Molecular cores of the high-latitude cloud MBM 40

New Astronomy 8 (2003) 795–803 www.elsevier.com / locate / newast

Molecular cores of the high-latitude cloud MBM 40 Young Chol Minh a , *, Hyun-Goo Kim a , Youngung Lee a , Hyeran Park b , Kwang-Tae Kim b , Yong-Sun Park c , Sang Joon Kim d a

b

Korea Astronomy Observatory, 61 -1, Hwaam, Yuseong, Daejeon 305 -348, South Korea Department of Astronomy and Space Science, Chungnam National University, Yuseong, Daejeon 305 -335, South Korea c SEES, Seoul National University, Sinlim, Gwanak 151 -742, South Korea d Department of Astronomy and Space Science, KyungHee University, Yong-In, Gyunggi-do 449 -701, South Korea Received 5 March 2003; received in revised form 13 May 2003; accepted 21 May 2003 Communicated by G. Setti

Abstract Towards the high-latitude cloud MBM 40, we identify 3 dense molecular cores of M | 0.2–0.5 M ( , and sizes of | 0.2 pc in diameter embedded in the H I cloud of | 8 M ( which is observed to be extended along the northeast–southwest direction. The molecular cloud is located almost perpendicularly to the H I emission. We confirm the previous result of Magnani et al. that MBM 40 is not a site for new star formations. We found a very poor correlation between the H I and the IRAS 100 mm emissions, but the CO (1–0) and 100 mm emissions show a better correlation of WCO /I100 5 160.2 K km s 21 (MJy sr 21 )21 . This ratio is larger by a factor of $ 5 than in dense dark clouds, which may indicate that the CO is less depleted in MBM 40 than in dense dark clouds.  2003 Elsevier B.V. All rights reserved. PACS: 95.85.Bh Keywords: ISM: abundances; ISM: individual: MBM 40; ISM: molecules

1. Introduction More than 100 of high-latitude clouds (HLCs) have been identified, which typically have low-visual extinctions A tot v | 1–2 mag and mass range of 10– 100 M ( . HLCs appear to be extended sometimes to

*Corresponding author. Tel.: 182-42-865-3263; fax: 182-42861-5610. E-mail addresses: [email protected] (Y.C. Chol Minh), [email protected] (H.-G. Kim), [email protected] (Y. Lee), [email protected] (H. Park), [email protected] (K.-T. Kim), [email protected] (Y.-S. Park), [email protected] (S. Joon Kim).

a few square degrees mainly because of their proximity to us (distances | 100 pc) (Low et al., 1984; Blitz et al., 1984). These HLCs span a wide range of physical properties, i.e., from diffuse to dense clouds, and belong to the translucent cloud category (van Dishoeck and Black, 1988). Although the HLCs are observed at the high galactic latitudes (ubu $ 208), their actual z-distances from the Galactic plane are mostly less than 100 pc because of their proximity to us. Therefore their environments should not be very different from those for dark clouds observed usually toward the galactic plane. The molecular abundances of HLCs seem to be similar to those of dense dark clouds but the

1384-1076 / 03 / $ – see front matter  2003 Elsevier B.V. All rights reserved. doi:10.1016 / S1384-1076(03)00068-X

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physical and chemical properties of HLCs are still controversial (Turner, 2000). The proximity to us and their relatively low masses make the HLCs an ideal object to study, for instance, the clumpy structures, the penetration of UV radiation, and the evolution of molecular clouds. The morphologies of most HLCs suggest that shock compressions are prevalent and the hydrodynamical turbulence seems to dominate the line broadening (e.g. Joncas et al., 1992; Minh et al., 1996). In many HLCs, there exist molecular cores which may have formed after shock compressions followed by the cooling and recombination phase. In general the diffuse H I gas (n H I | 30–40 cm 23 ) surrounds denser molecular cores (n H 2 $ 10 3 cm 23 ), but they do not seem to be in pressure equilibrium with each other. The pressure of the dense cores appears to be large enough to disrupt the whole cloud in the time scale of about 10 5 2 10 6 yrs (Magnani et al., 1985; Keto and Myers, 1986; Minh et al., 1996). Some of dense cores in HLCs are probably unstable to gravitational collapse, and there are a few observational evidences for recent star formations, such as T Tauri stars or outflows (Sandell et al., 1987; Sato and Fukui, 1989). Most of the HLC cores, however, appear to have much less gravitational energy than the internal turbulent energy (e.g. Magnani et al., 1985; Minh et al., 1996). Thus, star-formation is thought not to be an efficient process in most HLCs, but it is still necessary to investigate low-mass star-formations inside the molecular cores (Magnani et al., 1996). MBM 40 is a typical HLC located at a distance of | 100 pc in the position (l, b)5(378.6, 448.7) with the peak A v of 1–2 mag which is easily identifiable away from the background confusion in the IRAS 100 mm emission map. This cloud has been identified as a cirrus cloud in the IR-excess clouds list of Desert et al. (1988). While other cirrus clouds show widespread wispy and filamentary FIR emission features, it shows a well confined compact feature. Recently, Lee et al. (2002) have mapped the 12 CO and 13 CO 1–0 lines toward the molecular cores of MBM 40 and derive the total mass 17 M ( . The previous studies on MBM 40 have been summarized by Magnani et al. (1996) and they have reported no

