Faint Emission Lines and Temperature Fluctuations in M8

THE ASTROPHYSICAL JOURNAL SUPPLEMENT SERIES, 120 : 113È129, 1999 January ( 1999. The American Astronomical Society. All rights reserved. Printed in U.S.A.

FAINT EMISSION LINES AND TEMPERATURE FLUCTUATIONS IN M8 CE SAR ESTEBAN Instituto de Astrof• sica de Canarias, E-38200 La Laguna, Tenerife, Spain ; cel=ll.iac.es

MANUEL PEIMBERT AND SILVIA TORRES-PEIMBERT Instituto de Astronom• a, UNAM, Apdo. Postal 70-264, Mexico 04510 D.F., Mexico ; peimbert=astroscu.unam.mx, silvia=astroscu.unam.mx

AND JORGE GARCI A-ROJAS AND MO NICA RODRI GUEZ Instituto de Astrof• sica de Canarias, E-38200 La Laguna, Tenerife, Spain ; jgr=ll.iac.es, mrodri=ll.iac.es Received 1998 May 8 ; accepted 1998 August 10

ABSTRACT We present echelle spectroscopy in the 3500È10300 AŽ range of the Hourglass Nebula, which is embedded in the Galactic H II region M8. The data were obtained using the 2.1 m telescope at Observatorio Astronomico Nacional in San Pedro Martir, Baja California. We have measured the intensities of 274 emission lines, in particular 88 permitted lines of C`, N0, N`, O0, O`, Ne0, S0, S`, Si0, Si`, and Si``, some of them produced by recombination only and others mainly by Ñuorescence. We have determined electron temperatures and densities using di†erent line intensity ratios. We derive the He`, C``, O`, and O`` ionic abundances as well asÈfor the Ðrst time in a nebular objectÈthe total O abundance from recombination lines ; these nebular values are independent of the temperature structure of the nebula. We have also derived abundances from collisionally excited lines for a large number of ions and elements ; these abundances do depend on the temperature structure. Accurate t2 values have been derived by comparing the C``, O`, and O`` ionic abundances obtained making use of both collisionally excited lines and recombination lines. A comparison of the solar, Orion Nebula, and M8 chemical abundances is made. Subject headings : H II regions È ISM : abundances È ISM : individual (M8) È line : identiÐcation 1.

INTRODUCTION

We have obtained long-exposure CCD high spectral resolution echelle spectrograms in several spectral ranges of the Hourglass zone of M8 in order to obtain accurate measurements of recombination lines and other permitted lines of heavy element ions. The main aim of this work is to determine the O`` abundance from individual recombination linesÈavoiding the problem of line blendingÈand to compare it with the value derived using forbidden lines from the same spectrum. A similar comparison can be performed between the O` abundance obtained from O I lines produced by transitions between quintet terms (which are expected to be produced mainly by recombination and observable in the red part of the optical spectrum) and that obtained from the bright collisionally excited lines of this ion. Finally, we can carry out a similar comparison for the C`` abundance, but in this case, we have to compare our determinations using recombination lines with those obtained by PTD from UV C III] jj1906]1909 emission lines.

Messier 8 (M8), also known as the Lagoon Nebula, NGC 6523, S25, or W29, is one of the brightest and most studied Galactic H II regions. The ionized nebula forms a blister on the surface of a giant molecular cloud that is located behind the H II region (see, e.g., Lada et al. 1976 ; Lynds & OÏNeil 1982). The brightest zone of M8 is the Hourglass Nebula (hereafter HG). Most of the ionization of M8 is due to three O stars : 9 Sgr (HD 164794) of spectral type O5 V ((f)) (Auer & Mihalas 1972) ; HD 165052, an O6.5 V star (Walborn 1973) ; and Herschel 36 (H36, HD 164740) of spectral type O7.5 V (n) (Clayton & Cardelli 1988). There have been several previous works devoted to determining the chemical abundances in this nebula, all based on low-resolution spectroscopy (e.g., Peimbert & Costero 1969 ; Peimbert, TorresPeimbert, & Dufour 1993b, hereafter PTD). 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. Alternatively, recombination lines are almost independent of such variations and, in principle, they should be more precise indicators of the true chemical abundances. Very recently, Esteban et al. (1998, hereafter EPTE) have obtained O``/H` values from recombination O II line intensities for the Orion Nebula from echelle spectrophotometry, Ðnding that they are a factor of 1.5 larger than the O``/H` values obtained using forbidden lines. This discrepancy can be interpreted in terms of the e†ect of temperature Ñuctuations, t2, of the order of 0.024. A similar value of t2 has been estimated by Rubin et al. (1998) from Hubble Space T elescope (HST ) STIS spectra of the Orion Nebula covering the regions from the UV to the optical.

2.

OBSERVATIONS AND REDUCTIONS

The observations were carried out with the 2.1 m telescope in its f/7.5 conÐguration of the Observatorio Astronomico Nacional at San Pedro Martir, Baja California, Mexico in 1995 August and 1996 June. High-resolution CCD spectra were obtained using the REOSC Echelle Spectrograph ; its general characteristics are reported by Levine & Chakrabarty (1994). This instrument gives a resolution of 0.234 AŽ pixel~1 at Ha using the University College London (UCL) camera and a CCD-Tek chip of 1024 ] 1024 pixels with a 24 km2 pixels size. The spectral resolution is of 0.5 AŽ FWHM, and the accuracy in the 113

114

ESTEBAN ET AL.

wavelength determination of emission lines is of 0.1 AŽ . We obtained spectra in four overlapping wavelength intervals covering a very wide spectral range from 3500 to 10300 AŽ . Typically, three or four individual exposures were added to obtain the Ðnal spectra in each range. Slits covering 13A. 3 ] 2@@ in the blue exposures, 26A. 6 ] 2@@ in the red ones, and 39A. 9 ] 2@@ in both near-infrared (NIR) exposures were used to avoid overlapping between orders. The center of the single slit position taken on M8 was 12A south of the center of the Hourglass (zone HGS as deÐned by Sanchez & Peimbert 1991). The slit orientation was eastwest in all cases. Additional blue spectra were taken for the stars 9 Sgr and H36. A journal of observations is presented in Table 1. We used a Th-Ar lamp for wavelength calibration in all spectral ranges and a tungsten bulb for internal Ñat-Ðeld images. The absolute Ñux calibration of the blue and red spectra was achieved by taking echellograms of the standard stars HR 7596, HR 7950, and HR 8634 ; for both NIR spectra, the standard stars were HR 4963, HR 5501, and HR 7596. All the stars are from the list of Hamuy et al. (1992), which includes bright stars with Ñuxes sampled at 16 AŽ steps. An average curve for atmospheric extinction was used (Schuster 1982). The spectra were reduced using the IRAF1 echelle reduction package following the standard procedure of bias subtraction, aperture extraction, Ñat-Ðelding, wavelength calibration, and Ñux calibration. 3.

LINE INTENSITIES

Line intensities were measured integrating all the Ñux in the line between two given limits and over a local continuum estimated by eye. In the cases of line blending, a multiple Gaussian proÐle Ðt procedure was applied to obtain the line Ñux of each individual line. All these measurements were made with the SPLOT routine of the IRAF package. The brightest line of the red and the Ðrst NIR range spectrum, Ha, was close to saturation in the individual 1200 s spectra. Therefore, an additional 60 s exposure was taken with the same conÐguration for both spectral ranges, in order to obtain a useful Ha line Ñux as well as to check the linearity of the chip for the Ñuxes of the other bright lines measured in the 1200 s spectra. All the line intensities of a given spectrum have been normalized to a particular bright recombination line present in each wavelength interval. For the blue range, the reference line was Hb ; for the red range, He I j5876 ; for the Ðrst NIR range, He I j7065 ; and for the second NIR range, H I P13 j8665. To have a Ðnal homoge1 IRAF is distributed by NOAO, which is operated by AURA, under cooperative agreement with NSF.

neous set of line intensity ratios referred to the same line, the relative line intensities of each spectra were rescaled to Hb by means of the di†erent Ñux ratios of those recombination lines measured in consecutive spectral ranges. The di†erent spectral orders covered have overlapping regions at the edges. In these regions the optical sensitivity drops, and the line intensities are not accurate, and, therefore, they have not been considered. For the lines in common in two consecutive orders with good Ñux measurement, i.e., not detected at the edge of the order, the line intensity was the average of the values obtained in both orders. We estimate that the accuracy of the relative Ñux calibration achieved among the di†erent orders is about 2%È4% from the comparison of pairs of well measured Ñuxes of the same line and their underlying continua. In addition, each consecutive pair of spectral intervals covered has a common region in which it is possible to measure the same lines. The di†erences between the relative Ñuxes of a given lineÈwith respect to the same reference lineÈ obtained in the spectra of two consecutive spectral ranges do not present any systematic trend that could be related to anomalies in the relative Ñux calibration except for the second NIR range. In this case, the di†erences between the line intensities measured in both the Ðrst and second NIR spectra can be as large as 20%, being systematically lower in the second NIR spectrum. This fact could be due to Ñux calibration problems a†ecting only this last spectral range. The Ðnal intensity of a given line in the overlapping regions was the average of the values obtained in the di†erent consecutive spectra. For the lines present in both NIR spectra, the intensity measured in the Ðrst NIR range was the only one considered (owing to the more reliable Ñux calibration of this last range). The Ðnal 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 Ñux intensities (including all the possible sources of uncertainties in line intensity measurement and Ñux calibration) are estimated to be 2%È5% (0.01È0.02 dex) if the ratio F(j)/F(Hb) º 0.1, about 10% (0.04 dex) when 0.01 \ F(j)/F(Hb) \ 0.1, and about 20% (0.08 dex) when 0.001 \ F(j)/F(Hb) \ 0.01. For lines weaker than 0.001 ] F(Hb), the uncertainty could be ¹30% (0.12 dex) ; colons indicate uncertainties ¹40% (0.16 dex). For a given line, the observed wavelength (Table 2, fourth column) is determined by the center of the baseline chosen for the Ñux integration procedure or the centroid of the line when a Gaussian Ðt is used (in the case of line blending). For a line observed in di†erent spectra and orders, we have adopted the average wavelength of the di†erent measurements. The values given in the fourth column are observed wavelengths relative to the heliocentric reference frame.

TABLE 1 JOURNAL OF OBSERVATIONS

Object

Date

*j (AŽ )

Orders

Exposure Time (s)

M8 . . . . . . . .

1995 Aug 24 1995 Aug 23 1996 Jun 10 1996 Jun 13 1995 Aug 24 1995 Aug 24

3500È5950 4600È7075 6450È9100 8450È10300 3500È5950 3500È5950

38È63 30È46 25È34 22È26 38È63 38È63

600, 3600 60, 4800 60, 3600 2700 200 600

9 Sag . . . . . . H36 . . . . . . .

a Near-infrared.

Designation Blue Range Red Range First NIRa Range Second NIRa Range

TABLE 2

TABLE 2ÈContinued

OBSERVED AND REDDENING-CORRECTED LINE RATIOS AND IDENTIFICATIONS j 0 (AŽ ) 3512.51 . . . . . . 3530.49 . . . . . . 3554.46 . . . . . . 3587.26 . . . . . . 3613.64 . . . . . . 3634.28 . . . . . . 3662.26 . . . . . . 3663.41 . . . . . . 3664.68 . . . . . . 3666.10 . . . . . . 3667.68 . . . . . . 3669.47 . . . . . . 3671.48 . . . . . . 3673.76 . . . . . . 3676.37 . . . . . . 3679.36 . . . . . . 3682.81 . . . . . . 3686.83 . . . . . . 3691.56 . . . . . . 3697.15 . . . . . . 3703.86 . . . . . . 3705.04 . . . . . . 3711.97 . . . . . . 3721.83 . . . . . . 3721.94 . . . . . . 3726.03 . . . . . . 3728.82 . . . . . . 3734.37 . . . . . . 3750.15 . . . . . . 3770.63 . . . . . . 3797.90 . . . . . . 3805.77 . . . . . . 3819.61 . . . . . . 3833.57 . . . . . . 3835.39 . . . . . . ? ............. 3856.02 . . . . . . 3862.59 . . . . . . 3867.53 . . . . . . 3868.75 . . . . . . 3871.82 . . . . . . 3888.65 . . . . . . 3889.05 . . . . . . 3918.98 . . . . . . 3920.68 . . . . . . 3926.53 . . . . . . 3964.73 . . . . . . 3967.46 . . . . . . 3970.07 . . . . . . ? ............. 4009.22 . . . . . . 4026.19 . . . . . . 4068.60 . . . . . . 4069.62 . . . . . . 4069.89 . . . . . . 4072.15 . . . . . . 4075.80 . . . . . . 4076.35 . . . . . . 4101.74 . . . . . . 4120.81 . . . . . . 4143.76 . . . . . . 4153.30 . . . . . . 4168.97 . . . . . . 4169.22 . . . . . . 4243.98 . . . . . .

