Midcourse Space Experiment Spectra of the Orion Nebula and the ...

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ABSTRACT. Spectra of the Orion Nebula were obtained with the Midcourse Space Experiment Spirit III interfer- ometer from 370 to 2000 cm~1 with 2 cm~1 ...
THE ASTROPHYSICAL JOURNAL, 508 : 268È274, 1998 November 20 ( 1998. The American Astronomical Society. All rights reserved. Printed in U.S.A.

MIDCOURSE SPACE EXPERIMENT SPECTRA OF THE ORION NEBULA AND THE IMPLICATIONS FOR ABUNDANCES IN THE INTERSTELLAR MEDIUM J. P. SIMPSON,1,2 F. C. WITTEBORN,1,2 S. D. PRICE,3 AND M. COHEN4 Received 1998 February 18 ; accepted 1998 June 25

ABSTRACT Spectra of the Orion Nebula were obtained with the Midcourse Space Experiment Spirit III interferometer from 370 to 2000 cm~1 with 2 cm~1 resolution in a 6@ ] 9@ Ðeld of view (FOV) in 1996 November. Lines were detected of [S III] 534.4 cm~1, [Ne III] 642.9 cm~1, [Ne II] 780.4 cm~1, [S IV] 951.4 cm~1, [Ar III] 1112.2 cm~1, [Ar II] 1431.6 cm~1, H (7È6) 808.3 cm~1, H (8È6) 1332.9 cm~1, H (6È5) 1340.5 cm~1, H (S1) 587.0 cm~1, H (S2) 814.4 cm~1, H (S3) 1034.7 cm~1, H (S4) 1246.1 cm~1, and 2 2 2 2 H (S5) 1447.3 cm~1. The following abundances were determined from these lines : Ne/H \ 2 9.9 ^ 1.1 ] 10~5, S/H \ 8.1 ^ 1.1 ] 10~6, and Ar/H \ 2.5 ^ 0.2 ] 10~6. These abundances are all less than solar and conÐrm that the Sun is overabundant in heavy elements without the need for correction for the composition of interstellar dust. The low sulfur abundance compared with solar is an indication that a signiÐcant amount of the sulfur in Orion is in dust grains. The FOV-averaged molecular hydrogen column density is D1.6 ] 1020 cm~2 for an excitation temperature of D670 K and an extinction correction corresponding to an optical depth of 1.5 at 9.7 km. The unidentiÐed infrared emission features at 6.2, 7.7, 8.6, 11.3, and 12.7 km, attributable to polycyclic aromatic hydrocarbons, were also detected. A prominent, broad silicate feature centered near 18 km and additional weak features were detected and are discussed. Subject headings : infrared : ISM : continuum È infrared : ISM : lines and bands È ISM : individual (Orion Nebula) 1.

INTRODUCTION

solid silicate and carbonaceous material (dust). Additional weak lines and continuum features were also detected. Section 2 discusses the observations and data reduction ; ° 3.1 uses the forbidden line and hydrogen recombination line measurements to determine the abundances of neon, sulfur, and argon and compares the derived abundances with previously measured abundances in Orion and the Sun. Section 3.2 uses the observed molecular hydrogen lines to estimate an average extinction, excitation temperature, and column density for H , and ° 3.3 discusses the carbon2 containing molecular features.

The Midcourse Space Experiment (MSX) is a Ballistic Missile Defense Organization satellite that has a number of experimental objectives, one of which was to characterize the celestial sky, including midinfrared (IR) spectroscopy of the brightest sources (Mill et al. 1994). The Orion Nebula (M42, NGC 1976) is the brightest optically visible H II region and is ionized by the Trapezium cluster, the brightest star of which is the O6p star h1C Ori. The H II region is the front surface of the giant molecular cloud OMC1. Within the cloud is the optically obscured star-forming region, called BN/KL for its brightest near-IR and far-IR members (see Genzel & Stutzki 1989). Because it is the closest (450 pc) region of massive star formation and contains some of the brightest and best-resolved sources, the Orion region has long been one of the Ðrst targets to be studied with any new telescope or instrument. This paper presents the MSX Spirit III interferometer observations of the Orion Nebula taken on 1996 November 11. The spectra cover the frequency range from D390 to 1800 cm~1 (6È26 km) at 2 cm~1 resolution. Thus, these are the Ðrst observations of the six bright forbidden lines from Ne`, Ne``, S``, S3`, Ar`, and Ar`` and adjacent hydrogen recombination lines ever measured with the same instrument in the Orion Nebula. The wide wavelength coverage also permits the observation of the broad, strong, IR features attributed to

2.

