines and the first one, He I 3889 Å, is blended with H8 in our data. We have compared the predicted intensity for the sum of He I 3889 and H8 with the observed value, finding that the combined line flux observed only accounts for 72 and 76 per cent of the predicted flux for Positions 1 and 2, respectively. The most likely explanation for this low intensity in the blend is the effect of self-absorption on He I 3889 Å. In this case, the observed I(He I 3889)/I (He I 4471) value becomes 41 and 49 per cent of the theoretical one for Positions 1 and 2, respectively. The intensities given by OTV for the blend are also consistent with our values. Moreover, OTV find a very

8 408 C. Esteban et al. high intensity for He I 7065 Å. These results indicate the presence of significant optical depth effects in the helium triplet spectrum of the nebula (e.g. Robbins 1968); in fact, OTV find τ(3889 Å) of the order of 12. On the other hand, the ratios presented in Table 6 indicate the absence of significant line-transfer effects in the helium singlet spectrum (e.g. Robbins & Bernat 1973). 6 IONIC ABUNDANCES FROM FORBIDDEN LINES Ionic abundances of O, O, N, Ne, S, S, Cl, Ar and Ar 3 have been obtained from collisionally excited lines, using the five-level atom program of Shaw & Dufour (1995) and the atomic parameters referenced in it. We have assumed a two-zone scheme and t 2 0, adopting T ([N II]) for the low-ionization-potential ions O, N and S, and T ([O III]) for the high-ionization-potential ions O, Ne, S, Cl, Ar and Ar 3. The densities assumed are 4000 and 5700 cm 3 for Positions 1 and 2, respectively. The ionic abundances derived are listed in Table 7. To derive the S abundance we make use of the auroral lines which, unlike the nebular ones, are not affected by collisional de-excitation in high-density regions (N e 10 6 cm 3 ). We have a similar situation for deriving the O abundance but, in this case, we have made use only of the nebular [O II] λ3726, 3729-Å lines because the auroral lines are outside the spectral range covered in our observations. The Ar abundances were derived from the λ5192-å auroral line, which is very weak, so these abundances are more uncertain than those derived by other authors from the considerably brighter λ7136-å nebular line. 6.1 [Fe II] spectrum and Fe abundance 18 [Fe II] lines belonging to mutiplets (4F), (6F), (7F), (17F), (18F), (19F), (20F) and (21F) have been identified and measured in both slit positions. The presence of many [Fe II] lines in the spectrum of the Orion nebula was previously reported by Kaler et al. (1965). Moreover, OTV Table 7. Ionic abundances from forbidden lines in units of 12 log(x m /H ). analysed optical and near-ir [Fe II] emission lines to obtain the Fe abundance. These authors constructed a simple scheme to obtain the level populations of this ion instead of solving the very large number of equilibrium equations necessary for the full solution. Recently, Bautista & Pradhan (1996) have developed new extensive collisional radiative models for [Fe II], including new collisional excitation rates and radiative transition probabilities for a large number of transitions and atomic levels of this ion. In our case, we have interpolated this set of data to derive the Fe abundances. We have assumed T([N II]) as the most appropriate electron temperature for this low-ionization-potential ion and two different values of N e for each slit position; in the first case, we have assumed the densities obtained from forbidden lines of the other ions (4000 and 5700 cm 3 for Position 1 and 2, respectively) and in the second case we have assumed a density of 10 6 cm 3 as suggested by the diagnostic diagrams involving [Fe II] line ratios. T ([S II]) is probably not appropriate for the region where Fe originates because the [S II] lines might be affected by a lowvelocity shock component or by a high-density component. The comparison between the observed and predicted emission line ratios and also the Fe /H values derived are presented in Table 8. The results obtained using the lower and higher value of the electron density are separated by a slash in several columns of Table 8. We have selected [Fe II] 5262 Å as a reference line because it is a single line, well measured in both slit positions, that belongs to a multiplet presenting several measurable lines in our spectra. Columns 5 and 8 of Table 8 show the ratio between observed and predicted values. The large deviation that some lines present with respect to the theoretical values is remarkable. The difference is especially important for the single line of the multiplet (4F) observed. Fluorescence excitation is unlikely to occur for this line because the quartet levels cannot be pumped by dipole allowed transitions from the 6 D ground state. Moreover, a hypothetical fluorescence excitation would also affect other lines of the multiplet, which should be brighter and measurable. We therefore consider that this line is a misidentification. On the other hand, multiplet (7F) presents