Ion

Multiplet

He I He I He I He I He I He I HI HI HI HI HI HI HI HI HI HI HI HI HI HI HI He I HI [S III] HI [O II] [O II] HI HI HI HI He I He I He I HI

(38) (36) (34) (31) (6) (28) H30 H29 H28 H27 H26 H25 H24 H23 H22 H21 H20 H19 H18 H17 H16 (25) H15 (2F) H14 (1F) (1F) H13 H12 H11 H10 (63) (22) (61) H9

Si II Si II He I [Ne III] He I He I HI C II C II He I He I [Ne III] HI

(1) (1) (20) (1F) (60) (2) H8 (4) (4) (58) (5) (1F) Hv

He I He I [S II] O II O II O II O II [S II] HI He I He I O II He I O II [Fe II]

(55) (18) (1F) (10) (10) (10) (10) (1F) Hd (16) (53) (19) (52) (19) (21F)

j (AŽ )

F(j)

I(j)

3512.37 3530.74 3554.71 3587.41 3613.84 3634.37 3662.30 3663.46 3664.76 3666.17 3667.73 3669.54 3671.55 3673.83 3676.46 3679.45 3682.95 3686.98 3691.59 3697.21 3703.87 3705.03 3711.98 3721.85

0.071 0.081 0.123 0.189 0.247 0.208 0.164 0.177 0.218 0.228 0.240 0.247 0.300 0.328 0.370 0.420 0.440 0.488 0.577 0.649 0.747 0.331 1.02 1.69

0.107 0.122 0.185 0.282 0.368 0.310 0.242 0.262 0.322 0.336 0.354 0.365 0.444 0.484 0.546 0.619 0.649 0.718 0.849 0.954 1.10 0.486 1.50 2.47

3726.06 3728.81 3734.37 3750.10 3770.59 3797.85 3805.56 3819.56 3833.52 3835.33 3854.69 3855.96 3862.55 3867.32 3868.57 3871.58 3888.91

66.9 44.5 1.50 1.84 2.30 3.31 0.062 : 0.632 0.048 4.33 0.034 0.125 0.105 0.033 1.87 0.072 8.32

97.8 64.9 2.19 2.68 3.33 4.75 0.089 : 0.904 0.068 6.16 0.048 0.178 0.148 0.047 2.64 0.101 11.7

3918.83 3920.57 3926.44 3964.65 3967.31 3970.00 3998.64 4009.14 4026.11 4068.51 4069.67

0.065 0.127 0.086 0.602 0.643 10.8 0.027 0.137 1.29 0.770 0.040

0.090 0.177 0.119 0.827 0.882 14.8 0.037 0.185 1.74 1.02 0.052

4072.05 4076.24

0.031 0.307

0.041 0.407

4101.64 4120.73 4143.69 4153.27 4168.79

19.1 0.152 0.208 0.029 0.015

25.1 0.198 0.269 0.038 0.019

4243.67

0.024

0.030

j 0 (AŽ )

Ion

Multiplet

4267.26 . . . . . . . 4287.46 . . . . . . . 4317.14 . . . . . . . 4340.47 . . . . . . . 4345.56a . . . . . . 4349.43 . . . . . . . 4359.34 . . . . . . . 4363.21 . . . . . . . 4366.89 . . . . . . . 4368.25 . . . . . . . 4387.93 . . . . . . . 4416.27 . . . . . . . 4437.55 . . . . . . . 4471.48 . . . . . . . 4590.97 . . . . . . . 4596.17 . . . . . . . 4601.48 . . . . . . . 4607.06 . . . . . . . 4607.16 . . . . . . . 4613.87 . . . . . . . 4621.35 . . . . . . . 4630.54 . . . . . . . 4638.85 . . . . . . . 4641.81 . . . . . . . 4643.09 . . . . . . . 4649.14 . . . . . . . 4650.84 . . . . . . . 4658.10 . . . . . . . 4661.64 . . . . . . . 4667.00 . . . . . . . 4676.23 . . . . . . . 4701.62 . . . . . . . 4711.34 . . . . . . . 4713.14 . . . . . . . 4716.36 . . . . . . . 4733.90 . . . . . . . 4752.73b . . . . . . 4754.70 . . . . . . . 4769.40 . . . . . . . 4777.70 . . . . . . . 4779.71 . . . . . . . 4788.13 . . . . . . . 4803.29 . . . . . . . 4814.55 . . . . . . . 4815.51 . . . . . . . 4861.33 . . . . . . . 4881.00 . . . . . . . 4921.93 . . . . . . . 4924.50 . . . . . . . 4930.50 . . . . . . . 4931.00 . . . . . . . 4958.91 . . . . . . . 4985.90 . . . . . . . 4987.20 . . . . . . . 4991.94a . . . . . . 4994.36 . . . . . . . 5001.13b . . . . . . 5006.84 . . . . . . . 5011.30 . . . . . . . 5015.68 . . . . . . . 5032.43 . . . . . . . 5035.50 . . . . . . . 5041.03 . . . . . . . 5045.10 . . . . . . . 5047.74 . . . . . . . 5055.98 . . . . . . . 5056.31 . . . . . . .

C II [Fe II] O II HI O II O II [Fe II] [O III] O II OI He I [Fe II] He I He I O II O II N II [Fe III] N II N II N II N II O II O II N II O II O II [Fe III] O II [Fe III] O II [Fe III] [Ar IV] He I [Fe II] [Fe III] Ne I [Fe III] [Fe III] [Fe III] N II N II N II [Fe II] S II HI [Fe III] He I [Fe III] [Fe III] [O III] [O III] [Fe III] [Fe III] S II N II N II [O III] [Fe III] He I S II [Fe II] Si II N II He I Si II Si II

(6) (7F) (2) Hc (2) (2) (7F) (2F) (2) (5) (51) (6F) (50) (14) (15) (15) (5) (3F) (5) (5) (5) (5) (1) (1) (5) (1) (1) (3F) (1) (3F) (1) (3F) (1F) (12) (5F) (3F) (21) (3F) (3F) (3F) (20) (20) (20) (20F) (9) Hb (2F) (48) (2F) (1F) (1F) (1F) (2F) (2F) (7) (24.64) (19) (1F) (1F) (4) (7) (4F) (5) (4) (47) (5) (5)

j (AŽ )

F(j)

I(j)

4267.01 4287.24 4317.08 4340.37 4345.58 4349.33 4359.36 4363.13 4366.75 4368.16 4387.86 4416.21 4437.52 4471.43 4590.81 4596.11 4601.26 4607.08

0.184 0.040 0.027 38.7 0.071 0.027 0.030 0.312 0.021 0.018 0.401 0.010 0.056 3.41 0.020 0.009 0.028 0.034

0.228 0.049 0.033 46.6 0.086 0.032 0.036 0.373 0.025 0.022 0.475 0.012 0.065 3.91 0.022 0.010 0.030 0.037

4614.02 4621.31 4630.48 4638.75 4641.76 4643.02 4649.06 4650.68 4658.08 4661.48 4666.99 4676.11 4701.48 4711.31 4713.07 4716.18 4733.85 4752.81 4754.72 4769.41 4777.64 4779.65 4788.10 4803.35 4814.48 4815.53 4861.30 4880.95 4921.82 4924.50 4930.58 4931.16 4958.90 4985.78 4987.16 4991.84 4994.23 5001.11 5006.74 5011.34 5015.66 5032.19 5035.76 5040.95 5045.07 5047.68 5055.96

0.018 0.019 0.043 0.030 0.040 0.016 0.050 0.033 0.467 0.033 0.018 0.015 0.136 0.012 0.445 0.017 0.049 0.016 0.088 0.049 0.022 0.010 0.011 0.017 0.027 0.029 100 0.198 1.08 0.033 0.013 0.021 42.8 0.032 0.043 0.016 0.021 0.045 130 0.069 2.21 0.043 0.022 0.138 0.016 0.165 0.234

0.020 0.021 0.047 0.033 0.043 0.018 0.053 0.035 0.501 0.035 0.020 0.016 0.144 0.013 0.468 0.017 0.051 0.017 0.091 0.050 0.023 0.010 0.011 0.017 0.027 0.030 100 0.197 1.07 0.032 0.013 0.021 41.4 0.031 0.042 0.015 0.020 0.043 124 0.066 2.10 0.041 0.021 0.130 0.016 0.155 0.220

TABLE 2ÈContinued j 0 (AŽ ) 5084.80 . . . . . . . 5103.30 . . . . . . . ? .............. 5158.81 . . . . . . . 5191.82 . . . . . . . 5197.90 . . . . . . . 5200.26 . . . . . . . 5261.61 . . . . . . . 5270.40 . . . . . . . 5273.38 . . . . . . . 5299.00 . . . . . . . ? .............. 5412.00 . . . . . . . 5428.60 . . . . . . . 5432.77 . . . . . . . 5453.81 . . . . . . . 5495.70 . . . . . . . 5495.82 . . . . . . . 5512.77 . . . . . . . 5517.71 . . . . . . . 5537.88 . . . . . . . 5551.95 . . . . . . . 5555.03 . . . . . . . 5666.64 . . . . . . . 5676.02 . . . . . . . 5679.56 . . . . . . . 5686.21 . . . . . . . 5710.76 . . . . . . . 5739.73 . . . . . . . 5754.64 . . . . . . . ? .............. 5875.67 . . . . . . . 5891.65 . . . . . . . ? .............. 5927.82 . . . . . . . 5931.79c . . . . . . 5941.65 . . . . . . . 5952.39c . . . . . . 5957.56 . . . . . . . 5958.58 . . . . . . . 5978.93 . . . . . . . 6000.20 . . . . . . . 6046.44c . . . . . . 6142.53b . . . . . . ? .............. 6286.35c . . . . . . 6300.30c . . . . . . 6312.10 . . . . . . . 6347.09 . . . . . . . 6363.78c . . . . . . 6371.36 . . . . . . . 6401.50 . . . . . . . 6402.25b . . . . . . 6482.07b . . . . . . 6527.40 . . . . . . . 6533.80 . . . . . . . 6548.03 . . . . . . . 6562.82 . . . . . . . 6578.05 . . . . . . . 6583.41 . . . . . . . ? .............. 6678.15 . . . . . . . 6716.47 . . . . . . . 6730.85 . . . . . . . ? .............. ? .............. 7002.23 . . . . . . .