OBSERVATIONS

The primary instrument during the cryogenic phase of the MSX mission was the Spatial IR Imaging Telescope (SPIRIT) III, an o†-axis telescope cooled with solid hydrogen. SPIRIT III had Ðve linear arrays of mid-IR detectors, each Ðltered for a di†erent color (Mill et al. 1994 ; Egan et al. 1998), and a Ðeld-shared Michelson interferometer with six Si : As impurity-band conduction detectors. The e†ective primary aperture used by the interferometer was 24.68 cm. The detector and Ðlter characteristics were tailored for observing the emission from the EarthÏs atmosphere, and all but one detector were Ðltered to limit the spectral response. A helium-neon laser was used to time the sampling such that the maximum frequency was 3949.5 cm~1. The focal plane of the interferometer contained a column of three [email protected] ] [email protected] (nominal) detectors on the centerline and three 13@ ] 13@ detectors Ðltered with much narrower spectral response on the sides. Of the three smaller detectors, the middle detector covered the 3È5 km region, the top detector (Detector 4) had a response from 2.7 to 26 km, while the bottom detector (Detector 6) covered the same wavelength but had a notched Ðlter to reject radiation from the terres-

1 NASA/Ames Research Center, Astrophysics Branch, MS 245-6, Mo†ett Field, CA 94035-1000. 2 SETI Institute. 3 Air Force Research Laboratory, VSBC, Hanscom AFB, MA 017313010. 4 Radio Astronomy Laboratory, University of California, Berkeley, CA 94720-3411.

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MSX SPECTRA OF THE ORION NEBULA trial 15 km CO band. The satellite was oriented in space so 2 that the Sun was always toward the bottom. MSX was launched on 1996 April 24 and ran out of cryogen in 1997 February. The Orion Nebula observations were obtained by sequentially centering the 6@ ] 9@ Ðeld of view (FOV) (right ascension ] declination) of each of the three smaller MSX interferometer detectors on the Trapezium for 8 minutes each. The FOV was aligned in right ascension such that the Sun and Detector 6 were east of Detector 4 at all times. The pointing accuracy was [email protected]. The FOV of the detectors included almost all the ionized nebula (e.g., Felli et al. 1993), the BN/KL region of the molecular cloud immediately to the northwest, and the photodissociation region (PDR) between the H II region and the molecular cloud. Approximately 110 double-sided interferograms were taken by each of the three detectors. We report on the results from the two detectors with the largest spectral bandwidth (nominally D2.5È28 km, Detectors 4 and 6). The detector responsivity and source intensity were such that essentially no signal was detected in the 3È5 km detector and at wavelengths less than 5 km in the other detectors. The interferograms were processed with the standard MSX data reduction package, Convert 5.0, and transformed to spectra with the default Kaiser-Bessel apodization at 2 cm~1 resolution. For 4186 samples on either side of zero path, the apodization function used, A(i), for sample number i is A(i) \ I Mb[1 [ (i/4186)2]N/I (b), where 0 0 and the I is the modiÐed Bessel function of the Ðrst kind 0 value for b (7.283997681) was chosen so that the Ðrst sidelobe in the spectral domain is smaller than the peak by a factor of 100. The averaged spectra are plotted in Figures 1 and 2. The following lines are identiÐed in the spectra : strong atomic forbidden lines of neon, sulfur, and argon (Fig. 1a), atomic hydrogen recombination lines, and molecular hydrogen rotational lines (Figs. 1d and 1e). The broad solid features at 790, 890, 1160, 1300, and 1610 cm~1 are attributable to ionized polycyclic aromatic hydrocarbons (PAHs). The iron/magnesium silicates that are known to exist in the Orion Nebula have broad features at 10 and 18 km ; the 10 km feature is not obvious in these spectra because of a combination of multiple dust temperatures and optical