a very clear effect of fluorescence excitation. The four observed lines of this multiplet show very similar enhancements relative to the predicted intensities. For Position 1, the average enhancements are of the order 7 and 14 in the low- and high-density cases, respectively. For Position 2, the enhancements appear to be slightly larger, a factor of 11 and 21 for low and high densities, respectively. In fact, in a very recent work, Bautista et al. (1996) predict that multiplet (7F) should suffer the largest fluorescence effects in either optical or IR [Fe II] lines because the atomic levels involved in this transition are sextets and, therefore, pumping from the 6 D ground state can be produced by allowed dipole transitions. The intensity of an emission line affected by fluorescence from the continuum stellar emission is proportional to the intensity of the local radiation field; the fact that the apparent fluorescence effect is larger in Position 2 can therefore be naturally explained, because this position is closer to θ 1 Ori C, the brightest star of the Trapezium. All the other multiplets of [Fe II] observed do not show apparent fluorescence excitation effects, although the correspondence between observed and theoretical line

9 Chemical composition of the Orion nebula 409 Table 8. Comparison of observed over predicted [Fe II] line ratios and Fe abundances. The numbers separated by a slash represent the values obtained for different electron densities (4000 and 5700 cm 3 for Positions 1 and 2, respectively) and 10 6 cm 3. ratios is not very good in most cases. The use of low or high densities does not appear to increase the global consistency between observations and theoretical calculations. The similarity between the intensity measured for the same lines in both slit positions suggests that the discrepancy is real. The Fe abundances derived do not show a substantial scatter if multiplets (4F) and (7F), as well as the lines with large observational errors marked with colons in Table 8, are not considered. In this case, the average values of N (Fe )/N(H ) obtained are ( ) 10 7 and ( ) 10 7 for Positions 1 and 2, respectively, in the low-density case, and ( ) 10 7 and ( ) 10 7 in the high-density case (10 6 cm 3 ). These values are consistent with those obtained previously by Bautista et al. (1994), who make a re-assessment of the Fe abundance based on OTV data applying complete calculations for level populations. These authors find that the scatter in the Fe abundance obtained from the different [Fe II] lines they observed, defined as σ x / x, where σ x and x are the standard deviation and the mean value of the abundance, respectively, becomes lower when densities of the order of 10 6 cm 3 are considered. We have performed the same calculation for our data, again not considering multiplets (4F) and (7F) and those line intensities with colons, and we find only marginal differences, σ x / x being slightly larger for N e For Position 1 σ x / x is 0.62 and 0.75 when we consider low and high densities, respectively. For Position 2, σ x / x yields 0.49 and 0.58 for each density. We therefore do not find compelling evidence in favour of the high-density case. 6.2 [Fe III] spectrum and Fe abundance A number of 19 and 17 [Fe III] lines of multiplets (1F), (2F) and (3F) have been measured in Positions 1 and 2, respectively. The presence of a large number of [Fe III] lines in the spectrum of the Orion nebula was previously reported by Kaler et al. (1965). OTV derived the Fe /H ratio based on 12 [Fe III] lines. We have used the tables of level populations for Fe III computed by Keenan et al. (1992) to obtain the theoretical line ratios and the Fe abundance. We have used T ([N II]) as representative for Fe and the electron densities assumed are 4000 and 5700 cm 3 for Positions 1 and 2, respectively. The comparison between the observed and predicted emission line ratios, as well as the Fe /H values, is presented in Table 9. From the fifth column, the consistency between the observed and the predicted line ratios can be noted, except in the cases where line-blending is present. The good agreement between theoretical and observed line ratios implies little scatter in the Fe /H ratios derived. The average Fe /H is ( ) 10 7 and ( ) 10 7 for Positions 1 and 2, respectively. OTV obtain a similar value, Fe /H ( ) 10 7, using their spectroscopic data and the level populations calculated by Garstang et al. (1978). 6.3 [Ni II] and [Ni III] spectrum There is a single [Ni II] emission line in our spectra λ6366 Å, one of the brightest lines of multiplet (8F), corresponding to

