Observations of Molecular Hydrogen in the Carina Nebula

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DAE-HEE LEE,1 KYOUNG-WOOK MIN,1,2 W. VAN DYKE DIXON,3 MARK HURWITZ,3 KWANG-SUN RYU,2 KWANG-IL SEON,4. AND JERRY EDELSTEIN3.

THE ASTROPHYSICAL JOURNAL, 545 : 885È891, 2000 December 20 ( 2000. The American Astronomical Society. All rights reserved. Printed in U.S.A.

OBSERVATIONS OF MOLECULAR HYDROGEN IN THE CARINA NEBULA DAE-HEE LEE,1 KYOUNG-WOOK MIN,1,2 W. VAN DYKE DIXON,3 MARK HURWITZ,3 KWANG-SUN RYU,2 KWANG-IL SEON,4 AND JERRY EDELSTEIN3 Received 2000 February 21 ; accepted 2000 August 7

ABSTRACT We present a study of molecular-hydrogen absorption lines in the continuum spectra of three earlytype stars in the Carina Nebula, HD 93129a, HD 93250, and HD 303308, made with the Berkeley Extreme and Far-Ultraviolet Spectrometer (BEFS) on the ORFEUS telescope in 1993 September. The instrument possesses an intermediate spectral resolving power (j/*j \ 3000) over the wavelength range jj390È1170.We derive column densities for the H rotational levels up to [email protected]@ \ 6 and estimate the kinetic 2 temperature and the UV radiation Ðelds at the surface of the clouds in the nebula. All three stars show strong H absorption features, indicating that the Carina Nebula contains abundant molecular hydro2 gen. The UV radiation Ðeld strength seems to be correlated with the angular distance of the target star from g Car, implying that the clouds might be located close to g Car. We derive the absolute distances between the clouds and g Car, which are in good agreement with the angular positions of these stars. We also examine the morphology of the Carina Nebula based on the present result, as well as previous CO observations, and derive the conversion factor N(H )/W (12CO) in this region. 2 Subject headings : H II regions È ISM : individual (Carina Nebula) È ISM : molecules È ultraviolet : ISM 1.

INTRODUCTION

Interstellar clouds in the Carina Nebula have been extensively observed in the radio band of CO emission lines. De Graauw et al. (1981) showed that an extended molecular cloud appears to be wrapped partly around the brightest portion of the H II region (Car I and Car II) : part of the observed CO is located in front of the nebula, and part behind it. They argued that the situation in Carina is consistent with the sequential star formation model proposed by Elmegreen & Lada (1977), and the indication of a new heating center northwest of Car I might suggest the next step of star formation in the sequence. Whiteoak & Otrupcek (1984) concluded from their observations that CO is concentrated in the dust regions around the H II distribution rather than across the H II : similar locations were observed for H CO (Gardner, Dickel, & Whiteoak 1973) and OH clouds 2(Dickel & Wall 1974). The main CO emission has a velocity around [22 km s~1, and it is argued that the result is consistent with an H II region expanding into nearby interstellar molecular clouds, forming sharp H II region/interstellar-dust interfaces during the process. More detailed observation revealed many small CO clumps in the nebula (Brooks et al. 1998), toward one of which very faint reddened stars were detected (Megeath et al. 1996). It was suggested by Megeath et al. (1996) that this could be the Ðrst evidence of ongoing star formation in the Carina Nebula. The high-resolution map also shows that the Keyhole region and the associated Car II radio continuum source appear to have little interaction with the main cloud emission. This may support the idea of the surrounding giant molecular cloud being dispersed by the stellar winds and ionizing Ñuxes of the massive stars in the region (Cox & Bronfman 1995). Radio continuum and infrared observations have also been made to investigate the ionized gas and dust distribution in the Carina Nebula. IRAS images of the nebula at 60 km show a close correspondence with the 0.843 GHz radio in all parts of the nebula, except to the west of Car I, where bright infrared emission arises from a radio-dark channel (Whiteoak 1994). Whiteoak suggested that the dust/