depth e†ects, but the broad emission at D16È22 km may arise from the 18 km feature, which has a smaller optical depth (see the model III spectra of Simpson & Rubin 1990). The small emission feature at D980È1100 cm~1 is too narrow to be the 10 km silicate feature as seen in the Trapezium (Hayward, Houck, & Miles 1994). Weak PAH features include features at 910 and 1140 cm~1. All of these PAH features (sometimes called the ““ unidentiÐed IR emission features ÏÏ or UIR features) have been detected previously in other sources (Allamandola, Tielens, & Barker 1989 ; Witteborn et al. 1989 ; Beintema et al. 1996 and references therein). Unfortunately, there are ““ channel ÏÏ spectra throughout the whole spectrum ; these are most noticeable (compared with the noise) between 390 and 860 cm~1. Both the calibration and the data interferograms exhibit strong additional fringes at about 2800 samples from zero path : these are the cause of the channel spectra. Thus, the channel spectra can be largely eliminated by smoothing the spectra through truncating the interferograms or they can be partially removed by subtracting a cosine function ; Figures 1a, 1d, and 2 have an arbitrary cosine function subtracted (frequencies corresponding to 2792 and 2796 samples for Detectors 4 and 6, respectively, or 2.8292 and 2.8251 cm~1). However, the channel spectra contain many frequencies (from multiple fringes, not a single spike), and subtraction of a cosine is not adequate when one is measuring faint lines. Figures 1b and 1c are the result of truncating the interferograms to 2600 samples before the Kaiser-Bessel apodization. Now the lines can be more easily seen, but the Ñux determination is not good because of the low resolution and the curved baseline. The following weak lines are found in this range : [Fe III] 436.2 cm~1 (22.9 km), [Ar III] 458.1 cm~1 (21.8 km), H (8È7) 524.6 cm~1 (19.1 km), [P III] 559.6 cm~1 (17.9 km), and H (S1) 587.0 cm~1 (17.0 km). 2 The line intensities were measured by summing over the line proÐle and subtracting an average continuum and are given in Table 1. The errors in the table are statistical only and were determined from the formal statistics in the averaging of the approximately 100 spectra and the estimated uncertainties in determining the continuum. In the absence of a frequency-dependent calibration uncertainty, these statistical errors would be the appropriate ones to use in the determination of line ratios and derivative quantities

TABLE 1 LINE INTENSITIES Line

Frequency (cm~1)

Wavelength (km)

[S III] . . . . . . . . 534.4 18.71 [Ne III] . . . . . . . 642.9 15.56 [Ne II] . . . . . . . 780.4 12.81 [S IV] . . . . . . . . 951.4 10.51 [Ar III] . . . . . . . 1112.2 8.99 [Ar II] . . . . . . . . 1431.6 6.99 H (7È6) . . . . . . . 808.3 12.37 H (8È6)b . . . . . . 1332.9 7.50 H (6È5) . . . . . . . 1340.5 7.45 H (S1) . . . . . . . . 587.0 17.03 H2(S2) . . . . . . . . 814.4 12.28 H2(S3) . . . . . . . . 1034.7 9.66 2 H (S4) . . . . . . . . 1246.1 8.03 2 H (S5) . . . . . . . . 1447.3 6.91 2 a Errors are statistical only. b Blended with H (11È7) at 1331.9 cm~1.

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Detector 4 (W cm~2 sr~1)

Detector 6 (W cm~2 sr~1)

13.98 ^ 0.04([10)a 11.17 ^ 0.16([10) 12.93 ^ 0.71([10) 4.15 ^ 0.45([10) 4.72 ^ 0.07([10) 1.21 ^ 0.06([10) 2.03 ^ 0.28([11) 2.68 ^ 0.31([11) 8.31 ^ 0.41([11) 1.23 ^ 0.19([11) 1.28 ^ 0.15([11) 2.59 ^ 0.25([11) 2.29 ^ 1.56([11) 3.14 ^ 0.60([11)

15.20 ^ 0.03([10) 9.31 ^ 1.11([10) 12.98 ^ 0.92([10) 5.38 ^ 0.16([10) 5.52 ^ 0.06([10) 1.59 ^ 0.10([10) 2.05 ^ 0.18([11) 2.69 ^ 0.34([11) 8.30 ^ 0.37([11) 1.45 ^ 0.15([11) 0.63 ^ 0.18([11) 2.99 ^ 0.26([11) 2.21 ^ 0.30([11) 5.15 ^ 0.40([11)