Ion

Multiplet

[Fe III] S II

(1F) (7)

[Fe II] [Ar III] [N I] [N I] [Fe II] [Fe III] [Fe II] OI

(19F) (3F) (1F) (1F) (19F) (1F) (18F) (26)

[Fe III] S II S II S II N II [Fe II] OI [Cl III] [Cl III] N II OI N II N II N II N II N II Si III [N II]

(1F) (6) (6) (6) (29) (17F) (25) (1F) (1F) (63) (24) (3) (3) (3) (3) (3) (4) (3F)

He I C II

(11) (5)

N II N II N II N II Si II OI Si II [Ni III] OI Si I

(28) (28) (28) (28) (4) (23) (4) (2F) (22) (21.13)

S II [O I] [S III] Si II [O I] Si II [Ni III] Ne I N II [N II] [Ni III] [N II] HI C II [N II]

(19) (1F) (3F) (2) (1F) (2) (2F) (1) (8) (1F) (2F) (1F) Ha (2) (1F)

He I [S II] [S II]

(46) (2F) (2F)

OI

(21)

j (AŽ )

TABLE 2ÈContinued

F(j)

I(j)

5084.73 5103.48 5146.69 5158.72 5191.62 5197.78 5200.14 5261.61 5270.49 5273.34 5299.01 5342.24 5412.09 5428.61 5432.82 5453.82 5495.45

0.014 0.008 0.010 0.040 0.040 0.138 0.083 0.021 0.285 0.015 0.017 0.017 0.033 0.008 0.014 0.019 0.017

0.013 0.008 0.009 0.037 0.036 0.124 0.075 0.019 0.251 0.013 0.015 0.015 0.028 0.007 0.012 0.016 0.014

5512.81 5517.61 5537.79 5551.84 5554.90 5666.63 5675.97 5679.54 5686.22 5710.70 5739.61 5754.51 5865.15 5875.64 5891.56 5893.84 5927.73 5932.07 5941.53 5952.62 5957.45 5958.46 5978.97 6000.27 6046.37 6142.64 6151.26 6286.72 6300.09 6312.10 6347.15 6363.51 6371.37 6401.92

0.012 0.497 0.529 0.011 0.009 0.038 0.024 0.044 0.020 0.016 0.011 0.940 0.021 15.6 0.009 0.018 0.012 0.045 0.017 0.031 0.059 0.016 0.135 0.017 0.042 0.059 0.019 0.051 1.33 1.91 0.272 0.442 0.161 0.041

0.010 0.408 0.432 0.009 0.008 0.030 0.019 0.035 0.016 0.012 0.009 0.727 0.016 11.7 0.006 0.013 0.009 0.033 0.012 0.023 0.043 0.012 0.099 0.012 0.030 0.042 0.013 0.034 0.903 1.29 0.182 0.295 0.107 0.027

6482.10 6527.27 6533.67 6548.09 6562.80 6578.06 6583.43 6666.62 6678.23 6716.47 6730.85 6863.45 6989.38 7002.16

0.009 0.024 0.016 34.5 431 0.354 123 0.017 4.84 10.1 14.1 0.140 0.023 0.070

0.006 0.016 0.010 22.1 274 0.224 77.7 0.011 3.00 6.18 8.61 0.083 0.013 0.040

j 0 (AŽ ) 7065.28 . . . . . . 7135.78 . . . . . . 7155.14 . . . . . . ? .............. 7231.12 . . . . . . 7236.19 . . . . . . 7254.38 . . . . . . 7281.35 . . . . . . 7298.05 . . . . . . 7319.65 . . . . . . 7330.16 . . . . . . 7377.83 . . . . . . 7411.61 . . . . . . 7423.63 . . . . . . 7499.82 . . . . . . 7751.12 . . . . . . 7771.96 . . . . . . 7774.18c . . . . . . 7775.40 . . . . . . 7816.16 . . . . . . 7889.9 . . . . . . . . 8187.95 . . . . . . 8216.28 . . . . . . 8223.07 . . . . . . 8242.34 . . . . . . 8245.64 . . . . . . 8247.73 . . . . . . 8249.97 . . . . . . 8252.40 . . . . . . 8255.02 . . . . . . 8257.86 . . . . . . 8260.94 . . . . . . 8264.29 . . . . . . 8267.94 . . . . . . 8271.93 . . . . . . 8276.31 . . . . . . 8281.12 . . . . . . 8286.43 . . . . . . 8292.31 . . . . . . 8298.84c . . . . . . 8306.12 . . . . . . 8314.26 . . . . . . 8323.43 . . . . . . 8333.78 . . . . . . 8359.01 . . . . . . 8361.77 . . . . . . 8374.48 . . . . . . 8392.40 . . . . . . 8413.32 . . . . . . 8446.48 . . . . . . 8467.26 . . . . . . 8502.49 . . . . . . 8528.99 . . . . . . 8545.38 . . . . . . 8578.70 . . . . . . 8582.54 . . . . . . 8598.39 . . . . . . 8616.96 . . . . . . 8665.02 . . . . . . ? .............. 8845.38 . . . . . . 8862.79 . . . . . . 8914.74 . . . . . . 8996.99 . . . . . . 9014. 91 . . . . . . 9068.9 . . . . . . . . 9210.28 . . . . . .

Ion

Multiplet

He I [Ar III] [Fe II]

(10) (1F) (14F)

C II C II OI He I He I [O II] [O II] [Ni II] [Ni II] NI He I [Ar III] OI OI OI He I [Ni III] NI NI NI NI HI HI HI HI HI HI HI HI HI HI HI HI HI HI HI HI HI HI HI HI He I HI HI HI OI HI HI He I HI [Cl II] He I HI [Fe II] HI

(3) (3) (20) (45)

He I HI He I He I HI [S III] He I

(2F) (2F) (2F) (2F) (3) (1F) (1) (1) (1) (69) (1F) (2) (2) (2) (2) P42 P41 P40 P39 P38 P37 P36 P35 P34 P33 P32 P31 P30 P29 P28 P27 P26 P25 P24 P22 (68) P21 P20 P19 (4) P17 P16 P15 (1F) P14 (13F) P13

P11

P10 (1F) (83)

j (AŽ )

F(j)

I(j)

7065.26 7136.74 7155.24 7160.62 7231.17 7236.40 7254.39 7281.08 7296.96 7320.15 7330.38 7377.88 7411.70 7423.47 7499.66 7750.92 7771.73 7773.68 7775.22 7815.72 7890.18 8188.19 8216.42 8223.18 8242.11 8245.69 8247.76 8250.02 8252.34 8255.10 8257.95 8260.96 8264.36 8267.95 8271.84 8276.20 8281.28 8286.39 8292.21 8298.85 8305.95 8314.08 8323.21 8333.57 8358.70 8361.38 8374.18 8392.10 8413.27 8446.58 8467.46 8502.65 8529.23 8545.34 8578.31 8582.06 8598.01 8616.49 8664.35 8833.17 8845.32 8862.34 8914.01 8995.59 9013.40 9067.91 9210.61

5.87 18.1 0.045 0.034 0.182 0.289 0.115 0.955 0.088 6.48 5.33 0.051 0.015 0.013 0.056 5.29 0.041 0.085 0.014 0.114 0.214 0.055 0.045 0.051 0.059 0.050 0.060 0.064 0.076 0.106 0.117 0.126 0.154 0.134 0.134 0.169 0.280 0.251 0.233 0.479 0.226 0.264 0.307 0.360 0.438 0.188 0.481 0.544 0.615 0.843 0.544 0.973 0.040 1.13 0.360 0.080 1.43 0.094 1.93 0.390 0.094 2.87 0.044 0.138 4.12 58.1 0.145

3.33 10.1 0.025 0.019 0.099 0.158 0.061 0.511 0.047 3.46 2.85 0.027 0.008 0.007 0.029 2.56 0.019 0.042 0.007 0.054 0.100 0.024 0.020 0.022 0.026 0.022 0.055 0.028 0.033 0.046 0.050 0.054 0.067 0.058 0.058 0.073 0.121 0.108 0.100 0.206 0.097 0.114 0.132 0.155 0.185 0.079 0.203 0.230 0.260 0.349 0.225 0.403 0.017 0.459 0.150 0.030 0.581 0.038 0.769 0.150 0.037 1.10 0.017 0.052 1.55 21.8 0.052

ECHELLE SPECTROPHOTOMETRY OF M8 TABLE 2ÈContinued j 0 (AŽ ) 9229.02 . . . . . . . 9463.57 . . . . . . . 9545.97 . . . . . . . 9603.50 . . . . . . . 9672.56b . . . . . . 9850.24 . . . . . . . ? ..............

Ion

Multiplet

HI He I HI He I SI [C I]

P9 (67) P8 (71) (17) (1F)

j (AŽ )

F(j)

I(j)

9229.24 9462.93 9545.47 9602.03 9672.75 9849.86 10021.05

4.81 0.195 5.11 0.057 0.058 0.862 0.063

1.74 0.068 1.78 0.019 0.019 0.282 0.020

a Possibly blended with an unidentiÐed line. b Dubious identiÐcation. c Possibly a†ected by sky emission.

The identiÐcation and adopted laboratory wavelengths (Table 2, Ðrst column) of the lines collected in Table 2 were obtained following previous identiÐcations in the Orion Nebula by EPTE and Osterbrock, Tran, & Veilleux (1992) ; we also used the compilations of atomic data by Moore (1945, 1993) and Wiese, Smith, & Glennon (1966) and the papers of Hyung, Aller, & Feibelman (1993, 1994a, 1994b) on echelle spectrophotometry of the planetary nebulae NGC 6567, IC 4497, and NGC 6572, respectively. Sky emission lines were not included in the list. Thirteen emission lines could not be identiÐed in any of the available references. Some of these unidentiÐed lines have been reported previously in Orion and in some planetary nebulae. The line at 5146.69 AŽ has been observed in the Orion Nebula (5146.88 AŽ , EPTE) and IC 4997 (5146 AŽ ). The intensity of this line varies almost an order of magnitude among these three objects ; therefore, if it is a true nebular line, it is probably excited by a collisional or Ñuorescence mechanism. The line at 5342.24 AŽ has been also observed in the Orion Nebula (5342.63 AŽ , Osterbrock et al. 1992 ; 5342.40 AŽ , EPTE), IC 4497 (5342.56 AŽ ), and NGC 6567 (5342.38 AŽ ). The line at 6151.26 AŽ has been reported in the Orion Nebula (6151.34 AŽ , EPTE), IC 4997 (6151.29 AŽ ), and NGC 6567 (6151.45 AŽ ). Finally, the line observed at 7160.62 AŽ has been reported in IC 4997 (7160.62 AŽ ) and NGC 6567 (7160.60 AŽ ). The intensity of each of these last three lines is very similar in all the objects, despite the large di†erences in excitation conditions of the nebulae, a fact that suggests the true nebular nature of these lines and recombination as their most likely production mechanism. It is known that the main ionization source of HG is the star H36 (see, e.g., Woodward et al. 1986) and that it shows a considerably higher extinction than the other zones of M8. For H36 the A /E(B[V ) ratio, R, has been determined as 4.6 (Hecht et al.V1982) and as 5.3 (Cardelli, Clayton, & Mathis 1989). Following Sanchez & Peimbert (1991) and PTD we have adopted a reddening function with R \ 5.0. The use of this reddening law instead of the standard one TABLE 3 PHYSICAL CONDITIONS

Lines

n e (cm~3)

Lines

[O II] . . . . . . . [Cl III] . . . . . . [S II] . . . . . . . . [Fe III] . . . . . . [Fe II] . . . . . .

1600`1300 ~700 3000 : 1900`;2000 ~900 2000È10000 103È106

[O III] . . . . . . [O II] . . . . . . . [N II] . . . . . . . [S II] . . . . . . . . [Ar III] . . . . . .

T e (K) 8050 ^ 700 9600`2000 ~1600 8250`1050 ~900 8800`2600 ~1800 8100`1300 ~1200

117

(Whitford 1958) produces a di†erent C(Hb) but very similar line intensity ratios in the optical region. The reddening coefficient, C(Hb), was determined by Ðtting iteratively the observed Balmer decrement to the theoretical one computed by Storey & Hummer (1995) for the physical conditions of HGS. The Ðnal C(Hb) for each slit position was the average of the values obtained from the Ha/Hb, Hc/Hb, and Hd/Hb ratios. The Ðnal adopted value of C(Hb) is 0.85 ^ 0.05. This value is slightly lower than the C(Hb) \ 1.00 ^ 0.10 obtained by Sanchez & Peimbert (1991) and PTD for HGS. Table 2 shows the reddeningcorrected line intensity ratios, I(j)/I(Hb), for each line and slit position. The reddening-corrected Hb surface brightness from the blue exposure spectra is 1.72 ] 10~12 ergs cm~2 s~1 arcsec~2. 4.