The Carina Nebula (NGC 3372, RCW53) is one of the prominent features of the southern Milky Way, extending over an area of about 4 square degrees. Visually, it consists of four lobes separated by dark lanes. The brightest of these lobes has a triangular shape and coincides with a radio emission peak (de Graauw et al. 1981). The radio peak itself has a complex structure, with two prominent components known as Car I and Car II. They are believed to be excited by the open clusters Trumpler 14 (Tr 14) and Trumpler 16 (Tr 16), respectively, located about 2.9 and 2.3 kpc away from our Sun (Davidson & Humphreys 1997). As these clusters host several of the rare spectral type O3 Wolf-Rayet stars, and the peculiar g Car is a member of Tr 16, the Carina H II region/molecular cloud complex is an excellent region for studying the interaction of massive stars with their parental giant molecular cloud (Brooks, Whiteoak, & Storey 1998). For example, the ionizing photons and energetic winds from these bright young star clusters might penetrate deeply into the nearby molecular clouds, altering their structures through subsequent heating and shocks and triggering star formation in molecular clumps located at the boundary of the H II region and the main molecular cloud (Megeath et al. 1996). Hence, investigation of the kinematics, excitation, and chemical composition of the interstellar clouds in the Carina Nebula will provide valuable information regarding these physical processes (Walborn et al. 1998). 1 Department of Physics, Korea Advanced Institute of Science and Technology, 373-1 Gusong-dong, Yusong-gu, Taejon 305-701, Korea ; dhlee=space.kaist.ac.kr, kwmin=space.kaist.ac.kr. 2 Satellite Technology Research Center, Korea Advanced Institute of Science and Technology, 373-1 Gusong-dong, Yusong-gu, Taejon 305-701, Korea ; ksryu=satrec.kaist.ac.kr. 3 Space Sciences Laboratory, University of California, Berkeley, California 94720-7450, USA ; vand=ssl.berkeley.edu, markh= ssl.berkeley.edu, jerrye=ssl.berkeley.edu. 4 Korea Astronomy Observatory, 61-1 Hwaam-dong, Yusong-gu, Taejon 305-348, Korea ; kiseon=hanul.issa.re.kr.

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molecular cloud lying in the dark lanes has a low density and is intermixed with the ionized gas to the southeast of Car II, but increases signiÐcantly in density near Car I, where its eastern face lies in contact with the ionized material, before wrapping around the back of the nebula. In optical bands, high-resolution Ca II and Na I interstellar absorption-line proÐles have been observed (Walborn 1982). The results may imply at least two kinds of motion in the nebula : a systematic expansion of the entire nebula at 15È20 km s~1 and smaller-scale ““ bubbles ÏÏ or ““ bullets ÏÏ with much higher velocities. Hydrogen is the most abundant element in interstellar clouds, and thus in its molecular form may contribute signiÐcantly to the evolution of chemical and physical states of the interstellar clouds (Spitzer 1978). Systematic measures of the rich spectrum of H in the far ultraviolet give the 2 values of N(J), the column density of H molecules in the 2 rotational level J of the ground vibrational and electronic state. The sum of these gives N(H ), which may be compared with predictions of theories 2for H formation and 2 disruption. The rotational excitation depends on physical conditions in the clouds and can be used to determine the gas temperature T , the volume density n(H I) of neutral H, and b , the probability per unit time that an H molecule at 0 surface absorbs a photon in any of the 2 Lyman or the cloud Werner lines (Spitzer & Jenkins 1975). Evidently, b is a 0 measure of the ultraviolet radiation Ñux incident on a cloud. The hydrogen in an interstellar cloud exposed to the UV radiation Ðeld is primarily atomic down to a certain depth, while the extinction (by dust and self-shielding) of the radiation Ðeld is large enough that photodissociation is unimportant further inside and molecular hydrogen forms. Molecules other than H form deeper into a cloud, and 2 some, such as CO, are capable of self-shielding when their lines become opaque to those photons in the narrow energy ranges capable of dissociating them (Reach, Koo, & Heiles 1994). In this paper, we present the results of observations of molecular hydrogen toward three early-type stars in the Carina Nebula, HD 93129a, HD 93250, and HD 303308, made with the Berkeley Extreme and Far-Ultraviolet Spectrometer (BEFS) on the ORFEUS telescope. We derive the kinetic temperatures and the total H column densities of the molecular clouds in the nebula 2and estimate the density and size of the clouds as well as the radiation Ðelds at their surfaces. We also obtain the conversion factor between N(H ) and W (12CO), the integrated antenna tem2 perature of 12CO, of the clouds using published CO emission data. 2.

OBSERVATIONS AND DATA REDUCTION

The Berkeley Extreme and Far-Ultraviolet Spectrometer (BEFS), located at the prime focus of the 1 m ORFEUS Telescope, Ñew aboard the space platform ASTRO-SPAS during the 1993 September mission of the space shuttle Discovery (Hurwitz & Bowyer 1995). With an e†ective area of about 4 cm2 and a resolution j/*j \ 3000 (D 100 km s~1 velocity resolution) over the wavelength range 390È 1170 AŽ , the BEFS is ideally suited for absorption-line studies of bright, far-UV sources. We have analyzed the spectra of three early-type stars in the Carina Nebula. Table 1 lists these target stars and their characteristics. General procedures for data extraction, background subtraction, and wavelength and Ñux calibration are described in Dixon

Vol. 545 TABLE 1 STAR CHARACTERISTICS Star

HD 93129a

HD 93250

HD 303308

l ............................ b ........................... Spectral Type . . . . . . . . . . . . . E(B[V ) . . . . . . . . . . . . . . . . . . . log N(H I) (cm~2) . . . . . . . . d (kpc)b . . . . . . . . . . . . . . . . . . . . v sin i (km s~1) . . . . . . . . . . . Integration Time (s) . . . . . . S/N . . . . . . . . . . . . . . . . . . . . . . . .