10 410 C. Esteban et al. Table 9. Comparison of observed over predicted [Fe III] line ratios and Fe abundances. the a 4 F 3 a 4 P 2 transition. On the other hand, OTV measured two lines of the (2F) multiplet, with wavelengths out of our observed spectral range, finding very large Ni /H ratios. Lucy (1995) and Bautista et al. (1996) have studied the excitation mechanism of the lines of multiplet (2F), finding that fluorescence excitation by the UV stellar continuum can produce the observed behaviour. In our case, we have derived the Ni abundance from the single line measured by interpolating the most recent calculations of line emissivities and ratios provided for this ion by Bautista et al. (1996). The physical conditions adopted are the same as those for Fe. With all these ingredients, we derive N(Ni )/ N (H ) and , for Positions 1 and 2 respectively, larger than the value of given by OTV. As the lines used by OTV probably suffer from strong fluorescence effects, the large Ni abundances obtained from [Ni II] λ6366 Å also suggest a strong enhancement in its intensity. As this line comes from a transition between quartet levels, it does not include a fluorescence effect from the a 2 D ground level. Three weak [Ni III] lines of multiplet (2F) have been measured in our spectra, two of them observed in both slit positions and presenting very similar intensities in both spectra. One of these lines, λ Å, has been previously detected by Thackeray (1975). OTV measure two brighter lines of [Ni III] in the red part of the spectrum and outside our spectral coverage. Unfortunately, there are no calculations of the level populations of this ion because no published collision strengths are available for [Ni III]. OTV estimate these atomic parameters from available calculations for [Fe VII], with an analogous electron ground configuration. These authors calculate the Ni /H ratio using the collision strengths estimated as above and the transition probabilities obtained by Garstang (1958). Of the three lines of multiplet (2F) that we measured, there is a broad consistency between the intensities we found for λλ6000 and 6534 Å and those predicted by OTV, but λ Å is a factor of 4.5 larger than predicted. 7 He ABUNDANCE As was noted in Section 5.3, the intensity of He I emission lines may have contributions from collisional and selfabsorption effects. The comparison between observed and predicted emission lines shown in Table 6 indicates that only He I lines coming from 2 3 P n 3 S transitions could present substantial self-absorption effects and, therefore, that they are not suitable for deriving an accurate He /H ratio. On the other hand, collisional effects may affect all the measured lines, but, unfortunately, we have collision-torecombination factors (C/R) for only 6 of the lines available. In Table 10, we present the He abundances for these 6 lines uncorrected and corrected by collision contribution as well as their corresponding C/R factors. The uncorrected He /H ratios have been obtained using the pre- Table 10. He /H abundance ratios.

11 Chemical composition of the Orion nebula 411 dicted line emissivities calculated by Smits (1996) and assuming the same physical conditions as in Section 5.3. The C/R factors have ben obtained from the most recent calculations by Kingdon & Ferland (1995), which are based on new collision strengths from a 29-state quantal calculation of He I extending to n 5. The appropriate values of C/R have been derived assuming the same physical conditions as those used to obtain the uncorrected He /H ratios. The average He /H values for each slit position, as well as their associated statistical dispersion, are given in Table 10. These values compare well with the results obtained by OTV (0.089) and those given by Peimbert & Torres-Peimbert (1977) for their Positions 1a and 2b, which amount to and 0.089, respectively. 8 PERMITTED HEAVY-ELEMENT LINES We have measured 81 permitted lines of heavy-element ions such as O I, O II, C II, Ne I, S II, N II, N III, Si II and Si III, many of them at both slit positions; of these, 69 were detected by Kaler et al. (1965) and 22 by OTV. 8.1 Excitation mechanisms The excitation mechanisms of the permitted lines observed in the Orion nebula and other gaseous nebulae have been extensively studied by Grandi (1975a,b; 1976). For the Orion nebula, this author uses the compilation of line intensities by Johnson (1968), based on photographically determined line strengths. The collection of permitted lines observed in our echelle spectra is shown in Table 11. In this table we include the ion, the multiplet number, the transition associated with each multiplet, the laboratory wavelength of each line and its most probable excitation mechanism. The advantage of our set of data is that it contains high-quality line strength measurements of very faint lines with enough spectral dispersion to resolve several in some cases all lines of the same multiplet. An accurate determination of the excitation mechanism of some lines requires realistic model computations of the nebula that give details of the local radiation field and take into account all the processes of absorption and recombination to excited states of heavy elements simultaneously. We will present such computations for ions for which the neces- Table 11. Permitted lines of heavy-element ions in the Orion nebula and their most probable excitation mechanism.