evidence of recent or ongoing star formation, but suggested that the molecular cores are likely bounded gravitationally. In this paper we report observational results for MBM 40 with the H I 21 cm line and, using the CO data of Lee et al. (2002) and IRAS 100 and 60 mm data, discuss the physical properties of the molecular cores of MBM 40 as one of the typical HLCs. The observational methods are in Section 2, the cloud properties are derived in Section 3 and the discussions on the cloud stability and the comparison with the IRAS 100 mm data are given in Section 4.

2. Observations The H I 21 cm line was observed for 25 days with normal mode (complete survey) at June 1994 with the Synthesis Telescope of the Dominion Radio Astrophysical Observatory at Canada (Roger et al., 1973), which consists of 7 antennas of 9-m diameter (field of view52.68 for 20% response). The synthesized beam size is 19 3 2.79 (EW 3 NS). The 256 channel cross-correlation spectrometer was used for both left and right-circularly polarized emissions, which gives the channel resolution of 0.412 km s 21 . The map center is (a,d ) 2000 5 (16 h 10 m 14 s , 218549140). The system temperatures are about 65 K and the typical rms value (1 s ) of the spectra is 1.5 K. The 26-m single dish results were included to compensate the missing short spacing data. This prime focused antenna has a broad band receiver and a 256 channel spectrometer for each polarization identical to those for the Synthesis Telescope. Details of the telescopes and observing parameters can be found at http: / / www.drao.nrc.ca / facilities / telescopes. The 12 CO and 13 CO J51–0 lines were observed in 1997 January using the 14 m telescope of Taeduk Radio Astronomy Observatory at Daejeon, Korea. The HPBW and main beam efficiency of the antenna are 450 and 0.42, respectively, at 115 GHz. The data have been resampled with the 0.1 km s 21 resolution and their typical rms values (1 s ) are about 0.05 K. The observational method and parameters are summarized at Section II of Lee et al. (2002).

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Fig. 1. Sample spectra of the H I 21 cm line. On the top left of each panel the observed position is indicated as an offset from the map center (a,d ) 2000 5 (16 h 10 m 14 s ,218549140). Their line parameters are listed in Table 1.

peaks at vlsr | 3 km s 21 and 2 2 km s 21 . The 3 km s 21 emission is slightly more confined centrally compared to the 2 2 km s 21 emission, but they are difficult to be separated and seem to be associated together. The total velocity-integrated intensity map is shown in Fig. 3. Fig. 2 shows sample spectra of the 12 CO and 13 CO J51–0 lines and their Gaussian fit line parameters are included in Table 1. The CO observed region is indicated as a box in Fig. 3. There appears no

3. Cloud properties and discussion

3.1. Emission distributions Fig. 1 shows sample spectra of the observed 21 cm lines and their line parameters are included in Table 1. We show the spectra taken at the same offset positions with CO spectra in Fig. 2, but the H I emission is extended relatively uniformly throughout the observed area. In most places there exist two Table 1 Line parameters of the observed H I and CO spectra in Figs. 1 and 2 a Position offset b (Da, Dd )

Species

T (K)

FWHM (km s 21 )

vlsr (km s 21 )

H I 21 cm CO 1–0 13 CO 1–0

32.9 3.2 1.4

5.9 1.1 0.9

2.77 2.76 2.69

( 2 29, 29)

H I 21 cm 12 CO 1–0 13 CO 1–0

33.0 4.7 1.6

8.9 1.0 0.7

2.35 3.30 3.29

( 2 89, 129)

H I 21 cm 12 CO 1–0 13 CO 1–0

31.6 5.1 1.4

7.4 0.8 0.8

2.4 3.09 3.05

( 2 29, 209)

12

a b

CO data are Gaussian fit results and CO antenna temperatures are T R* . From the map center (a,d ) rm 2000 5 (16 h 10 m 14 s ,218549140).