287.4 [0.6 O3 If* 0.54 21.4a 2.9 120 3045 25.4

287.5 [0.5 O3.0 V ((f)) 0.47 21.39 2.3 107 1426 15.7

287.6 [0.6 O3.0 V ((f)) 0.45 21.45 2.3 111 1771 13.7

NOTE.ÈBasic stellar parameters are from Walborn (1995). a Log N(H I) estimated from E(B[V ) using the relation of Diplas & Savage 1994. b Distance of Tr 16, which includes HD 93250, HD 303308, and g Car, is 2.3 ^ 0.2 kpc. Tr 14, which includes HD 93129a, is about 30% farther away (Davidson & Humphreys 1997).

et al. (2000, in preparation). For the present analysis, the data have been binned by 9 pixels, corresponding to 0.13 AŽ . Statistical errors are calculated from the signal and background arrays. Table 1 includes the total ORFEUS integration time and the statistical signal-to-noise ratio (S/N) in a 0.13 AŽ bin averaged over the 1045È1060 AŽ band. We Ðt the interstellar absorption features using the SPECFIT (Kriss 1994) program of IRAF to determine the column density of molecular hydrogen along the line of sight toward each star. Fitting is done via s2 minimization using the Simplex algorithm. To provide the stellar continuum for HD 93250 and HD 303308, we make use of reference spectra from the Copernicus Spectral Atlas (Snow & Jenkins 1977). Atlas counterparts, listed in Table 2, are selected according to their spectral similarity to the observed stars. Spectral types of these atlas counterparts may not quite match those of our program stars. We scale and relocate the continuum references and let them Ðt freely to reduce the possible errors caused by the choice of the continuum proÐles. As the total H column density N(H ) 2 the uncertainty in the 2 in each reference spectrum is less than column densities for the levels [email protected]@ \ 0 and 1 of the corresponding program star, we do not attempt to make corrections to account for H absorption in the reference 2 an O3 If* star, we assume a spectrum. For HD 93129a, linear continuum, since the photospheric lines are believed to be quite weak (Taresch et al. 1997). To determine the number of absorption components along each line of sight, we use high-resolution (FWHM \ 0.3È1.2 km s~1) Na I optical data, as Na I is considered to be a tracer of hydrogen atoms and molecules in the di†use ISM (Ferlet & Vidal-Madjar 1985). We use Ca II data when Na I data are not available. Table 3 lists the TABLE 2 CONTINUUM REFERENCES Star

HD 93129a

HD 93250

HD 303308

Atlass star . . . . . . . . . . . . . . None 15 Mon k Col Spectral type . . . . . . . . . . . ... O7 V ((f)) O9.5 V E(B[V ) . . . . . . . . . . . . . . . . ... 0.07 0.02 log N(H ) (cm~2) . . . . . . ... 15.55 15.51 2 NOTE.ÈAll of the atlas star parameters are from Snow & Jenkins 1977.

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MOLECULAR HYDROGEN IN THE CARINA NEBULA

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TABLE 3 ABSORPTION COMPONENTS 1 AND 2 FOR EACH STAR Star

HD 93129a

HD 93250

HD 303308

Components . . . . . . 1 2 1 2 1 2 Species . . . . . . . . . . . . Ca II Ca II Na I Na I Na I Na I v (km s~1)a . . . . . . . [27 4 [10 0 [33 0 i W (mAŽ )b . . . . . . . . . . 1 1.35 1 1 1 1.55 i NOTE.ÈThe parameters are from Walborn 1982 and Walborn & Hesser 1975. a Heliocentric radial velocity. b Relative equivalent widths of each velocity component.

absorption components toward our program stars with their heliocentric radial velocities and equivalent widths. For the present analysis, we select the 1045È1060 AŽ region, in which most of the v \ 0 ] 4 vibrational band in the Lyman series lie. Here the H rotational absorption lines are well separated, and stellar 2lines are not prominent. Two signiÐcant atomic interstellar lines, Ar I at 1048.22 AŽ and Fe II at 1055.26 AŽ , are included in our model. We Ðt each H rotational level up to [email protected]@ \ 6 independent2 ly. Given the continuum proÐle, Doppler b value, and the information about the absorption components, we make a database of SPECFIT input parameters for each observed spectrum, with the continuum placement and the scaling of column densities as free parameters for the Ðt. The whole model is convolved with a Gaussian of 0.31 AŽ FWHM, the width of the instrument point-spread function (Hurwitz et al. 1998), before the program computes the s2 value. After locating and scaling the stellar continuum to a proper level, we begin Ðtting the low values of J, then move to higher values, since high-J lines often lie on the wings of lower-J features. We repeat this manual procedure until we Ðnd a reasonable Ðt. Once the absorption lines and corresponding column densities are determined in this way, we run the computer program to Ðne tune the model by minimizing s2. Figure 1 shows the results for each spectrum : the observed spectrum, continuum proÐle, and the best-Ðtting model. 3.