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FIG. 1.ÈThe MSX spectra of the Orion Nebula are plotted as a function of frequency. (a) The spectrum from Detector 4 is plotted as the solid line, and the spectrum from Detector 6 is the dashed line. The detected forbidden lines are marked. (b) and (c) The spectra were smoothed by truncating the interferograms before apodization (see text). In order to make the weak lines more noticeable on the plot, the spectra were Ñattened by multiplying by 10~4 times the frequency squared. The spectrum from Detector 6 has the larger Ñux. (d) and (e) The hydrogen recombination lines and the molecular hydrogen pure rotational lines are marked. The spectrum from Detector 6 is displaced upward by adding the constant, 5.0 ] 10~11 W cm~2 sr~1 cm. The dotted lines are the standard deviation of the mean from the coaddition of the D100 spectra for each detector.

such as abundances and excitation temperatures. However, the MSX calibration certiÐcation (S. Mazuk & R. Russell 1998, private communication) gives a slope uncertainty in the calibration of 20% from 4 to 23 km for Detector 4 and

15% for Detector 6. This corresponds to an uncertainty of 12% for the ratio of the 534 cm~1 [S III] line to the 1341 cm~1 Pfa line and a smaller uncertainty for lines closer together in frequency. This additional uncertainty was

MSX SPECTRA OF THE ORION NEBULA

No. 1, 1998

FIG. 2.ÈThe MSX spectra of the Orion Nebula are plotted as a function of wavelength.

added quadratically to the statistical errors in Table 1 in calculating the errors for ratios of lines. For comparison with other observations, an additional calibration uncertainty of D10% (S. Mazuk & R. Russell 1998, private communication) must be added to all the intensities. 3.

ANALYSIS AND DISCUSSION

3.1. Abundances of Neon, Sulfur, and Argon Abundances of elements heavier than helium are typically determined by comparing collisionally excited forbidden lines with hydrogen recombination lines or the radio freefree continuum (see Osterbrock 1989 for a complete discussion). The result of the comparison is the ionic abundance ratio with respect to H`. Since most elements in H II regions (deÐned to be a region in which all hydrogen is H`) are found in more than one ionization state, it is necessary to measure or estimate the ionic abundances of the several ionization states in order to determine the total abundance. The Orion Nebula is the standard for interstellar abundances, just like the Sun is the standard for stellar abundances (e.g., Wilson & Rood 1994). The most complete recent sets of abundance determinations for the Orion Nebula are those of Osterbrock, Tran, & Veilleux (1992) and Esteban et al. (1998). These determinations are based on optical measurements of an enormous number of lines ; however, the only element where all of the major ionization states have optical forbidden lines is oxygen. Consequently, the best-determined elemental abundance in H II regions is O/H, although even for oxygen the abundance determination is somewhat controversial because it requires a good determination of both the electron temperatures in the different ionization zones of the H II region and the mean squared temperature Ñuctuations (Peimbert 1993 ; Mathis 1996 ; Esteban et al. 1998). Many elements have one ionization stage observable in the optical and another in the midor far-IR. In this case, the Ðelds of view must sample exactly the same volume elements of the H II region ; this is a difficult task unless the H II region is either compact or distant. Otherwise, one must estimate the abundances of the unseen ionization states either through the use of detailed models (e.g., Baldwin et al. 1991 ; Rubin et al. 1991) or through the use of ionization correction factors based on ionization