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Fig. 2. Sample spectra of 12 CO and 13 CO J51–0 lines observed towards the cores of molecular cores A, B, and C in Fig. 4b. The position offsets of the cores are shown in each panel. Their line parameters are listed in Table 1.

specific structural relation between H I and CO emissions. The velocity-integrated intensity maps of the 12 CO and 13 CO (1–0) lines are shown in Fig. 4a

and 4b, respectively, and the IRAS 1 100 mm intensity in Fig. 4c. The molecular and IRAS emissions show similar distributions and appear as localized clumps elongated along the northwest–southeast direction which is roughly perpendicular to the H I emission distribution. The shape of the CO-bright region may suggest a shock compression as an origin of the molecular component. The three 13 CO cores are labelled from A to C and indicated in Fig. 4b. Their physical parameters will be discussed further in Section 3.3.

3.2. Temperature 12

CO has been considered a good thermometer in many molecular clouds because of the large opacity in its emitted lines, arising from its large abundance and small dipole moment. We derive the CO 1–0 excitation temperature using the 12 CO and 13 CO (1–0) lines assuming LTE conditions and the result is shown in Fig. 5. We found T ex | 10 K around the Fig. 3. Total integrated intensity map of the H I emission. The (0, 0) position is the map center in Table 1. The lowest contour level is 250 K km s 21 , and the increment is 15 K km s 21 . The box indicates the region observed with the CO lines in Fig. 4.

1

The IRAS Sky Survey Atlas (ISSA) image has been acquired from the Infrared Processing and Analysis Center (IPAC) using Skyview Virtual Observatory (http: / / skyview.gsfc.nasa.gov / skyview.html).

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Fig. 4. (a) Integrated intensity map of the 12 CO (1–0) line (eT *R dv). The contours increase by 1 K km s 21 from 2 K km s 21 . (b) eT *R 13 CO (1–0) dv map. The contours increase by 0.4 K km s 21 from 0.5 K km s 21 . The three molecular cores are indicated as A to C. (c) The IRAS 100 mm intensity map. The contours increase by 1 MJy sr 21 from 2 MJy sr 21 .

cores, which is the canonical value for cold dark clouds in the galactic plane and also for the dense cores of HLCs in general (Turner, 1993). The linewidth of CO lines Dvobs (CO)|1 km s 21 are also similar to those of the typical cold and quiescent dark clouds in the Galactic plane. We also derive the dust color temperature using the IRAS data (Fig. 6). Since there are no embedded sources and bright visible stars in the MBM 40 region (Magnani et al., 1996), we assume that the general interstellar radiation field heats the dust in the region and derive the dust color temperature

Fig. 5. The excitation temperature of the CO 1–0 line. The contours increase by 0.5 K from 6 K.

(T dust ) from the IRAS 100 and 60 mm intensities using the relation,

F

hc 1 1 T dust . ] ] 2 ]] k l60 l100

GY F S D G I100 l100 ln ] ]] I60 l60

31 b

if the simplified conditions are assumed, including a constant dust temperature along the line of sight and the dust absorption efficiency factor Q abs ~ l 2 b . The

Fig. 6. Dust color temperature distribution derived from the IRAS 60 mm and 100 mm data. The contour levels are indicated in the figure.

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value of b is thought to be in the range 1 # b # 2 depending on the dust materials. Assuming b 51.5 (see Kim, 1997), we derive about T dust | 22 K in the dense cores in Fig. 5, but no significant temperature structure is found. A small decrease in T dust toward the denser region is expected as mainly due to shielding of the ultraviolet photons of the interstellar radiation field. These temperatures should be regarded as an upper limit of the mean T dust . The small grains in the region can contribute to the 60 mm more effectively than to the 100 mm, which would result in the overestimation of the dust color temperature. Therefore, in MBM 40 we may expect that the cold dense cores are embedded in the warmer | 20 K envelope, which is also a typical value for dense warm clouds in the Galactic plane.