ANALYSIS

In their optical studies of interstellar Ca II lines toward stars in the Carina Nebula, Walborn & Hesser (1975) concluded that there is a strong low-velocity component between the Sun and the Carina Nebula. Recently, Walborn et al. (1998) also noted the existence of a cool (T \ 100 K), low-velocity gas near the Sun, in the direction of the Carina Nebula, similar to the nearby clouds in the directions of f Oph and m Per. Using the Copernicus telescope, Savage et al. (1977) observed HD 92740, located at the outer boundary of the Carina Nebula, and derived the H column 2 density and the kinetic temperature T . The results were 01 N(H ) \ 9.33 ] 1019 cm~2 and T \ 83 K, which is in 2 with the assertion of Walborn 01 et al. (1998) for the accord existence of a cool, low-velocity gas near the Sun. Sembach, Danks, & Savage (1993) also found a strong low-velocity component in the optical Na I observations toward HD 94493, which is located about 1¡.6 away from the Carina Nebula. Dixon et al. (2000) used the results of Sembach et al. to obtain the H column densities of the components toward HD 94493, 2observed by the ORFEUS telescope. The total column density and the kinetic temperature of the low-velocity component toward HD 94493 are 5.01 ] 1019

FIG. 1.ÈThe v \ 0 ] 4 vibrational band in the Lyman series of molecular hydrogen from the ORFEUS spectra of (top) HD 93129a, (middle) HD 93250, and (bottom) HD 303308. Overlaid are continuum proÐles and the best-Ðtting models (dashed line).

cm~2 and 83 K, respectively (Dixon et al. 2000). It is interesting to note that the results of HD 94493 are very similar to those obtained by Savage et al. (1977) for HD 92740. As there is only one component observed in the highresolution optical data of HD 92740 (Walborn & Hesser 1975), we believe the low-velocity clouds toward HD 92740

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and HD94493 are probably the same foreground cloud spread over the line of sight toward the Carina Nebula. The results of high-resolution optical observations of our program stars, shown in Table 3, indicate that each star has two absorption components. Of these two, we believe the common components near zero velocity represent the foreground cloud. The components ranging from [10 to [33 km s~1 may represent the systematic expansion of the H II region/molecular clouds in the Carina Nebula. While the high-resolution optical data show clear evidence of the existence of the two components, the ORFEUS spectra, with D100 km s~1 velocity resolution, cannot separate them. Hence, we must estimate the H column density of the fore2 star before we obtain a ground cloud toward each target reasonable H column density of the corresponding molec2 nebula. ular cloud in the We make two basic assumptions to estimate the H 2 column density of the foreground cloud. First, we utilize the result of high-resolution optical data and adopt the Na I (or Ca II) equivalent width ratio of the foreground cloud to each Carina Nebula cloud to be the H column density ratio. Hence, the total H column density2toward each target star is divided between2 the foreground cloud and the Carina Nebula cloud according to the Na I (or Ca II) equivalentwidth ratio. Second, we assume that the kinetic temperatures and the radiation environments of the foreground cloud are the same toward each of the three target stars. More speciÐcally, we adopt the ORFEUS result of HD 94493 by Dixon et al. (2000) and Ðx the relative ratios of the J-level column densities of H for the foreground cloud. The J-level column densities of 2the Carina Nebula cloud are then Ðtted by SPECFIT while the total H column density 2 to the result of of the foreground cloud is scaled according the optical observation. We use the Doppler velocity b \ 4 km s~1 for the foreground cloud as observed in the optical spectrum of HD 94493 (Sembach et al. 1993). No Doppler velocity measurements toward our program stars are available, so we assume b \ 5 km s~1 for the Carina Nebula clouds, as most of the available interstellar Na I data show a b value in the range of 2È8 km s~1. Table 4 shows the results of our analysis : molecularhydrogen column densities of the Carina Nebula clouds for the various J levels and the total column density for the foreground cloud. Molecular hydrogen is abundant in the Carina Nebula and the foreground cloud. The total column