12 412 C. Esteban et al. sary atomic data are available in a separate paper (Escalante et al., in preparation). Nevertheless, many lines that we have observed are clearly produced by recombination, and ionic abundances can be calculated directly from their observed intensities. We have measured 10 O I emission lines, the majority of them at both slit positions. Individual components of multiplets are not resolved for neutral species because of the small fine-structure splitting of their terms. The intensities of the O I lines in Table 2 therefore correspond to the total intensity of the multiplet. The excitation mechanism of the O I lines is mainly fluorescence resulting from the stellar continuum radiation, as shown quantitatively by Grandi (1975a, b). The upper levels of the lines of multiplets 21 to 27 can be excited directly by absorption of stellar photons longward of 912 Å from the ground term, followed by permitted downward transitions. In the case of multiplet 5, corresponding to a 3s 3 S 0 4p 3 P transition, the population of the upper level can be fed from transitions from s and d levels of higher principal quantum number. Therefore, starlight excitation may contribute largely to the observed strengths of this line. On the other hand, the identification of O I λ4803 Å is dubious because it is a transition between quintet terms that cannot be produced by fluorescence from the ground state. Excitation by recombination is also unlikely, because in that case other optical quintet lines with larger statistical weights, such as multiplet 18 at λ Å (3p 5 P 8d 5 D o ), which we did not detect, should be more intense than λ4803 Å. The line at λ4803 Å is very weak and is very close to the N II λ å line; thus a misidentification or confusion with the N II emission feature is possible. We have another line conflict, O I λ Å, which takes place between the 3p 1 G o core-excited levels with energies 0.5 and 2.5 ev above the Lyman limit. If strict LS coupling holds, however, production of O I singlets by fluorescence or dielectronic recombination at low densities is not possible because of the spin selection rule. This line is rather weak and only observed at Position 2; it could therefore be a misidentification. We have measured a large number of O II lines in our spectra. We have obtained good intensities for almost all individual lines of multiplets 1 and 2. In Fig. 3, we show a section of the echelle spectrum for Position 1, containing the individual emission lines of multiplet 1 of O II, to illustrate the good quality of the spectra. The absence of blends and the relatively high signal-to-noise ratio (S/N) of the O II lines in this spectral range are evident from the figure. The excitation mechanism of the O II spectrum was investigated by Grandi (1976). The presence of many emission lines cannot be explained by resonance fluorescence. Moreover, Grandi (1976) estimates that starlight excitation of the O II λ430-å line contributes only 20 per cent as much as recombination to multiplet 19 line intensities, the effect being considerably smaller for the rest of the lines. We have observed lines of several multiplets not analysed previously by Grandi (1976): 53, 66, 67, 68 and 93.01, all of them corresponding to 4f 3d transitions; these lines should also be excited by recombination because 4f levels cannot be excited by starlight from the 2p 34 S 0 ground level. The measurement of several lines of the same multiplet permits comparison of their relative intensities with those expected in an appropriate angular momentum coupling Figure 3. Section of the echelle spectrum of Position 1 showing all the individual emission lines of multiplet 1 of O II (observed fluxes). under the assumption that the populations of the fine-structure levels within a term are proportional to the statistical weights. The results of this comparison are shown in Table 12. In this table we include those multiplets with more than one emission line measured. Columns 1 and 3 indicate the ion, multiplet number and laboratory wavelength corresponding to each line. Column 4 gives J values of the lower and upper levels of the transitions. Columns 5 and 7 give the observed intensity of each line relative to the strength of the brightest line of the multiplet. Columns 6 and 8 include the ratios of the observed and the predicted relative intensities. For O II there is good agreement between the observations and the predictions of LS coupling for a significant fraction of the lines. The strengths of λλ and Å of multiplet 1 are discrepant by a similar fraction at both slit positions, and are the only observed lines involving the 3s 4 P 1/2 level in that multiplet. They do not present any apparent line-blending, as can be seen in Fig. 3. It is worth noting that Liu et al. (1995b) observed a similar difference for λ Å in their observations of the planetary nebula NGC For multiplet 2, λ Å differs by a factor larger than 2 from its predicted intensity, and this line apparently does not suffer from line-blending. The relative intensities of the two lines of multiplet 19 show important deviations from their predicted values. Liu et al (1995b) and Wiese, Fuhr & Deter (1996) have pointed out