PHYSICAL CONDITIONS

The large number of emission lines identiÐed and measured in the spectra allows the derivation of physical conditions using line ratios of di†erent ions. The values of n and e T given in Table 3 have been obtained using the Ðve-level e program for the analysis of emission-line nebulae of Shaw & Dufour (1995), except in the cases of the n derived from [Fe II] and [Fe III] lines. For [Fe II] we ehave used the diagnostic diagrams of line intensity ratios versus n published by Bautista & Pradhan (1998 ; mainly for thosee involving NIR [Fe II] lines) and additional diagrams obtained from the computations by Bautista, Peng, & Pradhan (1996). For [Fe III] we have used the diagnostic diagrams by Keenan et al. (1993), which are broadly consistent with the results obtained from the most recent diagnostic diagrams by Bautista & Pradhan (1998). The values of n obtained for the di†erent ions available e are very similar, except for [Fe II]. An in each slit position average value of 1750 cm~3 has been adopted as representative for HGS. Values of 2500È3700 cm~3 were obtained previously by PTD for the same zone of the nebula. On the other hand, the electron density derived for Fe` has a very high dispersion from 103 to 106 cm~3 depending on the line ratio used (excluding those lines probably a†ected by large Ñuorescence e†ects ; see ° 6.1). The same behavior was found by EPTE in the case of the Orion Nebula. It is interesting to note that the presence of [Fe II]Èemitting high-density regions in partially ionized zones in the Orion Nebula has been proposed by Bautista, Pradhan, & Osterbrock (1994) based on the comparison between observed and predicted [Fe II] line intensity ratios at di†erent physical conditions, whereas Lucy (1995) has suggested that photon pumping by UV continuum can explain the observed line ratios. Electron temperatures from forbidden lines have been derived from [O II], [O III], [N II], [S II], and [Ar III] line ratios. The values obtained from di†erent ions are similar within the uncertainties. However, T ([O II]) is signiÐcantly higher than T ([N II]). Similar values of T ([O II]) and T ([N II]) were obtained by Sanchez & Peimbert (1991) and PTD. Unfortunately, the electron temperatures derived from the Balmer and Paschen continua have very high uncertainties to be useful owing to the combination of the high spectral resolution and the relative weakness of the nebular continuum. 5. H I

AND He I RECOMBINATION SPECTRA

5.1. H I L ines The H I Balmer and Paschen spectra can be detected up

118

ESTEBAN ET AL.

to H30 and P42, respectively, in our data. In Table 4, we present the comparison between the observed relative intensities of the H I Balmer and Paschen lines with respect to the predicted ones using the machine-readable line-ratio tables of Storey & Hummer (1995). The theoretical line ratios have been evaluated for T \ 8250 K and n \ 1750 cm~3. The e e adopted value of T corresponds to the mean of the electron e temperatures assumed for high- and low-ionization ions weighted by their relative uncertainties (see ° 6). In any case, the H I line ratios are almost independent of the adopted temperatures and densities. The comparison between observed and predicted line ratios can not go further than H25 and P25 because this is the limit of the calculations of Storey & Hummer (1995). In the optical range, H8 and H14 have not been included in Table 4 because they are blended with He I j3888.65 and [S III] j3721.83, respectively. In the NIR ranges, P23, P18, P17, and P12 have not been included in Table 5, because P23 and P17 are severely a†ected by strong sky emission lines, while P18 and P12 are at wavelengths not covered between consecutive spectral NIR orders. Columns (2) and (4) of Table 4 gives the ratio of the observed to the predicted intensities. The di†erences between theoretical and observed line ratios are typically of the order of 10%È15% for the Balmer and Paschen lines except for those lines measured only in the second NIR spectral range (from P11 to P8), where we could have a Ñux calibration problem as was noted in ° 3. 5.2. He I L ines There are 42 He I emission lines identiÐed in our spectra. These lines arise mainly from recombination, although they may have contributions from collisional and self-absorption e†ects. In Table 5, 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 \ 8250 K and e TABLE 4 OBSERVED OVER PREDICTED H I BALMER AND PASCHEN EMISSION-LINE RATIOS Line (1)

I /I obs pred (2)

Line (3)

I /I obs pred (4)

H25 . . . . . . H24 . . . . . . H23 . . . . . . H22 . . . . . . H21 . . . . . . H20 . . . . . . H19 . . . . . . H18 . . . . . . H17 . . . . . . H16 . . . . . . H15 . . . . . . H13 . . . . . . H12 . . . . . . H11 . . . . . . H10 . . . . . . H9 . . . . . . . Hv . . . . . . . Hd . . . . . . . Hc . . . . . . . Hb . . . . . . . Ha . . . . . . .

0.93 1.02 1.00 1.01 1.01 0.93 0.90 0.91 0.88 0.84 0.95 0.92 0.88 0.84 0.90 0.85 0.94 0.98 1.00 1.00 0.95

P25 . . . . . . P24 . . . . . . P22 . . . . . . P21 . . . . . . P20 . . . . . . P19 . . . . . . P16 . . . . . . P15 . . . . . . P14 . . . . . . P13 . . . . . . P11 . . . . . . P10 . . . . . . P9 . . . . . . . P8 . . . . . . . ... ... ... ... ... ... ...

1.00 1.06 1.00 0.97 0.96 0.94 0.88 0.83 0.85 0.91 0.55 0.69 0.67 0.48a ... ... ... ... ... ... ...

a At edge of a spectral order.

Vol. 120 TABLE 5 OBSERVED OVER PREDICTED He I EMISSION-LINE RATIOS j (AŽ ) (1)

Transition (2)

I /I obs pred (3)

5015.68 . . . . . . 3964.73 . . . . . . 3613.64 . . . . . . 7281.35 . . . . . . 5047.74 . . . . . . 4437.55 . . . . . . 4168.97 . . . . . . 6678.15 . . . . . . 4921.93 . . . . . . 4387.93 . . . . . . 4143.76 . . . . . . 4009.22 . . . . . . 3926.53 . . . . . . 3871.82 . . . . . . 3833.57 . . . . . . 3805.76 . . . . . . 9603.50 . . . . . . 8914.74 . . . . . . 7065.28 . . . . . . 4713.14 . . . . . . 4120.81 . . . . . . 3867.48 . . . . . . 5875.67 . . . . . . 4471.48 . . . . . . 4026.19 . . . . . . 3819.61 . . . . . . 3705.04 . . . . . . 3634.28 . . . . . . 3587.26 . . . . . . 3554.46 . . . . . . 3530.49 . . . . . . 3512.51 . . . . . . 9463.57 . . . . . . 8361.70 . . . . . . 7816.16 . . . . . . 7499.82 . . . . . . 7298.05 . . . . . . 9210.28 . . . . . . 8996.99 . . . . . . 8845.38 . . . . . . 8582.54 . . . . . . 8528.99 . . . . . .

2 1SÈ3 1P 2 1SÈ4 1P 2 1SÈ5 1P 2 1PÈ3 1S 2 1PÈ4 1S 2 1PÈ5 1S 2 1PÈ6 1S 2 1PÈ3 1D 2 1PÈ4 1D 2 1PÈ5 1D 2 1PÈ6 1D 2 1PÈ7 1D 2 1PÈ8 1D 2 1PÈ9 1D 2 1PÈ10 1D 2 1PÈ11 1D 3 1SÈ6 1P 3 1SÈ7 1P 2 3PÈ3 3S 2 3PÈ4 3S 2 3PÈ5 3S 2 3PÈ6 3S 2 3PÈ3 3D 2 3PÈ4 3D 2 3PÈ5 3D 2 3PÈ6 3D 2 3PÈ7 3D 2 3PÈ8 3D 2 3PÈ9 3D 2 3PÈ10 3D 2 3PÈ11 3D 2 3PÈ12 3D 3 3SÈ5 3P 3 3SÈ6 3P 3 3SÈ7 3P 3 3SÈ8 3P 3 3SÈ9 3P 3 3DÈ9 3F 3 3DÈ10 3F 3 3DÈ11 3F 3 3DÈ14 3F 3 3DÈ15 3F

0.99 0.98 0.90 0.98 1.07 1.07 0.60a 0.96 1.02 1.00 1.03 1.17 1.14 1.41 1.31 2.29b 1.33 0.79 2.01 1.29 1.43 0.68 1.07 1.00 0.96 0.91 0.80 0.78 0.94 0.92 0.81 0.93 0.76c 1.31 3.74 1.03 2.33d 0.63 0.89 0.85 1.40 1.01

a Blended with O II j4169.22 AŽ . b Probably blended. c At edge of a spectral order. d Dubious identiÐcation.

n \ 1750 cm~3. The ratios between observed and predictede line intensity ratios are presented in column (3). As in the case of the H I Balmer spectrum, the observed ratios compare well with the predicted ones, taking into account that uncertainties of the order of 10%È20% are expected owing to the relative faintness of most of the He I lines as well as some contribution of collisional e†ects from the metastable 2 3S level (see Kingdon & Ferland 1996). An apparent trend is detected in the 2 3PÈn 3S and 3 3SÈn 3P lines (except for He I j3867, which is very weak and for He I j9463, which is just at the edge of an spectral order in the second NIR range) ; the observed intensities of these lines are systematically brighter than the predicted ones by more than 20%. This discrepancy can be explained by self-

No. 1, 1999

ECHELLE SPECTROPHOTOMETRY OF M8

absorption e†ects from the metastable 2 3S level, which produce large optical depths in the lines for the 2 3SÈn 3P and 3 3SÈn 3P transitions. The He I lines most a†ected by self-absorption are jj3889, 7065, and 10830. We have direct measurements of only He I j7065, which, in fact, shows a very large departure from the theoretical ratio. Unfortunately, He I j3889 is blended with H8, and He I j10830 is outside the spectral range covered by our data. On the other hand, the ratios presented in Table 5 indicate the absence of signiÐcant line-transfer e†ects in the helium singlet spectrum (see, e.g., Robbins & Bernat 1973). 6.

IONIC ABUNDANCES FROM FORBIDDEN LINES

Ionic abundances of N`, O`, O``, Ne``, S`, S``, Cl`, Cl``, Ar``, and Ar3` have been obtained from collisionally excited lines, using the Ðve-level atom program of Shaw & Dufour (1995) and the atomic parameters referenced in it. Additionally, we have determined the ionic abundances of Fe`, Fe``, Ni`, and Ni`` following the methods and data discussed in ° 6.1. For the low ionization potential ions N`, O`, S`, Cl`, Fe`, Fe``, Ni`, and Ni``, we have adopted a mean of T ([O II]), T ([N II]), and T ([S II]), weighted by their relative uncertainties, which is 8740 K. On the other hand, we have adopted T ([O III]), which is almost coincident with T ([Ar III]), for the high ionization potential ions O``, Ne``, S``, Cl``, Ar``, and Ar3`. The density assumed is 1750 cm~3 in all cases. Ionic abundances are listed in Table 6 and correspond to the mean value of the abundances derived from all the individual lines of each ion observed (which do not show large dispersion with respect to the mean). The values obtained are very consistent with those derived by PTD for the ions in common (di†erences not larger than 0.2 dex) and remarkably similar to those obtained by Pipher et al. (1984) for S`` and Ar`` from Ðne-structure lines in the IR. In contrast, the ionic abundances obtained for O`` and S`` are TABLE 6 IONIC ABUNDANCES FROM FORBIDDEN LINESa t2 Xm N` . . . . . . . . . O` . . . . . . . . . O`` . . . . . . . Ne`` . . . . . . S` . . . . . . . . . S`` . . . . . . . Cl` . . . . . . . . Cl`` . . . . . . Ar`` . . . . . . Ar3` . . . . . . . Fe` . . . . . . . . Fe` . . . . . . . . Fe`` . . . . . . Ni` . . . . . . . . Ni`` . . . . . .

0.000

0.032

7.30 8.18 8.02 6.76 5.77 6.86 4.28 5.06 6.24 3.65 5.42b 4.66c 5.73 4.20 5.15

7.42 8.31 8.27 7.03 5.89 7.04 4.39 5.29 6.44 3.85 5.55b 4.79c 5.98 4.31 5.26

a In units of 12 ] log (Xm/ H`). b Mean value of all [Fe II] lines except those of multiplets (4F), (5F), and (7F). c Considering only [Fe II] jj7155 and 8617.

119

very di†erent from those obtained by Afflerbach, Churchwell, & Werner (1997), which were also from Ðne-structure IR lines. 6.1. Iron and Nickel Forbidden L ines Twelve [Fe II] lines belonging to multiplets (4F), (5F), (6F), (7F), (13F), (14F), (18F), (19F), (20F), and (21F) have been identiÐed and measured in HGS. The presence of [Fe II] lines in the spectrum of M8 was reported by Rodr• guez (1996) who measured six lines of this ion. We have interpolated the data of Bautista & Pradhan (1996) to derive the Fe` emission-line ratios and abundances using the standard values of electron temperature and density assumed for low ionization potential ions (T \ 8740 K, n \ 1750 e e cm~3). Additionally, and only to compare line ratios, we have considered a density of 106 cm~3 as suggested by the results of using [Fe II] diagnostic diagrams (° 4). The comparison between the observed and predicted emission-line ratios as well as the Fe`/H` values derived are presented in Table 7. We have selected [Fe II] j5262 as a reference line because it is a single line at the center of the optical spectral range that belongs to a multiplet presenting two measurable lines in our spectra. Columns (5) and (6) of Table 7 show the ratio between observed and predicted line ratios for the lower and higher values of the electron density, respectively. The large deviation that some lines present with respect to the theoretical predictions is remarkable. The di†erence is especially large for the single lines of multiplets (4F) and (5F) observed, which are not expected to present a large contribution by Ñuorescence excitation because the quartet levels can not be pumped by dipole allowed transitions from the 6D ground state. Moreover, a hypothetical Ñuorescence excitation would also a†ect other lines of both multiplets that would be measurable. On the other hand, the [Fe II] j5035 line has also been observed in the Orion Nebula (EPTE) with I(j5035)/ I(j5262) between 0.66 and 0.34, which is not too far to the value of 1.11 observed in HGS. In any case, it is likely that these two lines have been misidentiÐed. Moreover, our data for the two measured lines of multiplet (7F) indicate either wrong theoretical intensities or the presence of large Ñuorescence excitation for this multiplet. In fact, it is expected that multiplet (7F) su†ers the largest Ñuorescence e†ects because the atomic levels involved in this transition are sextets and, therefore, pumping from the 6D ground state can be produced by allowed dipole transitions. However, Bautista et al. (1996) propose that the bulk of the [Fe II] emission comes from a high-density region (a partially ionized zone with 106È107 cm~3), much higher than the critical density for Ñuorescence excitation of [Fe II], and therefore, they do not expect strong Ñuorescence e†ects for this ion. EPTE also observed enhancements for the lines of this multiplet in the two slit positions observed in the Orion Nebula, but especially for position 1 (the zone closer to the ionizing star) where the enhancement factor is about 9.2, quite similar to the value of 9.9 observed in HGS. The presence of Ñuorescence as evidenced by the line intensities of multiplet (7F) is an indication that these lines are produced in a relatively low electron density medium and, therefore, in a place with conditions similar to those where the rest of the lines form. In most cases, the correspondence between observed and theoretical line ratios of [Fe II] is not very good and shows large dispersion (this large dispersion with respect to the

120

ESTEBAN ET AL.