Vol. 545

density of the foreground cloud ranges from 8.32 ] 1019 cm~2 to 9.77 ] 1019 cm~2, which is in good agreement with the results of HD 92740 and HD 94493 (Savage et al. 1977 ; Dixon et al. 2000). SPECFIT calculates only statistical uncertainties, which are relatively small when the Ðtting function is well provided. It is systematic errors that dominate the uncertainty in the present analysis. Line blending, inÑuences from other absorption lines, and uncertainties in the adopted stellar continuum and cloud model may constitute such systematic errors. Though systematic errors are difficult to quantify, we attempt to assess the most signiÐcant uncertainties in our analysis, those arising from our assumed cloud model. First, we estimate the e†ect of the errors in N(J) of the foreground cloud model determined by Dixon et al. (2000). Since the column densities for low-order J levels of the foreground cloud are about equal to those of the clouds in the Carina Nebula, it is reasonable to assume that the uncertainties in our results directly correspond to the uncertainties in the foreground cloud model. Dixon et al. (2000) estimate an overall uncertainty in their quoted column densities of 0.3 dex for the [email protected]@ \ 0È3 lines, which are applicable to the results of HD 94493, and we take those errors for our systemtatic uncertainties due to the foreground cloud model. In the case of high-order J levels, the e†ect of the foreground cloud model is not signiÐcant because the column densities of the high-order J levels of the foreground cloud are insigniÐcant in comparison to the column densities of the clouds in the Carina Nebula. Next, we determine the sensitivity of our result to the assumed Doppler velocity of the clouds in the Carina Nebula. This second source of systematic error is probably more serious than the Ðrst for the high-order J levels since there are no data available regarding the Doppler velocity of the Carina Nebula clouds, and the column densities for J º 4 are highly sensitive to the assumed Doppler parameter as they are located on or near the saturation portion of the curve of growth. Our test shows the column densities of J ¹ 1 lines vary less than 0.15 dex, while those of J º 2 decrease up to 0.4 dex when b is raised to 8 km s~1. The strongest UV absorption lines of H in the spectrum 2 R(1), and P(1) arise from the [email protected]@ \ 0 and 1 levels : R(0), Lyman lines. The relative populations of [email protected]@ \ 0 and 1 are established primarily by collisions with thermal protons, since line self-shielding reduces the e†ects of UV pumping

TABLE 4 MOLECULAR HYDROGEN COLUMN DENSITIES OF THE FOREGROUND CLOUD AND THE CLOUDS IN THE CARINA NEBULA Star Component . . . . . . N(H )c . . . . . . . . . . . 2 N(0)c . . . . . . . . . . . . . N(1)c . . . . . . . . . . . . . N(2)c . . . . . . . . . . . . . N(3)c . . . . . . . . . . . . . N(4)c . . . . . . . . . . . . . N(5)c . . . . . . . . . . . . . N(6)c . . . . . . . . . . . . .

HD 93129a 1a 19.80 18.87 19.74 16.98 17.76 16.64 17.73 13.81

2b 19.93 19.59 19.66 16.79 15.44 14.52 14.33 ...

HD 93250 1a 19.94 19.31 19.79 17.70 18.00 17.75 18.38 14.77

2b 19.92 19.58 19.65 16.78 15.43 14.51 14.32 ...

NOTE.ÈAll column densities are given as logarithms (cm~2). a The molecular clouds in the Carina Nebula. b The foreground cloud. c See text for the uncertainties of these values.

HD 303308 1a 19.81 18.89 19.67 18.00 18.09 18.37 18.71 17.77

2b 19.99 19.65 19.72 16.85 15.49 14.58 14.39 ...

No. 2, 2000

MOLECULAR HYDROGEN IN THE CARINA NEBULA TABLE 5

DERIVED PHYSICAL PARAMETERS OF THE CLOUDS IN THE CARINA NEBULA Star

HD 93129a

HD 93250

HD 303308

Component 1 1 1 f 0.137 0.122 0.108 R ] n (10~16 s~1) 2.8 ] 103 1.4 ] 104 1.0 ] 105 b (10~10 s~1) 3.5 ] 103 2.3 ] 104 1.6 ] 105 0 NOTE.ÈPossible error ranges are discussed in the text.

dramatically in saturated lines. For clouds with strong selfshielding [N(0) and N(1) º 1017 cm~2], Dalgarno, Black, & Weisheit (1973) argue that collisions of molecules with protons should outweigh other processes in determining the relative population of the [email protected]@ \ 0 and 1 levels, meaning that the [email protected]@ \ 0 to 1 population ratio is a measure of the proton kinetic temperature, and hence, the cloud kinetic temperature. Therefore, the ratio of column densities is given by

A B

A

B

[E [170 K N(1) g 01 \ 9 exp , (1) \ 1 exp kT T N(0) g 01 01 0 where g and g are the statistical weights of [email protected]@ \ 0 and 1 0 & Beckwith 1 levels (Shull 1982).