important departures from pure LS coupling for 3d 3p and 4f 3d transitions in O II. Both works give intermediatecoupling quantal calculations of multiplet 19 with similar results. In Table 12 we give the predictions of LS (multiplets 1 and 2) and intermediate coupling (multiplets 10, 15 and 19) of Liu et al. (1995b). It is important to note that the component of multiplet 19 at λ Å is blended with the He I line at λ Å. However, the observed 4156/4153 value is closer to the LS coupling predictions than to the intermediate-coupling ones. Liu et al. reported a similar situation for NGC Multiplet 15 (3s 2 D 3p 2 F o ) is the only multiplet with a 1 D core that we detected. Nussbaumer & Storey (1984) have predicted similar or greater dielectronic rates for other multiplets that we did not detect, like

13 Table 12. Comparison of the observed and predicted intensity ratios of permitted lines of heavy elements. a Values for LS coupling (left) and intermediate coupling (right). b Blended with He I λ Å. c Possible contamination with O II λ Å. d Possible contamination with [Fe III] λ Å. e Blended with O I λλ , , Å. Chemical composition of the Orion nebula 413 multiplet 16 (3s 2 D 3p 2 D o ) and multiplet 36 (3p 2 F o 3d 2 G). Grandi (1976) has studied the excitation mechanism of C II lines in gaseous nebulae belonging to multiplets 3, 4 and 6. We have observed 6 C II lines in our spectra, corresponding to multiplets 2, 4, 5 and 6. We have identified the line at λ Å with multiplet 60 of He I rather than multiplet of C II λ3871 Å, because the latter is a spin-forbidden transition. The most probable dominant excitation mechanism of these lines is shown in Table 11. Kaler (1972) found that multiplet 4 should be affected by strong starlight excitation because of the behaviour of the ratio of C II λ3921 to λ4267 Å relative to the effective temperature of the ionizing star. Grandi (1976) demonstrated this fact quantitatively, finding that the contribution of resonance fluorescence by starlight can be important for transitions with upper 2 S and 2 D levels. In particular, the contribution to the strength of C II λ3921 Å resulting from starlight is 2.5 times greater than recombination and, consequently, other lines of multiplet 4 should be similarly affected. In Table 12, we show that the relative intensities between the two observed lines of multiplet 4 are in very good agreement with the predictions of LS coupling. Grandi (1976) showed that recombination dominates the excitation of λ Å by an order of magnitude. This line comes from a transition involving terms with large l values, levels that cannot be reached by permitted resonance transitions from the ground term and, therefore, are excited mainly by recombination. On the other hand, C II λ Å has 3p 2 P 0 as upper term and can be populated from transitions excited by both resonance fluorescence by starlight (coming directly from 2 S and 2 D levels) and recombination from terms having large L values. Finally, λ5892 Å (3d 2 S 4p 2 P 0 ) has an upper term that can be populated mainly from 2 S and 2 D and, hence, the dominant mechanism should be resonance fluorescence by starlight, with a contribution from recombination. The excitation mechanism for multiplet 7 of S II is very uncertain. Recombination should produce other multiplets with greater intensity like multiplets 44 (4p 4 D o 4d 4 F) or 45 (4p 44 D o 4d 4 D). Continuum fluorescence from the ground state should also excite other transitions like multiplet 9 (4s 4 P 4p 4 S o ), but no other S II multiplets were detected. Grandi (1976) has shown that recombination is insufficient to account for the observed triplet line strengths of N II in the Orion nebula. Resonance fluorescence by the recombination line He I λ Å is the dominant mechanism to excite the 4s 3 P 0 term and hence is responsible for the strength of multiplets 3, 5 and probably 4. Grandi (1976) suggests that multiplet 28 could be excited by a combination of recombination and starlight. This suggestion can also be applied to multiplet 20, which has the same upper term as multiplet 28. The only observed line of multiplet 29, N II λ Å, corresponds to a 3p 3 P 3d 3 P 0 transition and is blended with [Fe II] λ Å. Resonance flourescence by starlight seems to be less appropriate to explain the intensity of this line because the population of the upper term should come from upper starlight-excited 3 P or 3 D terms. If the identification of N II λ Å is correct, therefore, a combination of recombination and starlight should be the dominant excitation mechanism. The λ å line corresponds to a transition between quintet terms. Neither fluorescence from the ground state nor recombination can