Vol. 120

TABLE 7 OBSERVED OVER PREDICTED [Fe II] EMISSION-LINE RATIOS AND Fe` ABUNDANCE I(j)/I(j5262) j (AŽ ) (1)

MULTIPLET (2)

TRANSITION (3)

Observed (4)

Obs./Pred.a (5)

Obs./Pred.b (6)

N(Fe`)/N(H`)a (]10~7) (7)

5035 . . . . . . 4716 . . . . . . 4416 . . . . . . 4287 . . . . . . 4359 . . . . . . 8617 . . . . . . 7155 . . . . . . 5273 . . . . . . 5159 . . . . . . 5262 . . . . . . 4815 . . . . . . 4244 . . . . . .

(4F) (5F) (6F) (7F) (7F) (13F) (14F) (18F) (19F) (19F) (20F) (21F)

a 6D Èb 4P 1@2 5@2 a 6D Èa 4H 7@2 9@2 a 6D Èb 4F 9@2 9@2 a 6D Èa 6S 9@2 5@2 a 6D Èa 6S 7@2 5@2 a 4F Èa 4P 9@2 5@2 a 4F Èa 2G 9@2 9@2 a 4F Èb 4P 9@2 5@2 a 4F Èa 4H 9@2 13@2 a 4F Èa 4H 7@2 11@2 a 4F Èb 4F 9@2 9@2 a 4F Èa 4G 9@2 11@2

1.11 0.89 0.63 2.58 1.89 2.00 1.32 0.68 1.95 1.00 1.42 1.58

1881 2282 0.28 9.92 9.95 0.07 0.08 0.20 0.68 1.00 0.81 0.27

3700 1079 0.83 24.6 24.9 0.68 0.25 0.39 1.34 1.00 2.37 1.23

11600 13400 1.75 61.5 61.8 0.42 0.50 1.24 4.23 6.20 5.00 1.67

a n \ 1750 cm~3. e b n \ 106 cm~3. e

theoretical values was also noticed by Rodr• guez 1996). The use of high electron densities does not improve the global consistency between observations and theoretical calculations. In this sense, the similarity of the discrepancies between observed over predicted [Fe II] line ratios found in HGS and those obtained for the same lines in the Orion Nebula by EPTE is striking. In fact, for the lines in common of multiplets (18F), (19F), (20F), and (21F), the di†erences between the ratios obtained for HGS and the Orion Nebula do not exceed 20%. On the other hand, the scatter in the Fe` abundance obtained from the di†erent [Fe II] lines (deÐned as p /x, where p and x are the standard deviation x and the mean value xof the abundance, respectively) obtained assuming low and high densities becomes somewhat lower when the high-density case is considered. We derive p /x \ 0.84 and 0.66 for low and high densities, x respectively. A similar result was obtained for the Orion Nebula by Bautista et al. (1994) and forms the basis of their suggestion of a high-density [Fe II]Èemitting region, but the high intensities we measure for the lines of multiplet (7F) imply that Ñuorescence e†ects can be important and a meaningful comparison of the scatter in the Fe` abundances would require their consideration. We obtain a mean Fe`/H` ratio of 2.63 ] 10~7 from each individual line excluding those of multiplets (4F), (5F), and (7F) (multiplets having large Ñuorescence contribution or being possible misidentiÐcations). On the other hand, if the bulk of the [Fe II] is in fact produced in the low-density zone, it is very probable that all optical lines are a†ected by Ñuorescence (see Lucy 1995 ; Baldwin et al. 1996 ; Bautista & Pradhan 1998). In this case, the most reliable value of Fe`/H` is that obtained from [Fe II] 7155 and 8617 AŽ lines, which give a mean Fe`/H` \ 4.6 ] 10~8. Both estimates of the Fe` abundance are included in Table 6. Seventeen [Fe III] lines of multiplets (1F), (2F), and (3F) have been measured in HGS. The presence of some [Fe III] lines in the spectrum of M8 was previously reported by PTD and Rodr• guez (1996). We have used the Fe III level populations tables computed by Keenan et al. (1992) to obtain the theoretical line ratios and the Fe`` abundance. The comparison between the observed and predicted emission line ratios as well as the Fe``/H` values are presented

in Table 8. The consistency between the observed and the predicted line ratios is remarkable, except in the cases where line blending is present. The good agreement between theory and observations implies little scatter in the Fe``/H` ratios derived from the di†erent [Fe III] emission lines. The adopted average Fe`` abundance of 5.35 ] 10~7 has been obtained from those lines not a†ected by line blending and is included in Table 6. Our value of the Fe``/H` is consistent with previous determinations by PTD and Rodr• guez (1996). Two [Ni II] emission lines are observed in our spectra, jj7379 and 7411, the brightest lines of multiplet (2F), corresponding to the 3d9a 2DÈ3d84s a 2F transition. There is no previous reference about any observation of [Ni II] lines in M8. We have derived the Ni` abundance from these lines, 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 used to derive the Fe`/H` ratio. With all these ingredients, we derive Ni`/H` \ 1.6 ] 10~8 and 5.7 ] 10~8, for [Ni II] jj7379 and 7411, respectively. This discrepancy arises from the abnormal jj7411/7379 ratio, which is a factor 3.65 higher than predicted. A similar value was observed by Osterbrock et al. (1992) and by Henry & Fesen (1988) in the Orion Nebula. Lucy (1995) explained quantitatively this behavior in the observed jj7411/7379 ratio as an e†ect of Ñuorescence by stellar UV continuum. On the other hand, Bautista et al. (1996) suggest that this high ratio is the result of both pure collisional excitation in a high-density zone with 106 cm~3 and continuum Ñuorescence from the lowdensity zone as for Fe`. The value of Ni`/H` \ 1.6 ] 10~8 obtained from [Ni II] j7379 must be considered as an upper limit to the Ni` abundance since Lucy (1995) found that this transition can also be considerably a†ected by continuum pumping ; for the Orion Nebula he estimated an intensity enhancement of a factor 5. Four [Ni III] lines of multiplets (1F) and (2F) have been measured in our spectra. It is possible that j6401.5 is blended with a Ne I line and its intensity might be overestimated. Unfortunately, there are no calculations of the level populations of this ion because no published collision strengths are available for [Ni III]. Osterbrock et al. (1992)

No. 1, 1999

ECHELLE SPECTROPHOTOMETRY OF M8

121

TABLE 8 OBSERVED OVER PREDICTED [Fe III] EMISSION-LINE RATIOS AND FE`` ABUNDANCE I(j)/I(j4658) j (AŽ ) 5270 . . . . . . 5011 . . . . . . 4931 . . . . . . 5412 . . . . . . 5085 . . . . . . 4986 . . . . . . 4925 . . . . . . 4987 . . . . . . 4881 . . . . . . 4658 . . . . . . 4702 . . . . . . 4734 . . . . . . 4607 . . . . . . 4667 . . . . . . 4755 . . . . . . 4769 . . . . . . 4778 . . . . . .

TRANSITION

MULTIPLET

Observed

Obs./Pred.

N(Fe``)/N(H`) (]10~7)

5D È3 P 3 2 5D È3 P 2 1 5D È3 P 1 0 5D È3 P 1 2 5D È3 P 0 1 5D È3 H 4 6 5D È3 H 4 5 5D È3 H 3 4 5D È3 H 4 4 5D È3 F 4 4 5D È3 F 3 3 5D È3 F 2 2 5D È3 F 4 3 5D È3 F 3 2 5D È3 F 3 4 5D È3 F 2 3 5D È3 F 1 2

(1F) (1F) (1F) (1F) (1F) (2F) (2F) (2F) (2F) (3F) (3F) (3F) (3F) (3F) (3F) (3F) (3F)

0.501 0.132 0.030 0.056 0.026 0.062 0.064 0.084 0.393 1.000 0.287 0.102 0.074 0.040 0.182 0.100 0.046

1.54 2.28 6.00 1.87 2.60 0.65 10.7 1.83 1.51 1.00 0.96 1.11 1.72 1.67 1.01 1.05 1.02

5.76 8.45 18.4 6.94 9.84 2.39 38.6 6.83 5.62 3.73 3.57 4.14 6.39 6.15 3.76 3.91 3.84

estimate these atomic parameters from available calculations for [Fe VII], with analogous electron ground conÐguration. We have calculated the Ni``/H` ratio following the procedure outlined by Osterbrock et al. (1992) that uses the collision strengths estimated by these authors and the transition probabilities obtained by Garstang (1958). We then Ðnd the abundance ratios Ni``/ H` \ 1.1 ] 10~7, 0.4 ] 10~7, and 1.4 ] 10~7 from jj6000, 6534, and 7890, respectively. We adopt the value of 1.4 ] 10~7 obtained from j7890, the brightest by far of the three [Ni III] lines. 7. He`

ABUNDANCES

The comparison between observed and predicted He I emission lines shown in Table 5 indicates that lines coming from 2 3PÈn 3S and 3 3SÈn 3P transitions could su†er substantial self-absorption e†ects and, therefore, that they are not suitable for deriving an accurate He`/H` ratio. On the other hand, collisional e†ects may a†ect all the measured lines, but unfortunately, we have collision-to-recombination factors (C/R) for only eight of the lines available. In Table 9, we present the He` abundance for these eight linesÈ uncorrected and corrected by the collisional contributionÈ as well as their corresponding C/R factors. The uncorrected TABLE 9 He`/H` ABUNDANCE RATIO j (AŽ )

Uncorrected

C/R

Corrected

4026 . . . . . . . . . . 4388 . . . . . . . . . . 4471 . . . . . . . . . . 4922 . . . . . . . . . . 5876 . . . . . . . . . . 6678 . . . . . . . . . . 7065a . . . . . . . . . 7282 . . . . . . . . . . Mean . . . . . .

0.0752 0.0777 0.0780 0.0800 0.0837 0.0752 0.1571 0.0767 0.0781 ^ 0.0030

0.0062 0.0041 0.0096 0.0056 0.0250 0.0120 0.2700 0.1007

0.0747 0.0774 0.0773 0.0795 0.0816 0.0743 0.1237 0.0761 0.0773 ^ 0.0026

a Not included in the derivation of the mean He`/H`.

NOTES

Blend

Blend

Blend

He`/H` ratios have been obtained using the predicted line emissivities calculated by Smits (1996) and assuming the same physical conditions as in ° 5.2. The C/R factors have been obtained from the most recent calculations by Kingdon & Ferland (1995). The average value He`/ H` \ 0.077 has been derived excluding He I j7065 (the line that su†ers the largest self-absorption e†ects) and is given in Table 9. This value is larger than the He`/H` ratio of 0.065 obtained by PTD for HG. This apparently discrepant result is not unexpected, taking into account the results of Sanchez & Peimbert (1991). These authors obtained a relatively high dispersion in the He`/H` fraction obtained for di†erent parts of HG (from 0.067 to 0.099). 8.