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We apply equation (1) to our present analysis of the Carina Nebula and obtain the kinetic temperatures T of 01 the molecular clouds. The results are 827, 158, and 416 K for HD 93129a, HD 93250, and HD 303308, respectively. Although these values are highly uncertain because T 01 varies logarithmically with the change in N(1)/N(0), of which each term contains a certain amount of systematic errors, it appears that the molecular clouds in the Carina Nebula are much heated by their active environments compared to the average temperature of interstellar molecular clouds, which is 77 ^ 17 K (Savage et al. 1977). Table 5 shows the fraction of molecular hydrogen f, the product of the volume density n , and the H formation H 2 rate R, and b , the photon-absorption rate in the Lyman 0 and Werner bands at the exterior of the clouds. We estimate these parameters from the measured rotational excitations of H and N(H I) observations. According to Jura (1975), 2 the levels [email protected]@ \ 4 and [email protected]@ \ 5 are populated both by direct formation into these levels of newly created molecules and by photon pumping. For densities less than 104 cm~3, the levels [email protected]@ \ 4 and [email protected]@ \ 5 are depopulated by spontaneous emission. Assuming the steady state condition, we derive Rn and b from N(4)/N(H), following Jura (1974). We0expect uncertainties in these values since the calculation is based on the column densities that contain system-

FIG. 2.ÈContour representations of the 12CO(1È0) emission integrated over a velocity range of [30 to 0 km s~1 (thin lines), and of the 60 km FIR emission (thick lines), superimposed on a red wavelength optical image. The Ðgure is taken from Brooks et al. (1998) and Whiteoak (1994). g Car and three stars observed by ORFEUS are indicated : 1 \ g Car, 2 \ HD 303308,3 \ HD 93250, and 4 \ HD 93129a.

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atic errors. However, since the derived values Rn and b are 0 linearly proportional to N(4)/N(H), the estimated error in the column density is about a factor of 2.5 at its maximum when b \ 8 km s~1 for the high JA values and smaller for the [email protected]@ \ 0,1 lines. We Ðnd that Rn ranges from 2.8 ] 10~13 s~1 to 1.0 ] 10~11 s~1. These values are extremely high compared to that of the average interstellar cloud from the Copernicus observations (Jura 1975). We obtain b \ 1.6 ] 10~5 for the cloud toward HD 303308, which is0 about 104.5 times larger than that in the solar neighborhood b \ 5 ] 10~10 s~1 (Jura 1974). As HD 303308 is located _ close to g Car on the celestial sphere, the cloud in the direction toward HD 303308 is presumably irradiated strongly by g Car. The values of b for the clouds toward HD 93250 0 and HD 93129a are 2.3 ] 10~6 s~1 and 3.5 ] 10~7 s~1, respectively. Accounting for the systematic errors in the column densities, the uncertainties in these values are also expected to be about a factor of 2.5. Nevertheless, it is interesting to note that, for each of our target stars, b decreases with increasing distance from g Car. Assuming0 that the clouds toward HD 303308, HD 93250, and HD 93129a are irradiated by g Car only and that the incident Ñux falls as the inverse square of the distance (e.g., no intervening dust extinction), we Ðnd that the relative distance of each cloud from g Car is approximately 1, 2.6, and 6.8, respectively. Our approximate relative distances derived from the H excitation conditions are in rea2 relative separation between g sonable agreement with the Car and the sight line to each program star on the celestial sphere. Because of three-dimensionalprojection e†ects and the presence of dust in the region, one would not expect a perfect correlation. From our measurement of the Ñux on each cloud and the intrinsic brightness of g Car, we can add an absolute scale to our relative distances. The strength of the UV continuum radiation is usually parametrized by its ratio to the average interstellar Ñux near 1000 AŽ , F \ 105 photons cm~2 s~1 _ H absorption rate in the AŽ ~1 (Jura 1974). The unshielded Lyman and Werner bands for a UV 2Ñux F may be written 0 b \ b (F /F ) , (2) 0 _ 0 _ where b \ 5 ] 10~10 s~1 (Jura 1974). With these relations, it is_ possible to estimate the photon Ñux at the boundaries of the three molecular clouds and, thus, the actual distances from g Car, if the luminosity in the Lyman bands of g Car is known. Davidson et al. (1995) estimate roughly that the apparent Ñux of g Car around 1000 AŽ would probably be between 4 ] 10~9 and 4 ] 10~8 ergs cm~2 s~1 AŽ ~1 if there were no circumstellar and interstellar extinction. This Ñux corresponds to F \ 200È2000 photons cm~2 s~1 AŽ ~1. From equation (2), we derive the radiation Ñux F near 1000 AŽ at the boundary of the cloud toward HD0 303308 to be 3.2 ] 109 photons cm~2 s~1 AŽ ~1. As the radiation Ñux is inversely proportional to the square of the distance when interstellar extinction is assumed to be negligible, a choice of the distance from the earth to g Car of D \ 2.3 kpc yields a distance between g Car and the cloud toward HD 303308 of D \ 0.58È1.82 pc. If we take D \ 1.2 pc for convenience, 0the distances from g Car to 0the clouds toward HD 93250 and HD 93129a are 3.1 pc and 8.2 pc, respectively. These absolute distances are in good agreement with the absolute angular position of our sight lines. For example, the absolute distance for cloud toward HD 93129a would correspond to D [email protected] at the distance of g Car, in excellent agreement with the map shown in Figure 2.

Vol. 545 4.