271

potentials (e.g., Peimbert & Torres-Peimbert 1977). However, because of di†erences in the choice of data to be modeled and other assumptions, the abundances that best Ðt the di†erent models have a surprisingly large range. The MSX interferometer is the Ðrst instrument to adequately observe the neon, sulfur, and argon lines. The detector FOV was large enough to include the entire source and the spectral range wide enough to include two ionization stages for each element plus several hydrogen recombination lines. Because the two ionization stages include almost all of the sulfur and argon and all of the neon, no modeling is necessary. This is in contrast to measurements at optical wavelengths where there are lines of Ne``, S`, S``, and Ar`` but not Ne`, S3`, or Ar`. The other advantages of mid-IR forbidden line observations are that the line emissivities are insensitive to both electron temperature, T , and electron density, N , except at the highest e , and the Orion Nebula e contains very little values of N e material with N [ 104 cm~3. Because extinction in the IR e is substantially lower than that in the optical and the Orion Nebula is optically visible, no corrections for extinction were deemed necessary. The result is that neither the optical uncertainties of extinction and electron temperature variations nor the far-IR uncertainties of density Ñuctuations (Simpson et al. 1995) make a large contribution to the systematic uncertainties of abundance analyses. There are still, of course, uncertainties because of any unknown systematic errors in the calibration. The ionic abundances were calculated from the ratios of the forbidden line intensities divided by the intensity of the strongest hydrogen line, Pfa (see Simpson & Rubin 1990 ; Simpson et al. 1995). Any absolute calibration uncertainties owing to the size of the FOV or the pointing cancel. Electron temperatures of 8500 K (Wilson & JaŽger 1987 ; Esteban et al. 1998 for the Balmer continuum) and electron densities of 4000 cm~3 (Osterbrock et al. 1992 ; Esteban et al. 1998) were used. The electron densities and temperatures were chosen to be those appropriate to the lines that are the most sensitive : [S III] for density and hydrogen Pfa for temperature. There should be a contribution to the abundance errors for the uncertainties in these assumptions : we assumed an uncertainty of 500 K in T and 1000 cm~3 in N e for T and 1%È8%e and added in quadrature errors of D9% e He I recomfor N depending on the line. Because there are e bination lines at the same frequencies as the hydrogen lines, the hydrogen line intensities were decreased by an estimated 8.5% (Smits 1991, 1996). The total abundances were calculated by summing the abundances in the two observed ionization stages. The ionic and total abundances are given in Tables 2 and 3, respectively. The hydrogen emissivities were taken from Storey & Hummer (1995), and the references for the collisional excitation cross sections for the forbidden lines are given in Table 2. The missing ionization stages that need to be considered are S` and Ar3`. Since the ionization potential of S0 is only 10.36 eV, at least some of the S` observed in the Orion Nebula comes from the PDR, where hydrogen is neutral. Consequently, the S`/H` \ 5.0 ] 10~7 observed by Osterbrock et al. (1992) and Esteban et al. (1998) is, in fact, an upper limit to the S`/H` that should be added to the (S`` ] S3`)/H` to get the total H II region S/H. Esteban et al. measured Ar3`/ H` \ 0.5 ] 10~7. Thus the S/H and Ar/H ratios in Table 3 are really lower limits ; the Ðnal values are probably no more than S/H \ 8.6 ] 10~6 and Ar/H \ 2.6 ] 10~6.

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Vol. 508

TABLE 2 IONIC ABUNDANCE RATIOS Ionic Ratio

Detector 4

Detector 6

Weighted Average

Reference

Ne`/H` . . . . . . . Ne``/H` . . . . . . S``/H` . . . . . . . S```/H` . . . . . . Ar`/H` . . . . . . . . Ar``/H` . . . . . .

71.9 ^ 8.4([6) 27.9 ^ 3.5([6) 7.1 ^ 1.2([6) 0.51 ^ 0.08([6) 0.40 ^ 0.04([6) 1.86 ^ 0.17([6)

72.3 ^ 8.5([6) 23.3 ^ 3.8([6) 7.8 ^ 1.2([6) 0.66 ^ 0.07([6) 0.53 ^ 0.06([6) 2.18 ^ 0.19([6)

72.1 ^ 7.6([6) 27.3 ^ 3.2([6) 7.5 ^ 1.2([6) 0.63 ^ 0.07([6) 0.46 ^ 0.05([6) 2.04 ^ 0.17([6)

1 2 3 4 5 3

REFERENCES.È(1) Saraph & Tully 1994 ; (2) Butler & Zeippen 1994 ; (3) Galavi s, Mendoza, & Zeippen 1995 ; (4) Johnson, Kingston, & Dufton 1986 ; (5) Pelen & Berrington 1995.

Table 3 also compares the MSX-observed abundances with previously measured abundances in Orion and in the Sun. The previously measured Orion abundances of Ne, S, and Ar exhibit great scatter, no doubt because of the reasons mentioned above. In recent years it has become accepted (e.g., Mathis 1996) that the reference interstellar abundance (usually taken as the C, N, and O abundances in the Orion Nebula ; e.g., Wilson & Rood 1994) is substantially less than solar. However, for C, N, and O, arguments have to be made about the fraction of the heavy elements in the interstellar medium that reside in dust. Our measurements in Table 3 show that the abundances of the noble gases, neon and argon, are also less than solar, although the di†erence is not as signiÐcant for neon (see Widing 1997) as it is for the other heavy elements. Thus our measurements conÐrm that the Sun is overabundant in heavy elements without the need for corrections for the composition of interstellar dust. It is also interesting that the Orion S/H ratio is as low as the Orion O/H ratio compared with the solar system ; this is possibly an indication that a signiÐcant amount of the sulfur in Orion is in dust grains if the ratios of the heavy element abundances are the same in the Sun and the Orion Nebula.