PERMITTED HEAVY ELEMENT LINES

We have measured 88 permitted lines of heavy element ions such as O I, O II, C II, Ne I, S I, S II, N I, N II, Si I, Si II, and Si III, most of them detected for the Ðrst time in M8. 8.1. Excitation Mechanisms Most of the permitted lines of heavy elements observed in M8 have been observed also in the Orion Nebula by EPTE. These authors carried out a detailed discussion of the excitation mechanisms of those lines that could be also applied to the lines in common with the present work. In what follows, we will focus mainly in the excitation mechanism of those lines not observed or not discussed in the Orion Nebula by EPTE. We have observed seven C II lines in our spectra corresponding to multiplets 2, 3, 4, 5, and 6. The most likely dominant excitation mechanism for the lines of multiplets 2, 4, 5, and 6 in the Orion Nebula was discussed by EPTE. In the case of M8, the same conclusions are applicable. Resonance Ñuorescence by starlight can be important for transitions with upper 2S and 2D levels, namely multiplets 3 and 4. On the other hand, multiplets 2 and 5 have np 2P0 upper terms that can be populated from transitions excited by both resonance Ñuorescence by starlightÈfrom excited 2S and 2D levelsÈand recombination from terms having large L values ; hence, for these lines, the dominant mechanism should be resonance Ñuorescence by starlight, with a contri-

122

ESTEBAN ET AL.

bution from recombination. Grandi (1976) showed that recombination dominates by an order of magnitude the excitation of j4267.26. This line comes from a transition involving terms with large L values, terms that cannot be reached by permitted resonance transitions from the ground term and, therefore, are excited mainly by recombination. The measurement of several lines of the same multiplet of a given ion permits the comparison of their relative intensities with those expected in an appropriate angular momentum coupling ; this comparison is made under the assumption that the populations of the Ðne-structure levels within a term are proportional to their statistical weights. The results of this comparison are shown in Table 10. In this table we include those multiplets with more that one emission line measured. Columns (1)È(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. Column (5) gives the observed intensity of each line relative to the strength of the expected brightest line of the multiplet. Column (6) includes the ratio of the observed and the predicted relative intensities. In Table 10, we show that the relative intensities between the two observed lines of multiplets 3 and 4 are in very good agreement with the predictions of L S coupling. We observed Ðve N I lines in our spectra of M8. Grandi (1975a) has shown that continuum Ñuorescence can feed the 3d 4P and 4s 4P terms and therefore a†ect the observed strength of the lines. The four lines of multiplet 2 are the brightest of the multiplet, and their ratios are not consistent with the predicted values for L S coupling as can be seen in Table 10. Grandi (1976) has shown that resonance Ñuorescence by the recombination line He I j508.64 is the dominant mechanism to excite the 4s 3P0 term of N II in the Orion Nebula, and hence it could be 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 be also applied to multiplet 20, which has the same upper term as multiplet 28, and perhaps to multiplet 24, which can be fed also by this mechanism from the ground level. The single line of multiplet 19 observed, N II j5001.13, has an upper 3d 3F0 term that cannot be reached by permitted resonance lines, and therefore, recombination is the only mechanism that can produce this line. Alternatively, N II j5001.13 is not the brightest line of the multiplet, and we expect to see other lines in the case of recombination. The only singlet line of N II observed, j6482.07, can be excited only by recombination, but if this identiÐcation is true, we should see other brighter singlet lines, namely j4447.03, for example. Therefore, the lines of multiplets 8 and 19 could have been misidentiÐed. In Table 10, we compare the relative intensities of the N II lines of the same multiplet with those predicted in L S coupling. All the multiplets show ratios between observed and predicted strengths far from unity, which indicates that the population of these levels is not proportional to their statistical weights. We have measured 12 O I emission lines in HGS. In most cases, the multiple components for neutral species are not resolved owing to the small Ðne-structure splitting of their terms. Thus the intensities of most of the O I lines in Table 2 correspond to the total intensity of the multiplet. The excitation mechanism of the O I lines is mainly Ñuorescence due

to stellar continuum radiation, as shown quantitatively by Grandi (1975a, 1975b). The upper levels of the lines of multiplets 20 to 26 correspond to 3p 3PÈns 3S0 or 3p 3PÈnd 3D0 transitions ; these transitions can be excited directly by absorption of stellar photons longward of 912 AŽ from the ground term followed by permitted downward transitions. The population of the upper level of multiplets 4 and 5È corresponding to 3s 3S0Ènp 3P transitionsÈcan be fed from transitions from s and d levels with higher principal quantum numbers. Therefore, starlight excitation may contribute signiÐcantly to the observed strength of these lines. On the other hand, we detect and resolve the three lines of multiplet 1 of O I at jj7772, 7774, and 7775, which are transitions between quintet terms that cannot be produced by Ñuorescence from the ground state. This multiplet is expected to be the brightest of the quintets and to be produced by recombination. The O I j7774.18 line su†ers from severe contamination by sky emission, probably producing its high observed intensity. As can be seen in Table 10, there is good agreement between the observations and the predictions of L S coupling for the ratio of the two other O I lines. We have measured 17 O II lines of multiplets 1, 2, 10, 15, and 19. The presence of many of these emission lines cannot be explained by resonance Ñuorescence as was demonstrated by Grandi (1976). Moreover, this author estimates that starlight excitation of the O II j430 line contributes only 20% as much as recombination to the line intensities of multiplet 19, with the e†ect considerably smaller for the rest of the lines. Therefore, we will assume that recombination is the dominant excitation mechanism of the observed O II lines. As can be seen in Table 10, there is no good agreement between observations and predictions of L S coupling for many of the O II lines. The ratios di†er by factors larger than 2 for several of them, contrary to the situation in the Orion Nebula (EPTE). The identiÐcation of the two Ne I lines (of multiplets 1 and 21) in our spectra is dubious because we did not detect other lines, such as those of multiplets 11 or 13, that should be of intensity comparable to or greater than that of the lines of multiplet 21. On the other hand, the Ne I j6402.25 line is blended or confused with the [Ni III] j6401.50 line. We detected one line identiÐed as S I belonging to multiplet 17. It corresponds to a 4s@ 3DoÈ4p@ 3D transition. The excitation mechanism of this multiplet is uncertain ; continuum Ñuorescence by absorption of j1278 UV photons from the 3p4 3P ground state could excite the higher level of the transition. 2Pure recombination is also possible, but, in this case, we expect to see other quintet lines in the optical region that are not detected. Therefore, the correct identiÐcation of this line is dubious. We detected eight lines of multiplets 6, 7, 9, and 19 of S II. The excitation mechanism of these lines is also very uncertain. Recombination should produce other multiplets with higher intensity such as multiplets 44 (4p 4DoÈ4d 4F) or 45 (4p4 4DoÈ4d 4D). Furthermore, continuum Ñuorescence from the ground state is unlikely because the resonance lines that could lead to the excitation of the observed lines have wavelengths shorter than the Lyman limit. In any case, the observation of several of the brightest lines of multiplets 6 and 7 indicates that their identiÐcation is probably correct. In Table 10 we can see that there is only moderate agreement between the observations and the predictions of L S coupling for the relative intensities of the S II lines of multiplets 6 and 7.

TABLE 10 OBSERVED OVER PREDICTED INTENSITY RATIOS OF PERMITTED LINES OF HEAVY ELEMENTS

Ion (1)

Multiplet (2)

C II . . . . . .

3 4

N I ......

2

N II . . . . . .

3

5

20 28

O I ...... O II . . . . . .

1

1

2

10

15 19 S II . . . . . .

6

7

Si II . . . . . .

1 2 4 5

j (AŽ ) (3)

JÈJ@ (4)

I obs (5)

I /I obs pred (6)

7236.19 7231.12 3920.68 3918.98 8216.28 8242.34 8223.07 8187.95 5679.56 5666.64 5676.02 5710.76 5686.21 4630.54 4613.87 4643.09 4621.35 4601.48 4607.16 4788.13 4779.71 5941.65 5931.79 5927.79 5952.39 7771.96 7774.18 7775.40 4649.14 4641.81 4638.85 4676.85 4661.64 4650.84 4349.43 4366.89 4345.56 4317.14 4072.15 4069.62 4069.89 4590.97 4596.17 4153.30 4169.22 5453.81 5432.77 5428.64 5032.41 4991.94 5103.30 3856.02 3862.59 6347.09 6371.36 5978.93 5957.56 5055.98 5056.31 5041.03

3/2È5/2 1/2È3/2 3/2È1/2 1/2È1/2 5/2È5/2 5/2È3/2 3/2È1/2 1/2È3/2 2È3 1È2 0È1 2È2 1È1 2È2 1È1 2È1 1È0 1È2 0È1 2È2 1È1 2È3 1È2 0È1 2È2 2È3 2È2 2È1 5/2È7/2 3/2È5/2 1/2È3/2 5/2È5/2 3/2È3/2 1/2È1/2 5/2È5/2 5/2È3/2 3/2È1/2 1/2È3/2 5/2È7/2 3/2È5/2 1/2È3/2 5/2È7/2 3/2È5/2 3/2È5/2 5/2È5/2 5/2È7/2 3/2È5/2 1/2È3/2 5/2È5/2 3/2È3/2 5/2È3/2 5/2È3/2 3/2È1/2 1/2È3/2 1/2È1/2 3/2È1/2 1/2È1/2 3/2È5/2 3/2È3/2 1/2È3/2

1 0.63 1 0.51 1 1.30 1.10 1.20 1 0.86 0.54 0.34 0.46 1 0.43 0.38 0.45 0.64 0.79 1 0.91 1 2.75 0.75 1.92 1 2.62 0.37 1 0.81 0.62 0.30 0.66 0.66 1 0.78 2.69 1.03 1 1.27

1 1.12 1 0.96 1 3.02 2.62 3.16 1 1.61 2.28 1.93 2.60 1 3.11 1.13 2.10 2.01 3.27 1 1.40 1 5.12 3.15 10.77 1 3.74 0.86 1 1.54 3.00 1.31 2.46 3.17 1 1.70 7.17 2.60 1 1.20

1 0.45 1 0.86 1 0.75 0.44 1 0.37 0.20 1 0.83 1 0.59 1 0.43 1

1 0.65 1 2.98/2.52 1 1.27 1.76 1 2.47 0.39 1 1.47 1 1.17 1 0.86 1

0.59

1.17

Notes (7)

Blend

Blend Blend Blend

Blend

Blend

124

ESTEBAN ET AL.

We detected one line identiÐed as Si I belonging to multiplet 21.13. It corresponds to a transition 3p 3Do 3È5f @[5/2]. Its identiÐcation is highly dubious. We have measured nine Si II lines belonging to multiplets 1, 2, 4, and 5 ; all of them correspond to doublets. Grandi (1976) found that starlight excitation of the 5s 2S and 4d 2D terms dominates over recombination for the strength of these lines in the Orion Nebula. Table 10 shows the good agreement between the observed intensity ratios and the L S coupling predictions for all the observed multiplets of Si II. Finally, the Si III j5793.73 line of multiplet 4 is the only permitted line of this ion observed in M8. Grandi (1976) found that the strength of this line is produced by a combination of starlight excitation and recombination in the Orion Nebula. 8.2. Ionic Abundances Permitted lines produced mainly by recombination can give accurate determinations of ionic abundances because under nebular conditions, their relative intensities depend weakly on temperature and density. Let I(j) be the intensity of a recombination line of an element i times ionized at wavelength j ; then the abundance of the ionization state i ] 1 of element X is given by N(Xi`1) j(AŽ ) a (Hb) I(j) eff \ , (1) N(H`) 4861 a (j) I(Hb) eff where a represents the e†ective recombination coefficient. Mosteffcalculations of a use term-averaged transition eff for total intensities of multiprobabilities, which give rates plets. However, in many cases, our echelle observations can resolve several lines of each multiplet. If we assume that the relative populations of levels within a term are approximately proportional to their statistical weights, g , the j to intensities of the lines of a multiplet will be proportional the gf-values, g f \ g f , which are proportional to the ji have i ijtested this hypothesis in ° 8.1 for line strength, s j. We ij multiplets with more than one observed component and have found large departures for most cases. This is in contrast with the results obtained by EPTE for the Orion Nebula, where there is considerable agreement between observations and predictions for L S coupling. This di†erence could be due to the lower signal-to-noise ratio in our M8 data as compared to that for the Orion Nebula. The comparison between Figure 3 of EPTE and Figure 1 of Garc• a-Rojas et al. (1998) illustrates the di†erent signal-tonoise ratio of the spectra of both objects in the zone of multiplet 1 of O II. We have determined the ionic abundances from individual lines (see Tables 11, 12, and 13) multiplying the a of the multiplet by a factor that takes eff

TABLE 12 O`/H` RATIO FROM O I LINES

Multiplet 1 .........