DISCUSSION

CO emission is believed to be thermalized in both lowand high-density gas and is therefore regarded as a suitable tracer for the overall distribution and velocity structure of the molecular cloud. Brooks et al. (1998) observed the 12CO(1È0) emission at 115.271 GHz in the Carina Complex. The result shows that CO emission extends over 3¡ in the northwest-southeast direction, in agreement with the Carina GMC identiÐed in the Columbia survey (Grabelsky et al. 1988). Their higher-resolution map reveals a clumpy morphology with two strong emission areas : one centered on Car II and the other toward a region northwest of the optical emission. Figure 2 is a contour representation (thin lines) of the 12CO(1È0) emission integrated over a velocity range of [30 to 0 km s~1, superimposed on the optical image of the area centered on the Keyhole Nebula (Brooks et al. 1998). Also shown on the map is the 60 km IR emission (thick lines) observed by IRAS (Whiteoak 1994). While the contour map shows a smaller amount of CO for the cloud toward HD 303308 than those toward the two other stars, our observation of H reveals that all three clouds contain more or less equal2 amounts of molecular hydrogen. This situation is similar to the argument of Magnani et al. (1998) in their analysis of the CO-to-H conversion factor, N(H )/W (CO), in two translucent high-2 2 where W (CO) is the integrated latitude molecular clouds, antenna temperature of CO molecules. The H column 2 while density is approximately constant across the clouds, the conversion factor, N(H )/W (CO), is primarily Ñuctuated by CO. The reason might2 be that CO molecules for the cloud toward HD 303308, located in the central Keyhole region near g Car, are dissociated more easily by the strong radiation environment. In fact, CO and its principal isotopic varieties are known to be more easily photodissociated than H . According to van Dishoeck & Black 2 (1988), the photodissociation rate of 12CO is three to four orders of magnitude higher than that of H inside the 2 clouds, and for clouds with N(H ) \ 1021 cm~2, 2 N(CO)/N(H ) decreases by a factor of 35 from its original value if the 2UV intensity increases 10 times. Our result is also consistent with the 60 km FIR data which show a high intensity near HD 303308. As Whiteoak (1994) mentioned, warm dust and H II are closely correlated over the entire nebula (except in the radio-dark channel to the west of Car I) based on the FIR data and 0.843 GHz radio continuum. In this regard, it is of interest to note the results of Cox & Bronfman (1995), who observed the molecular gas of the Keyhole Nebula in 12CO(2È1) emission and found that the molecular clumps detected around g Car are characterized by masses of about 10 M , 2 orders of magnitude lower _ masses. They argued that this than the corresponding virial large discrepancy may be due to the fact that the clouds are supported by the external pressure from warm and di†use gas that is ionized by the strong UV radiation Ðeld permeating this region, and the molecular clumps associated with the Keyhole Nebula are part of a much larger expanding structure with a velocity of 20 km s~1. We estimate directly the conversion factor N(H )/W (CO) in Table 6, using our H total column densities2 and CO 2 integrated antenna temperature from Brooks et al. (1998). Until recently, this important conversion factor has been obtained only by indirect methods, which yield values ranging from 0.3È7 ] 1020 cm~2 (K km s~1)~1. Analysis of the halo cloud toward HD 94473, also observed with

No. 2, 2000

MOLECULAR HYDROGEN IN THE CARINA NEBULA

891

TABLE 6 CO TO H CONVERSION FACTOR IN THE CARINA NEBULA 2 Star

HD 93129a

HD 93250

HD 303308

Component 1 1 1 N(H )a 19.80 19.94 19.81 2 W (CO) (K km s~1)b 32È44 32È44 \20 N(H )/W (CO) (1018 cm~2/K km s~1) 1.4È2.0 1.9È2.7 [3.2 2 a Column densities are given as logarithms (cm~2). From Table 4. b Integrated antenna temperature of 12CO with e†ective FWHM beam of [email protected] (Brooks et al. 1998).

ORFEUS, gives the direct estimation of (4 ^ 2) ] 1019 cm~2 (K km s~1)~1 (Ryu et al. 2000). Our results for the Carina Nebula clouds in Table 6 are smaller than that of the halo cloud toward HD 94473 by an order of magnitude. We can think of two reasons for these extremely low conversion factors. First, we might have attributed too much N(H ) to the foreground cloud. However, even if we assume that 2all of the N(H ) observed by BEFS belongs to the 2 Carina Nebula, the result is still surprisingly low. Second, it is possible that the target stars are embedded within, not beyond, the nebula. In an embedded geometry, molecular gas more distant than the star could contribute to the CO emission but not to the H absorption. Although our 2 N(H )/W (CO) ratios are probably dominated by the 2 amount of background CO, it is interesting to note that in the cloud toward HD 303308, where the CO photodissociation e†ect discussed above would be expected to be of greatest importance, our observed ratio is at its maximum. 5.