3.3. Gas abundance and LTE masses The H I emission in MBM 40 does not appear to be confused much by the foreground and background H I gas. The H I column densities may be estimated directly from the observed 21 cm lines, assuming optically thin emission, using the equation, N(H I)5 1.823 3 10 18 WH I cm 22 where WH I is the observed integrated intensity of the 21 cm line in the unit K km s 21 (e.g. de Vries et al., 1987). The H I column 20 22 densities are in the range of 1–5 3 10 cm in the whole observed region. The molecular component is mainly concentrated in clumps. The 13 CO column density, N( 13 CO)|13 10 15 cm 22 , was derived toward the peak positions of the clumps, assuming optically thin emission of the 13 CO 1–0 lines, LTE, and the rotational temperature of 10 K. The 50% change of the rotational temperature gives # 20% uncertainty in the total column density value and this estimation of the column density should be correct within a factor of 2–3 including all observational uncertainties. The parameters for the three 13 CO cores are summarized in Table 2. Using the conversion ratio H 2 / 13 CO56 3 10 5 (Dickman, 1978), we derive the H 2 column density, N(H 2 )¯6310 20 cm 22 at the peak position of the core C, and the mean volume density n(H 2 ) 5 4 3 10 2 cm 23 assuming the size of | 0.5 pc ( | 2 3 halfpower size). Although there should be a large

Table 2 Parameters of the molecular cores Cores

Position a (9, 9)

Size b (pc)

N( 13 CO)c (cm 22 )

MLTE d (M ( )

MVir (M ( )

A B C

( 2 2, 20) ( 2 2, 2) ( 2 8, 12)

0.16 0.23 0.25

6.8 3 10 14 6.8 3 10 14 7.7 3 10 14

0.16 0.34 0.46

13.6 11.8 16.8

a

Peak positions of the 13 CO cores in Fig. 4b. Offsets from the map center in Table 1. b Half-power size in diameter of the 13 CO 1–0 cores assuming a spherical geometry and the distance of 100 pc. c Beam-averaged column density of 13 CO at the peak position. Derived assuming optically thin emission and the rotational temperature of 10 K. d LTE mass of the clump. Derived assuming a Gaussian density distribution and H 2 / 13 CO56 3 10 5 (Dickman, 1978).

uncertainty in estimating the mean volume density, this value indicates that the molecular abundances are similar to the typical cold dark clouds in the galactic plane. The H I mass of MBM 40 has been derived to be MH I | 8 M ( by integrating the H I line intensities in the observed region. The total gas mass has been estimated to be 17 M ( for the observed region of the CO line using the conversion factor 2.3 3 10 20 (K km s 21 )21 (Lee et al., 2002). The LTE masses of the three cores in Fig. 4b are determined to be | 0.2–0.5 M( using M 5 Npeak (H 2 ) 3 2 3 mm H 3 Area(FWHM) where m 51.36 amu per H nuclei accounts for the mass of He and other elements. The Area(FWHM) of each clump has been decided from the half-power contour of the peak intensity in the 13 CO map. The parameters of the molecular cores are summarized in Table 2.

3.4. Virial mass of the cores In several translucent clouds dense cores similar to those of dark clouds have been found (see e.g. Turner et al., 1992; Reach et al., 1995; Magnani et al., 1996 and references therein), but the star formations in HLCs are still controversial. In discussing the cloud stability and future star-formations, it would be useful to test whether the clump is balanced in energy. In case one may ignore both of the magnetic field and external pressure, the turbulence can be considered as a main agent for supporting the clouds

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against the self-gravity. Then, the simplest scalar virial theorem stats: 3GM . KV L . ]] 5R 2 tur

(1)

Since we can observe only one component of the velocity dispersion along the line of sight, the turbulent velocity kV 2tur l 1 / 2 is related to the observed full line width at half maximum DVFWHM approximately as 3 GDV KV L . F]] 8 ln 2 2 tur

2 FWHM

.