since the ratios of di†erent previously observed pairs of lines correspond to di†erent excitation temperatures, T : ex the higher the energies of the upper levels, the higher the apparent T . The observers of the emission near BN/KL ex have compared the column densities in their beams with the predictions of shocks of di†erent velocities in order to characterize the outÑows from the stars (e.g., Burton & Haas 1997). On the other hand, the temperatures of the shocked gas are not so high as to thermally excite the higher (l º 2) vibrational levels ; copious emission from l \ 2È1 is a good indicator of Ñuorescent excitation (Usuda et al. 1994). The MSX interferometer FOV is so large that it includes both the shocked emission from BN/KL and the Ñuorescent emission from the Orion PDR. There are Ðve rotational lines of H in the ground vibra2 tional state seen in either or both spectra with a signal-tonoise ratio greater than 3 (Fig. 1). Although the H Ñuores2 cent emission from the PDR probably su†ers very little extinction, the shock-excited emission from the KL region possibly has an extinction of as much as 1È2 mag at 2.12 km (Parmar, Lacy, & Achtermann 1994 ; Burton & Haas 1997). Following the technique of Parmar et al. (1991) and Wesselius et al. (1996) with transition probabilities from Turner, Kirby-Docken, & Dalgarno (1977), one can estimate the excitation temperature, T , by plotting the ex number of molecules in the upper energy level, N , divided by the statistical weight, g , as a function ofupper the upper level energy, E , that is, upper upper N I upper \ l , (1) g (hl/4n)Ag upper upper where I is the intensity of the line and A is the transition l probability. A Ðt for T is easily obtained, since ex N /g is expected to decrease exponentially as a funcupper tion of Eupper /T . upper ex

3.2. Molecular Hydrogen The Orion Nebula and environs have been observed extensively in the rotational and vibrational-rotational lines of H (Burton & Haas 1997 ; Parmar, Lacy, & Achtermann 19942and references therein). The brightest emission comes from the region near BN/KL, where it is thought that the gas is excited by the shocks from the winds and outÑows of the young stellar objects in the region ; there is also substantial Ñuorescent emission from the PDR between the H II region and the molecular cloud (Usuda et al. 1996). The molecular gas clearly includes a large temperature range,

TABLE 3 ELEMENTAL ABUNDANCE RATIOS Element

MSX Determination

Ne/H . . . . . . S/H . . . . . . . Ar/H . . . . . . C/H . . . . . . . N/H . . . . . . . O/H . . . . . . .

9.9 ^ 1.1([5) 8.1 ^ 1.2([6) 2.5 ^ 0.2([6) ... ... ...

Other Orion 8.1,a 8.5,a 4.5,a 2.8,a 6.8,a 4.0,a

39.5,b 4.0,c 7.8d ([5) 13.3,b 9.4,c 14.8d ([6) 2.1,b 2.6,c 6.3d ([6) 2.1,b 2.5d ([4) 8.7,b 5.2,c 6.0d ([5) 3.8,b 3.1,c 4.3d ([4)

a Rubin et al. 1991 ; Rubin, Dufour, & Walter 1993. b Baldwin et al. 1991. c Osterbrock et al. 1992. d Esteban et al. 1998. e Widing 1997. f Grevesse, Noels, & Sauval 1996.