SUMMARY

We have observed H molecules along lines of sight 2 toward three O3 type stars in the Carina Nebula, HD 303308, HD 93250, and HD 93129a, with the Berkeley spectrometer on the ORFEUS telescope. As optical spectra of all three show the existence of a foreground cloud, we

assume the physical parameters of the foreground cloud using the results of Dixon et al. (2000) and Savage et al. (1977). We Ðnd H abundant in all three clouds in the Carina 2 Nebula. Estimation of the kinetic temperature using the column densities of the rotational levels [email protected]@ \ 0 and 1 indicates that the clouds are embedded in a hot environment. Column densities of higher rotational levels give UV radiation densities at the surface of the clouds. The cloud toward HD 303308, closest to g Car among our program stars, shows the strongest UV photon density. Assuming a simple inverse square law and the photon Ñux of g Car near 1000 AŽ given by Davidson et al. (1995), we estimate the distances from g Car to be 1.2, 3.1, and 8.2 pc for the clouds toward HD 303308, HD 93250, and HD 93129a, respectively. These distances are in good agreement with the angular positions of our program stars shown in Figure 2, which implies our estimation from UV excitation data is reasonable. Using the 12CO(1È0) survey by Brooks et al. (1998) and taking account of the background CO that may exist, we set the lower limit for conversion factor of N(H )/W (CO) \ 2 ] 1018 cm~2/K km s~1. 2 We thank K. J. Brooks and J. B. Whiteoak for their permission to use the Ðgures of the Carina Nebula in Figure 2.

REFERENCES Brooks, K. J., Whiteoak, J. B., & Storey, J. W. V. 1998, Proc. Astron. Soc. Magnani, L., et al. 1998, ApJ, 504, 290 Austrailia, 15, 202 Megeath, S. T., Cox, P., Bronfman, L., & Roelfsema, P. R. 1996, A&A, 305, Cox, P., & Bronfman, L. 1995, A&A, 299, 583 296 Dalgarno, A., Black, J. H., & Weisheit, J. C. 1973, ApJ, 14, 77 Reach, W. T., Koo, B.-C., & Heiles, C. 1994, ApJ, 429, 672 Davidson, K., et al. 1995, AJ, 109, 1784 Ryu, K. S., Dixon, W. V., Hurwitz, M., Seon, K. I., Min, K. W., & Edelstein, Davidson, K., & Humphreys, R. M. 1997, ARA&A, 35, 1 J. 2000, ApJ, 529, 251 de Graauw, T., Lidholm, S., Fitton, B., Beckman, J., Israel, F. P., NieuwenSavage, B. D., Bohlin, R. C., Drake, J. F., & Budich, W. 1977, ApJ, 216, 291 huijzen, H., & Vermue, J. 1981, A&A, 102, 257 Sembach, K. R., Danks, A. C., & Savage, B. D. 1993, A&AS, 100, 107 Dickel, H. R., & Wall, J. V. 1974, A&A, 32, 5 Shull, J. M., & Beckwith, S. 1982, ARA&A, 20, 163 Diplas, A., & Savage, B. D. 1994, ApJ, 427, 274 Snow, T. P., & Jenkins, E. B. 1977, ApJS, 33, 269 Dixon, W. V., Hurwitz, M., & Lee, D. H. 2000, ApJ, submitted Spitzer, L. 1978, Physical Processes in the Interstellar Medium (New York : Elmegreen, B. G., & Lada, C. J. 1977, ApJ, 214, 7251 Wiley) Ferlet, R., & Vidal-Madjar, A. 1985, ApJ, 298, 838 Spitzer, L., & Jenkins, E. B. 1975, A&A, 13, 133 Gardner, F. F., Dickel, H. R., & Whiteoak, J. B. 1973, A&A, 23, 51 Taresch, G., et al. 1997, A&A, 321, 531 Grabelsky, D. A., Cohen, R. S., Bronfman, L., & Thaddeus, P. 1988, ApJ, van Dishoeck, E. F., & Black, J. H. 1988, ApJ, 334, 771 331, 181 Whiteoak, J. B. 1994, ApJ, 429, 225 Hurwitz, M., & Bowyer, S. 1995, ApJ, 446, 812 Whiteoak, J. B., & Otrupcek, R. E. 1984, Proc. Astron. Soc. Australia, 5, Hurwitz, M., et al. 1998, ApJ, 500, 1 552 Jura, M. 1974, ApJ, 191, 375 Walborn, N. R. 1982, ApJS, 48, 145 ÈÈÈ. 1975, ApJ, 197, 581 ÈÈÈ. 1995, Rev. Mexicana Astron. AstroÐs., 2, 51 Kriss, G. A. 1994, in ASP Conf. Ser. 61, Astronomical Data Analysis Walborn, N. R., et al. 1998, ApJ, 492, 169 Software and Systems III, ed. D. R. Crabtree, R. J. Hanisch, & J. Barnes Walborn, N. R. & Hesser, J. E. 1975, ApJ, 199, 535 (San Francisco : ASP), 437

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