14 414 C. Esteban et al. produce such states in normal nebular conditions, and therefore it has probably been misidentified. In Table 12, we compare the relative intensities of the N II lines of the same multiplet with those predicted in LS coupling. Multiplets 3 and 5 present ratios between observed and predicted strengths close to unity, indicating that the population of these levels is roughly proportional to their statistical weights. In multiplets 20 and 28 there are large differences with the predicted relative intensities in LS coupling. The lines arising from the 3d 3 D o 3 level tend to be much weaker than expected relative to other lines of the same multiplet. We can also compare the line intensities of one multiplet with the other because both multiplets have the same upper term and can be assumed to be optically thin. Their relative intensities depend only on the transition probabilities and wavelengths, and can be compared with the predictions of atomic theory. Table 13 shows observed and predicted ratios of lines with the same upper level. The differences between the observed and theoretical line ratios for both positions are smaller than the observational errors, but the possibilities of a breakdown of LS coupling for the 3d 3 D o levels, or the levels not being proportional to their statistical weights, cannot be ruled out. Two N III lines of multiplet 2 are observed in the spectrum of Position 2. Grandi (1976) remarks that recombination is incapable of producing the observed line intensities because N 3 is not expected in significant amounts in the Orion nebula. In Table 12 we show that the observed ratio between the two N III lines of multiplet 2 are in good agreement with the predictions of LS coupling. We have measured nine Si II lines in the 2 slit positions. These lines belong to multiplets 1, 2, 4, 5 and 8, the first four corresponding to doublets and the last one to a quartet. Grandi (1976) has studied the excitation mechanism of the doublet lines, finding that starlight excitation of the 5s 2 S and 4d 2 D terms dominates over recombination by two orders of magnitude and reproduces the observed line strengths. If the only quartet line observed, λ Å, is correctly identified, it should be produced by recombination. However, as this line is not the most intense of multiplet 8, its identification is uncertain. Table 12 shows the good agreement between the observed intensity ratios and the LS-coupling predictions for all the observed doublet multiplets of Si II. Because multiplets 1 and 2 have the same upper term, their line ratios have been compared in Table 13 with theoretical values from the Opacity Project data base. Two Si III lines have been observed at both slit positions. The singlet one, λ Å, has been analysed by Grandi (1976), finding that its strength is produced by a combination of starlight excitation and recombination. The second line, λ Å, corresponds to a triplet involving high angular momentum states that should be produced by recombination. Nevertheless, this line is predicted to be the faintest of multiplet 9, and as two other brighter lines of this multiplet are not observed, the line has probably been misidentified. 8.2 Ionic abundances Only permitted lines produced mainly by recombination can give accurate determinations of ionic abundances, because under nebular conditions their relative intensities depend Table 13. Line ratios with common upper levels. a Nussbaumer & Storey (1984). b Victor & Escalante (1988). c Luo & Pradhan (1989). d Bell, Hibbert & Stafford (1995). e Cunto et al. (1993) (The Opacity Project data base). weakly on temperature and density, and in many cases are independent of optical depth effects. We give ionic abundances using published values of effective recombination coefficients from lines that can be excited by recombination in Tables 14 and 15. Most calculations of effective recombination coefficients use term-averaged transition probabilities, which give rates for total intensities of multiplets. In most cases, some components of a multiplet are not detected and the total intensity of the multiplet must be estimated. If we assume that the relative populations of levels within a term are approximately proportional to their statistical weights, g j, the intensities of the lines of a multiplet will be proportional to the gf values, g j f ji g i f ij, which are proportional to the line strength, s ij. We have tested this hypothesis in Section 8.1 for multiplets with more than one observed component, with positive results in most cases. Thus the estimated total intensity of the multiplet can be obtained by multiplying the sum of intensities of observed lines by the multiplet correction factor m cf all i, j obs i, j s ij s ij, (1) where the upper sum runs over all the components of the multiplet, and the lower sum runs over the observed components of the multiplet. If the radial part of the dipole transition element is assumed constant for all the lines of the multiplet, s ij can be calculated for any ion in a given coupling scheme. Tabulations for multiplets in LS coupling are given by Allen (1973), and for mixed couplings by Gongora-T. & Escalante (1991). The importance of assuming the correct coupling scheme in the interpretation of astrophysical observations has been shown by Escalante & Gongora-T. (1990) and Liu et al. (1995b). Some multiplets show important deviations from LS coupling and intermediate coupling calculations are needed. In those cases we have taken s i,j from the compilation of Wiese et al. (1996) or from Liu et al. (1995b). Abundances obtained from the estimated total intensity of a multiplet are labelled as sum in Table 14, along with their m cf in LS coupling.