C``/H` (]10~5) MULTIPLET 2 ........... 3 ...........

6 ...........

6578.05 7236.19 7231.12 Mean m cf Sum 4267.26

0.224 0.158 0.099 1.00 0.257 0.228

Case A

Case B

74 22 24 23

16 22 24 23

21 19

21 ...

I(j)/I(Hb) [I(Hb) \ 100]

O`/H` (]10~5) Case A

7771.96 7775.40 Mean m cf Sum

0.019 0.007

19 17 18

1.49 0.034

19

s ij ; (2) ; s all i,j ij where the sum runs over all the components of the multiplet. Additionally, we have determined the mean of the ionic abundances obtained from all the lines of a given multiplet. Abundances obtained in this way are labeled as ““ Mean ÏÏ in Tables 11, 12, and 13. An alternative approach is to estimate the total intensity of the multiplet by multiplying the sum of intensities of observed lines by the multiplet correction factor, which introduces the contribution of allÈobserved and TABLE 13 O``/H` RATIO FROM O II LINESa O``/H`(]10~5) MULTIPLET 1 ...........

2 ...........

10 . . . . . . . . . .

C``/H` RATIO FROM C II LINES I(j)/I(Hb) [I(Hb) \ 100]

j (AŽ )

into account the relative intensity of each line within the multiplet

TABLE 11

j (AŽ )

Vol. 120

15 . . . . . . . . . .

19 . . . . . . . . . .

j (AŽ )

I(j)/I(Hb) [I(Hb) \ 100]

4638.85 4641.81 4649.14 4650.84 4661.64 4676.23 Mean m cf Sum 4345.56 4349.43 4366.89 4317.14 Mean m cf Sum 4069.62 4069.89 4072.15 Mean m cf Sum 4590.97 4596.17 Mean m cf Sum 4153.30

0.033 0.043 0.053 0.035 0.035 0.016 1.02 0.219 0.086b 0.032 0.025 0.033 1.54 0.139 0.052 0.041 1.98 0.184 0.022 0.010 1.03 0.033 0.044

Case A

Case B

34 16 11 37 28 17 24

32 15 11 35 27 17 23

18 159 24 40 65 43

18 114 17 29 46 31

36 20

26 ...

17 19

... ...

19 132 86 109

... ... ... ...

113 1300/1300

... 48/56c

a All the values are calculated assuming L S coupling except those after a slash, which correspond to intermediate coupling. b Probably blended. c The same values apply for Case C.

No. 1, 1999

ECHELLE SPECTROPHOTOMETRY OF M8

unobservedÈlines ; s m \ all i,j ij , (3) cf ; s obs i,j ij where the upper sum runs over all the components of the multiplet, and the lower sum runs over the observed components of the multiplet. Abundances obtained from the estimated total intensity of a multiplet are labeled as ““ Sum ÏÏ in Tables 11, 12, and 13 along with their m in L S cf coupling. Abundances deduced from this last method are preferred over average abundances when the m factor is cf close to 1 because observational errors and the e†ects of departures from L S coupling are minimized. We have used the C II a values obtained by Pequignot, eff Petitjean, & Boisson (1991) to derive the C`` abundances. The physical conditions assumed are T \ 8050 K and n \ e e 1750 cm ~3. The C``/H` ratios obtained are shown in Table 11. The brightest line of this ion in the optical range, C II j4267.26, is case insensitive and can be used to derive a proper C``/H` ratio. On the other hand, C II j6578 appears to be quite sensitive to optical depth e†ects, and, therefore, inappropriate for abundance determinations. In any case, and in agreement with the results of EPTE for the Orion Nebula, for j6578.03, case B gives a C`` abundance consistent with multiplet 6. Therefore, the contribution of resonance Ñuorescence by starlight expected by Grandi (1976) for this line seems to be negligible in comparison with the contribution due to recombination. A similar situation has been found for the lines of multiplet 3, which is also case insensitive. For this multiplet, the C`` abundance is consistent with the value obtained from C II j4267.26, which indicates that Ñuorescence e†ects on the line strengths of this multiplet are not important for the physical conditions of HGS. Pequignot et al. (1991) give a Ïs for di†erent multiplets of eff we have considered multiO I. To derive the O`/H` ratio plet 1, which is the only one probably produced by pure recombination. For this multiplet no case B needs to be considered because the multiplicity of the levels involved in the transition di†er from the multiplicity of the ground state of O I. The physical conditions assumed are T \ 8740 and n \ 1750 cm ~3. The O`/H` abundances aree presented in e Table 12. The ““ Sum ÏÏ value of the O`/H` abundance is 19 ] 10~4. For comparison, abundances derived making use of O I j8446Èwhich is strongly a†ected by ÑuorescenceÈgive values between one and two orders of magnitude larger. Storey (1994) has computed the O II a Ïs for cases A, B, eff derive the O`` and C. We have used these coefficients to abundances listed in Table 13. We have assumed that the relative line intensities within a multiplet follow those predicted in L S coupling except in the case of multiplet 19, where we have also considered intermediate coupling following Liu et al. (1995). The physical conditions assumed are the same as for the C II calculations. The abundances determined from multiplets 1 and 2 are almost independent of the case assumed and virtually independent for multiplet 10 (the case is not included in Table 13 because a is the eff B is same for both cases). Liu et al. (1995) Ðnd that case more appropriate for quartets and case A for doublets, in agreement with the results of Peimbert, Storey, & TorresPeimbert (1993a). Inspection of the di†erent values of the abundance obtained for multiplet 19 indicates that there is

125

no signiÐcant di†erence between the result obtained assuming L S coupling or intermediate coupling calculations (values separated by a slash in Table 13). We used the dielectronic recombination rate of Nussbaumer & Storey (1984) to obtain an O`` abundance from multiplet 15. The O`` abundances obtained from the di†erent lines and multiplets observed show signiÐcant dispersion. The O``/H` values for multiplets 1, 2, and 10 are almost case insensitive and consistent. In contrast, abundance derived from multiplets 15 and 19 are signiÐcantly larger as it was also found for the Orion Nebula by EPTE. A possible problem with the derivation of O`` abundances is the e†ect of absorption by the underlying dustscattered stellar continuum. This possibility was discussed by Peimbert et al. (1993a) in the case of the Orion Nebula. However, EPTE do not Ðnd any trace of underlying absorption of the nebular O II lines by the stellar continuum of h1 Ori C, the brightest star of the Trapezium. In the case of M8, we have not detected the presence of O II absorption lines in the spectra of 9 Sgr and H36. 9.

TEMPERATURE FLUCTUATIONS IN M8

Ionic abundances from forbidden lines often di†er signiÐcantly from those derived from recombination lines (see, e.g., Peimbert et al. 1993a ; Liu et al. 1995 ; EPTE), probably because both abundance determinations have a di†erent dependence on the temperature structure of the nebula. In Table 14 we show the comparison between the abundances of O`, O``, and C`` determined from collisionally excited lines and recombination lines. As was found also for the Orion Nebula by EPTE, the ionic abundances determined from collisionally excited lines are always somewhat lower than those obtained from recombination lines. The abundances derived from standard forbidden line analysis for O` and O`` are taken from Table 6. The O`/H` ratio from recombination lines has been taken from the ““ Sum ÏÏ value of Table 12. In the case of O``/H`, the ratio for recombination is given by the average of the values obtained for multiplets 1, 2, and 10. We select the ““ Sum ÏÏ value for case B as representative of multiplets 1 and 2. Multiplet 10 is case insensitive, and the abundances derived assuming cases A and B are the same. Therefore we have also selected the ““ Sum ÏÏ value for this last multiplet. The C``/H` ratio obtained from collisionally excited lines has been taken from PTD. These authors have measured the UV C III] jj1906]1909 emission lines from IUE data for exactly the same zone of M8 that we have observed (HGS). The representative C``/H` from recombination lines is that obtained using the C II j4267 line of multiplet 6. Comparison between ionic abundances obtained from collisionally excited lines and recombination lines can TABLE 14 t2 PARAMETER

12 ] log (Xm/H`) O` . . . . . . . . . . . . . . . . O`` . . . . . . . . . . . . . . C`` . . . . . . . . . . . . . . Mean . . . . . . . . . . .

Forbidden Lines (t2 \ 0.00)

Recombination Lines

8.18 8.02 7.95

8.28 8.32 8.28

t2(R/C) 0.022 0.040 0.032 0.032

^ ^ ^ ^

0.030 0.023 0.015 0.019

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ESTEBAN ET AL.

provide an estimate of the spatial temperature Ñuctuations in this zone of M8. The presence of these Ñuctuations, t2, in gaseous nebulae was Ðrst proposed by Peimbert (1967, 1971). We have derived the t2 value that produces the agreement between the ionic abundances obtained from both recombination lines and forbidden lines, which are shown in Table 14. We have adopted a value of t2 \ 0.032 ^ 0.019 for HGS as the mean value (weighted by their relative uncertainties) of the Ñuctuations obtained from O`, O`` and C``. This value of the spatial temperature Ñuctuation is somewhat lower than (but within the uncertainties of) the value of 0.046 ^ 0.01 obtained by PTD from the comparison of the C`` abundance obtained from collisionally excited and recombination lines. Values of t2 ranging from 0.00 to 0.09 have been found for planetary nebulae, with a typical value around 0.04 (Peimbert 1971 ; Dinerstein, Lester, & Werner 1985 ; Liu & Danziger 1993 ; PTD ; Liu et al. 1995 ; Kingsburgh, Lopez, & Peimbert 1996). The values derived by EPTE for the Orion Nebula are smaller than the typical values derived for planetary nebulae and the value we have obtained for M8 ; alternatively, they are about 0.02 higher than those derived from photoionized models (i.e., see Baldwin et al. 1991). Accurate t2 values should be determined for other nebulae because they are of great importance to determining the mechanism that produces such t2 values and in the quest to obtain accurate abundances for gaseous nebulae. Very recently, Liu (1998) has obtained a di†erence of a factor of D15 between collisional and recombination C``/H` ionic abundance in the planetary nebula NGC 4361. Liu interprets that this large di†erence cannot be accounted for with temperature Ñuctuations and could be related to uncertainties in the e†ective recombination coefficients. In the case of our results for M8 and previous ones for the Orion Nebula (EPTE), the agreement of the t2 values obtained for di†erent ions (C``, O`, and O`` in the case of M8 and C`` and O`` in the case of the Orion Nebula) suggests strongly that these di†erences are not related to uncertainties in the e†ective recombination coefficients and more likely related to external physical reasons.

Vol. 120 10.

TOTAL ABUNDANCES

To derive the total gaseous abundances we need to adopt a t2 value and to correct for the unseen ionization stages by using ionization correction factors, i . The abundances cf derived from recombination lines are independent of t2 and consequently are more reliable than those derived from collisionally excited lines. The total abundances adopted for each element are presented in Table 15. 10.1. Forbidden L ines The abundances derived from collisionally excited lines have been computed assuming t2 values of 0.00 and 0.032 that can be used to interpolate or extrapolate for other t2 values. The absence of He II emission lines in the spectra (He II j4686 /Hb ¹ 0.00005) and the similarity between the ionization potentials of He` and O`` implies the absence of measurable O3` in HGS. Therefore, to obtain the total oxygen abundance we can simply assume that N(O) N(O` ] O``) \ . N(H) N(H`)

To derive the total nitrogen abundance we have used the following expression : N(N) N(N`) N(O` ] O``) N(N`) \ i (N`) ] \ ] , cf N(H) N(H`) N(O`) N(H`) (5) which is a good approximation for objects of low ionization such as HGS (see, e.g., Stasinska 1990 ; Mathis & Rosa 1991). The only measurable collisionally excited lines of Ne in the optical region are those of Ne``. The ionization potential of this ion is very high (63.4 eV), and we do not expect a signiÐcant fraction of Ne3` in HGS. In fact, photoionization models of Stasinska (1990) appropriate to HGS (B0B3 and B0C3, hereafter the reference models), do not predict

TABLE 15 M8 AND ORION GASEOUS ABUNDANCESa M8 This Work ELEMENT (1) He . . . . . . C ........ N........ O........ Ne . . . . . . S ........ Cl . . . . . . . Ar . . . . . . . Feg . . . . . .

t2 \ 0.00 (2)

t2 \ 0.032 (3)

PTD t2 \ 0.046 (4)

SM8Tb (5)

ORIONc t2 \ 0.024 (6)

^ ^ ^ ^ ^ ^ ^ ^ ^

10.99 8.49d 7.78 8.60e 7.95f 7.12 5.34 6.97f 6.13

11.00 8.55 7.80 8.74 7.99 7.38 ... 6.76 6.16

11.00 8.52 7.79 8.67 7.93 7.27 5.36 6.87 6.15

10.99 8.39 7.78 8.64 7.89 7.17 5.33 6.80 6.11

10.99 8.40 7.62 8.60 7.68 6.90 5.13 6.77 6.00

0.04 0.08d 0.18 0.15e 0.20f 0.10 0.22 0.11f 0.30

(4)

a In units of 12 ] log (X/H). b Mean value of cols. (3) and (4). c Esteban et al. 1998. d C`/H` from collisionally excited lines and C``/H` from recombination lines. e From recombination lines. f Uncertain i applied. cf g Adopting Fe`/H` from [Fe II] jj 7155 and 8617.