Solar System 1.2([4)e 1.6([5)f 3.5([6)f 3.5([4)f 9.3([5)f 7.4([4)f

No. 1, 1998

MSX SPECTRA OF THE ORION NEBULA

Even though the extinction is not the same for the di†erent H -emitting volumes, some correction must be estimated,2 since the S(3) line occurs at a wavelength very near the maximum of the 9.7 km silicate feature. Weighted leastsquares Ðts were calculated for the slope and intercept of the log (N /g ) versus E relation for di†erent test upper upper upper values of the extinction. An extinction law proportional to the extinction at 9.7 km was used : for S(1) to S(5), the proportionality constants are 0.37, 0.40, 1.0, 0.29, and 0.25 (see Draine & Lee 1984 with an additional contribution from icy mantles at 6.9 km). The best Ðt (minimum s2) was found for q \ 1.54, and the extinction-corrected values of 9.7 km N /g are plotted in Figure 3. The solid line correupper upper sponds to T \ 674 K and a column density of ex D1.6 ] 1020 cm~2. It was assumed that the ortho/para ratio is given by the statistical weights. Lower values of the extinction produce Ðts with higher T and lower column ex densities. It might seem that a value of q \ 1.5 is too large, 9.7the km spectra. However, since no 10 km absorption is seen in this absorption would occur in the Orion molecular cloud itself, which is behind the PDR that does have the 9.7 km feature in emission (e.g., Hayward et al. 1994). Moreover, a value of q \ 1.5 is close to what would be found if the 9.7were km Ðtted with the models of Simpson & Rubin continuum (1990). Clearly, the H emission that is seen is the sum of extinguished emission2from the molecular cloud and nonextinguished emission from the PDR. Thus, these FOVaveraged values for extinction and T may not be especially meaningful. However, they shouldex be representative of other, more distant H II regions that include both a PDR and a star-forming region with outÑows. 3.3. PAHs The spectrum of the Orion Nebula (Fig. 1) is the Ðrst obtained of such a large portion (6@ ] 9@) at such high spectral resolution (2 cm~1). The PAH features (see Allamandola et al. 1989 for a full discussion) are very distinct and appear to be well resolved. The relative strengths of the bands at 6.2, 7.7, 8.6, and 11.3 km are 4, 10, 3, and 3, approximately the same in the large FOV of MSX as in the

FIG. 3.ÈH excitation diagram for an assumed extinction at 9.7 km of q \ 1.54. 2The diamonds are the Detector 4 H lines and the squares 9.7 km 2 density of 1.6 ] 1020 are the Detector 6 lines. The solid line is for a column cm~2 at T \ 674 K. ex

273

much smaller (21A diameter) FOV used by Bregman et al. (1989) in their study of the Orion Bar at its brightest location, ““ position 4.ÏÏ In a study of several H II regions as well as planetaries and reÑection nebulae, Cohen et al. (1986 and 1989) found signiÐcant di†erences in the relative strengths of the bands in di†erent classes of objects. Bregman et al. modeled the geometry of the Bar as a PDRÈH II region seen partially edge-on. This model (their Fig. 8) suggests that the PAH emission comes from a very extended PDR between the molecular cloud and the H II region. The optical depth of the PDR is uniform over a large angular area, but a bend in the surface increases our line-of-sight optical depth by roughly a factor of 3. This gives rise to the Bar. The strength measured by Bregman et al. of the 7.7 km feature in the Bar is 2.4 times that of a region of similar angular size in our spectra of a 6@ ] 9@ region surrounding the Bar. Since the Bar spectrum was chopped against the surrounding emission, 2@ distant, it is really 0.3 brighter in comparison with the data obtained with MSX. Thus, the strength in the Bar is about 3 times the speciÐc intensity in the PAH bands of the surrounding region, which is in good qualitative agreement with the model. The underlying continuum at wavelengths beyond about 8.5 km is stronger than would be expected from the spectrum of the Bar alone. The large area covered clearly includes knots of emission from other sources such as the Trapezium region (Hayward et al. 1994), which have speciÐc intensities at 10 km in small Ðelds of view up to 200 times what we measured over a 6@ ] 9@ region. The PAH features in Figures 1 and 2 may be compared with recent Infrared Space Observatory (ISO)/Short Wavelength Spectrometer spectra of the circumstellar PAH emission bands in HR 4049, a cool post-asymptotic giant branch (AGB) star, the C-rich planetary nebulae NGC 7027 and IRAS 21282]5050 (Beintema et al. 1996 ; Molster et al. 1996), and a number of H II regions (Roelfsema et al. 1996). The 7.7 km feature itself peaks at 7.6 km ; this was noted earlier by Bregman (1989) to be characteristic of H II regions and is also seen in the Orion Bar (Bregman et al. 1989) and in the H II regions studied with ISO. In the AGB stars and the planetary nebulae, the feature peaks at 7.8 km. In the Orion spectra, the 7.7 km feature is clearly separated from the 8.6 km feature with relatively little continuum between them. The 6.2 km feature is similar to the feature as seen in the other sources (Molster et al. 1996 ; Roelfsema et al. 1996). The interesting weak feature that occurs at 9.6 km in Orion is also seen in HR 4049 and the H II region, IRAS 18434[0204, and would be lost even to airborne telescopes because of terrestrial ozone. Features in the 11È13 km region are attributable to the CÈH out-of-plane bending mode of hydrogen atoms attached to the carbon rings (Allamandola et al. 1989 ; Witteborn et al. 1989). Ionization of the PAH molecules tends to shift the location of the feature to shorter wavelengths or higher frequencies (Hudgins & Allamandola 1995a, 1995b ; Hudgins, Allamandola, & Sandford 1998). Thus, the 11.3 km band (890 cm~1) could be due to either a single hydrogen atom in a neutral molecule or two adjacent hydrogen atoms in an ionized molecule. The feature at 11.0 km (910 cm~1) is due to a single hydrogen atom on an ionized PAH molecule. The features at 11.9 km (840 cm~1) and 12.7 km (790 cm~1) are due to one and two adjacent hydrogen atoms, respectively (Allamandola et al. 1989 ; Witteborn et al. 1989). Additional hydrogen atoms attached