No. 1, 1999

ECHELLE SPECTROPHOTOMETRY OF M8

appreciable ionic fractions of Ne3`. In contrast, we expect to Ðnd a large amount of Ne`. Usually, the total Ne abundance for H II regions of a high degree of ionization is obtained from N(Ne) N(Ne``) \ i (Ne``) ] cf N(H) N(H`) \

N(O` ] O``) N(Ne``) ] . N(O``) N(H`)

(6)

In the case of HGS this i gives a correction of 1.91, which cf implies Ne`/Ne`` \ 0.91. On the other hand, the reference models give a mean Ne`/Ne`` \ 3.0 and imply that equation (6) does not apply to M8. Both ionic fractions and the Ne``/H` ratio obtained from our spectra give 12 ] log (Ne/H) \ 7.04 and 7.24, respectively, values too low compared with those typically observed in Galactic ionized nebulae. Pipher et al. (1984) observed the Ðnestructure line [Ne II] j12.81 km. They derived 12 ] log (Ne`/H`) \ 7.63 for HG but in a zone north of our position. Assuming this result as valid for HGS, we obtain Ne`/Ne`` \ 7.47. It is obvious that the ionization degree could not be the same in the zones in which both ions are observed, but the Ne`/Ne`` ratio obtained from the comparison between optical and IR observations is more consistent with the expected Ne abundance. Taking into account this, we adopt this highly uncertain i (Ne``) \ 8.47 to derive the total Ne abundance. The cf large discrepancy between the Ne`/Ne`` ratios obtained from the usual ionization correction factor and the IR lines indicates that large a fraction of Ne` coexists with O`` in HG. It must be taken into account that abundances derived from collisionally excited IR lines are essentially independent of t2 because the energy to excite the lines is considerably smaller than the kinetic energy of the electron. We have measured forbidden lines from two ionization stages of S, giving S`/S`` \ 0.08. An ionization correction factor, i (S` ] S``), to take into account the presence of S3` hascfto be considered : N(S` ] S``) N(S) . \ i (S` ] S``) ] cf N(H`) N(H)

(7)

The reference models predict an i (S` ] S``) \ 1.1. On cf observed the Ðnethe other hand, Pipher et al. (1984) structure line [S III] j \ 18.71 km and obtained an upper limit for [S IV] j \ 10.51 km. These authors found S3`/ S`` ¹ 0.03. We adopt this upper limit to estimate the S abundance. For Cl we have measurements of Cl` and Cl`` lines. The Cl`/Cl`` ratio is 0.17, which is very close to the value of 0.20 obtained by Osterbrock et al. (1992) for the Orion Nebula. Based on the models of Mathis & Rosa (1991) and our O`/O`` and S`/S`` ratios, we expect a Cl3`/Cl value of 0.22. This result implies that [Cl IV] j8046 of multiplet (1F) is observable, but we do not detect this line. With an estimated lower limit for detection of about 0.0001 ] I(Hb) in this spectral zone, we derive a Cl3`/Cl`` ¹ 0.02. Taking into account this estimate, we do not consider it necessary to adopt a substantial value for i (Cl` ] Cl``) to derive cf assume the total Cl abundance ; therefore, we N(Cl) N(Cl` ] Cl``) \ . N(H) N(H`)

(8)

127

From our data we obtain Ar3`/Ar`` \ 0.003 for HGS. This value is in contradiction with the predictions from the reference models, which give ratios between 0.22 and 0.25. A similar discrepancy was obtained also by PTD in the upper limit they estimated for HGS and by Peimbert, TorresPeimbert, & Ru• z (1992) when comparing M17 with the i cf values of Mathis & Rosa (1991). Pipher et al. (1984) observed [Ar II] j 6.99 km and [Ar III] j 8.99 km, obtaining an Ar`/Ar`` \ 3.56 ; this values is larger than the ratio of about 1.15 given by the reference models, which would imply a very high Ar abundance. In our case, we adopt an intermediate value between both estimates : Ar`/ Ar`` \ 2.36. We measure lines of two stages of ionization of Fe, with a Fe`/Fe`` value of 0.09. The contribution for the presence of Fe3` is expected to be important ; therefore, N(Fe) N(Fe` ] Fe``) \ i (Fe` ] Fe``) ] . cf N(H) N(H`)

(9)

We have adopted the i (Fe` ] Fe``) \ 1.70 given by the cf reference models. 10.2. Permitted L ines Helium has to be corrected for the presence of He0. Peimbert & Torres-Peimbert (1977) obtained an empirical equation for i (He`) based on the observed S`/S`` and cf O`/O`` ratios. From our observations and the expression of Peimbert & Torres-Peimbert we obtain i (He`) \ 1.27 cf identical to for HGS and, therefore, a He/H ratio of 0.098, the value derived by EPTE for the Orion Nebula. For C we have only direct determinations of C``. From IUE observations, PTD obtain the C`/H` ratio making use of the [C II] j2326 line. The reference models predict negligible contributions of ionization stages higher than C``. No correction factor is needed for this element, and, therefore, N(C) N(C` ] C``) . \ N(H`) N(H)

(10)

10.3. Comparison with Other Objects In Table 15 we present the gaseous abundances of M8 derived by us for t2 \ 0.000 and t2 \ 0.032 ; we also present the values obtained by PTD for t2 \ 0.046 and the abundances obtained for the Orion Nebula by EPTE for t2 \ 0.024. From the values obtained by us for t2 \ 0.032 and those obtained by PTD for t2 \ 0.046, we have deÐned an average value, SM8T, which is also presented in Table 15. In Table 16 we present the solar abundances and the abundances for B stars in clusters of the solar vicinity. With the exception of Fe it can be seen that the other heavy elements are about 0.2 dex lower in the B stars than in the Sun, a well-known result that implies that the Sun is superÈ metal rich for its age and its position in the Galaxy. To compare the SM8T abundances with those of the Orion Nebula and the Sun, it is necessary to estimate the fraction of heavy elements embedded in dust grains. We have assumed that the fraction of heavy elements trapped in dust is the same for M8 and Orion ; therefore, following EPTE we have added 0.10 dex, 0.08 dex, and 1.37 dex to the gaseous C, O, and Fe abundances, respectively. For He, Ne,

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Vol. 120

TABLE 16 M8,a ORION,a AND STELLAR ABUNDANCESb Element He . . . . . . C ....... N ....... O ....... Ne . . . . . . S........ Cl . . . . . . Ar . . . . . . Fe . . . . . .

Sunc

B Starsd

SM8T

SM8T [ Orion

SM8T [ Sun

^ ^ ^ ^ ^ ^ ^ ^ ^

... 8.35 7.80 8.69 ... 7.11 ... ... 7.51

11.00 8.62 7.79 8.75 7.93 7.24 5.36 6.87 7.52

]0.01 ]0.13 ]0.01 ]0.03 ]0.04 ]0.10 ]0.03 ]0.07 ]0.04

]0.01 ]0.07 [0.18 [0.12 [0.15 [0.06 [0.14 ]0.35 ]0.02

10.99 8.55 7.97 8.87 8.08 7.33 5.50 6.52 7.50

0.035 0.05 0.07 0.07 0.06 0.11 0.30 0.11 0.04

a Gas ] dust values. b In units of 12 ] log (X/H). c Grevesse, Noels, & Sauval 1996. d Snow & Witt 1996.

and Ar, no correction was applied since they are noble gases. For N, S, and Cl, no dust correction was applied since they are not signiÐcantly depleted in the neutral ISM (Savage & Sembach 1996). The SM8T abundances corrected for dust are presented in Table 16. There is a very small dispersion in the di†erences between SM8T and Orion for N, O, Ne, S, Cl, Ar, and Fe (see Table 16) ; all the di†erences are positive, which probably indicates the presence of a small composition gradient. The distance determinations for M8 and Orion (see PTD and references therein) indicate a radial galactocentric di†erence of 1.9 ^ 0.2 kpc, M8 being located at 6.5 kpc and Orion at 8.4 kpc from the Galactic center. An O/H gradient in the [0.05 to [0.09 dex kpc~1 range is expected from observations of Galactic H II regions (see Esteban & Peimbert 1996, and references therein), while from Table 16 the O/H gradient amounts to only [0.016 dex kpc~1, and the average value for the N, O, Ne, S, Cl, Ar, and Fe gradients amounts to [0.024 dex kpc~1. Considering the errors in the abundance determinations, the di†erences are not signiÐcant. We can also compare the gradients with Galactic chemical evolution models. From the model by Carigi (1996), based on the yields by Maeder (1992), a gradient of [0.045 dex kpc~1 for O/H is obtained. On the other hand, the predicted gradients for C/O and Fe/O amount to [0.053 dex kpc~1 and [0.058 dex kpc~1, respectively, in excellent agreement with the observed value for C/O that amounts to [0.053 dex kpc~1 but in poor agreement with the observed Fe/O gradient which is [0.005 dex kpc~1. The predicted large C/O gradient is due mainly to the increase of the C to O yields ratio with metallicity in massive stars. For the Fe/O ratio, in addition to the yields by Maeder, Carigi took into account the production of Fe owing to Type I supernovae. In Table 16 we also compare SM8T with the solar values. SM8T shows an average deÐciency in N, O, Ne, S, and Cl of 0.14 dex relative to the Sun. The large Ar/O di†erence should be investigated further and probably is due mainly to the uncertainty in the collisional cross section for the IR [Ar II] lines, which have been used to estimate the i cf adopted for this element. The large C/O and Fe/O excesses between SM8T and the Sun probably are due to Galactic chemical evolution. While the observed di†erences amount to 0.19 dex and 0.14 dex, respectively, the model by Carigi (1996) predicts 0.25 dex and 0.25 dex for both ratios. In addition to the radial gradient e†ect mentioned before,

there is the more important age e†ect, due to the age di†erence between the Sun and M8, which goes in the direction of increasing the C/O and Fe/O di†erences. 11.

SUMMARY

We present echelle spectroscopy in the 3500È10300 AŽ range of the Hourglass Nebula in M8. We have measured the intensities of about 274 emission lines ; 88 of them are permitted lines of heavy elements. We have determined physical conditions of M8 and its chemical composition. We derive the He`, C``, O`, and O`` ionic abundances based on recombination lines. These abundances do not depend on the temperature structure of the nebula ; therefore, they are more reliable than those derived from collisionally excited lines. Taking into account the low ionization degree of M8 that indicates the presence of O in the O` and O`` stages only, we have obtainedÈ for the Ðrst time in an ionized nebulaÈthe total O abundance based exclusively on recombination lines. We have obtained an average t2 \ 0.032 ^ 0.019 by comparing the C``, O`, and O`` ionic abundances obtained, making use of both collisionally excited lines and recombination lines. This t2 value has been used to determine chemical abundances from forbidden lines. The chemical abundances of N, O, Ne, S, Cl, Ar, and Fe in M8 are slightly higher than in Orion, which is consistent with the presence of a small composition gradient. The gradient in C/O is somewhat higher than and in excellent agreement with the prediction of a recent chemical evolution model for the Galaxy. An underabundance of 0.14 dex is found for most heavy elements in M8 relative to the Sun. The large C/O and Fe/O excesses between M8 and the Sun could be due to Galactic chemical evolution. We would like to thank M. Bautista for providing us his latest calculations for Fe and Ni and L. Carigi for several fruitful discussions. We are grateful to the referee, Gary Ferland, for his useful suggestions. C. E. would like to thank all the members of the Instituto de Astronom• a, UNAM, for their warm hospitality during his stays in Mexico. This research was partially funded through grant PB94-1108 from the Direccion General de Investigacion Cient• Ðca y Tecnica of the Spanish Ministerio de Educacion y Ciencia and grants DGAPA-UNAM IN109696 and CONACyT 25451-E (Mexico).

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