274

SIMPSON ET AL.

to the ring should produce features at wavelengths longer than 13 km (Allamandola et al. ; Hudgins et al.). There is a hint of such a feature at D13.5 km (740 cm~1) in the Detector 4 spectrum ; however, it cannot be conÐrmed in the Detector 6 spectrum because of a low signal-to-noise ratio at the CO notch in the Ðlter. We Ðnd intensities for the 2 weak bands, relative to the 7.7 km band strength, at 9.6, 11.0, 12.7, and 13.5 km to be 0.15, 0.03, 0.10, and 0.04, respectively. The asymmetry of the emission feature near 6.2 km, which is attributed to CÈC stretch (Leger & Puget 1984), and the strong 11.3 km emission feature, which is attributed to out-of-plane CÈH bending (Duley & Williams 1981), is very evident in Figures 1a, 1d, and 1e. The asymmetry of the 11.3 km feature has been observed before (e.g., by Aitken & Roche 1982) and is ascribed to anharmonicity of the out-ofplane CÈH vibrations (Barker, Allamandola, & Tielens 1987). Asymmetry in the 6.2 km feature was observed by Cohen et al. (1989) and is ascribed to anharmonicity of the CÈC stretching modes by Allamandola et al. (1989). Since the CÈC stretch feature appears at longer wavelengths (e.g., 6.3 km) in many laboratory PAHs, some of the observed width may be caused by blending of emissions from di†erent PAHs. 4.

SUMMARY

The present observations make an important contribution to the knowledge of the abundances of the interstellar standard, the Orion Nebula : the Orion abundances are signiÐcantly lower than the solar system abundances. The abundances measured and quoted here are all gas phase for Orion, although at least some of the C, O, and S must be in the solid phase in dust and do not contribute to the gas-

phase abundances The MSX observations show that the noble gases Ne and Ar also have low abundances compared with solar. There are abundance gradients in the Galaxy (e.g., Simpson et al. 1995) ; however, at a distance of 500 pc from the Sun, the Orion Nebula is not signiÐcantly farther from the Galactic center compared with the Sun. Moreover, observations of stars in the Orion region give C, N, and O abundances similar to the Nebula (Cunha & Lambert 1994). We conclude that the Sun is overabundant in heavy elements compared with the local interstellar medium. Five rotational lines of molecular hydrogen were observed. Since these lines arise from both the shockexcited region around BN/KL and the Ñuorescent excited gas of the PDR, the observed excitation temperature and column density represent an average of what would be expected from combining the two regions in the same beam. Numerous PAH features are observed ; these probably also arise in the PDR between the H II region and the molecular cloud. The features are well resolved : the 7.7 km feature contains two components, the stronger of which is at 7.6 km. The 8.6 km feature is relatively narrow and well separated from the 7.7 km feature. Weak features at 9.6, 11.0, 12.0, 12.7, and possibly 13.5 km are detected. We thank N. Bonito, A. El Fakih, and K. Farnham for their help with running Convert and R. Russell for many helpful discussions regarding the data processing. The MSX is sponsored by the Ballistic Missile Defense Organization (BMDO). We thank the referee for his helpful comments on the presentation. We gratefully acknowledge the steadfast support of the BMDO. J. P. S. and F. C. W. acknowledge support from NASA/Ames Research Center Research Interchange grant NCC 2-900.

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