(2)

And from this equation, we can rewrite the simplest virial theorem in terms of observable quantities only,

F

M pc DVFWHM . 0.069 km s 21 ] ] M( R

G

2

,

(3)

and from this equation, the virial mass can be estimated simply as

F GF

R Mvir . 210 M ( ? ] pc

DVFWHM ]]] km s 21

G

2

(4)

The virial masses for three clups are tabulated in the last column of Table 2. Compared with LTE masses, the virial masses exceed the LTE masses by more than one order of magnitude. Although the uncertain-

801

ties in LTE mass estimation and the effect of ignoring magnetic field and surface pressure term were taken into account, we can hardly conclude that these clumps are in virial equilibria. In other words, MBM 40 seems not to be a possible site for new star formations, at least, for the three core in the central part of the cloud.

3.5. Correlations between gas components and 100 m m emission For a comparison between the atomic component and the infrared emission, the correlation of N(H I) versus I100 is shown in Fig. 7a for the observed points of WCO ( ; eT *A ( 12 CO 1–0) dv) # 2 K km s 21 . We found, however, ‘no’ relation between these two parameters. These observed points belong to the outer envelope of the molecular component which could have been formed by shock compression. There have been several previous reports for poor correlations in some infrared cirrus clouds (Low et al., 1984; Reach et al., 1994), which may be resulted from the early evolutionary phase of shocked HLCs and transiently heated dust grains. Since the H I and IRAS 100 mm emissions are extended far beyond the observed region, we may need larger area data for further studies.

Fig. 7. (a) N(H I) vs. I100 for the observed position of WCO # 2 K km s 21 . (b) WCO vs. I100 . The first-order regression was derived for the points WCO $ 2 K km s 21 .

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The CO (1–0) and the 100 mm emissions show a better correlation than that for H I and I100 , as shown in Fig. 7b, but there still exist large uncertainties. The first-order regression gives the WCO /I100 5 160.2 K km s 21 (MJy sr 21 )21 for the points WCO $ 2 K km s 21 . The uncertainty is given for the 1 s value of the fitting procedure for the observed points. Previous studies suggest that the ratio WCO /I100 is $ 1 for HLCs, 0.1–0.2 for Galactic plane clouds, and | 0.05 for clouds in the central part of the Galaxy (e.g. Hauser et al., 1984; Weiland et al., 1986; de Vries et al., 1987; Strong et al., 1988; Meyerdierks et al., 1990; Herbstmeier et al., 1993; Reach et al., 1994).This value for MBM 40 locates at the lower end among those of HLCs toward the galactic plane clouds values. Since the I100 values (per hydrogen nucleus) could be larger in HLCs than in dense Galactic plane clouds because of possible smaller dust optical depths in HLCs, the systematic increase of the ratio from dense clouds to HLCs may not be explained with the decrease of the I100 values in HLCs. Thus we explain the WCO /I100 ratios either as the enhancement of CO fractional abundances relative to H 2 in HLCs or as the decrease of CO abundances in dense dark clouds. The enhancement of the CO abundance in the cirrus cloud Draco nebula has been suggested as resulted from the shock by a collision with high velocity H I clouds although the shock fronts is not sufficiently proven to exist (Mebold et al., 1985; Rohlfs et al., 1989; Herbstmeier et al., 1993). It is not clear, however, that shock chemistry can actually enhance the production of CO molecules. No significant relation has been found between the CO and the shocked CH 1 abundances in diffuse clouds (van Dishoeck and Black, 1988). Also the molecular envelope around the MBM 40 cores, the fractional CO abundances are relatively low and show a poor correlation with I100 (for the points WCO # 2 K km s 21 ). In dense molecular clouds, model calculations suggest that the CO fractional abundance decreases steeply after | 10 6 yrs in the cores like TMC-1, which is mainly because of the freeze out onto grains (Ruffle et al., 1997). And there have been several reports that CO could be depleted by accretion onto grain mantles in dense cores or by conversion to complex molecules (cf. Williams, 1993; Kramer et

al., 1999), which can lead to smaller values of WCO /I100 ratio in dense Galactic plane clouds compared to the ratio of HLCs. But the very low depletion or undepletion of oxygen and sulfur has been suggested in HLCs (Heithausen et al., 1998). The CO and other heavy molecular abundances in HLCs are still controversial but a part of their abundance differences with those for dense clouds may be resulted from different degrees of depletion of heavy elements in these two cloud types.

4. Summary MBM 40 is one of the typical high-latitude clouds, located at a distance of | 100 pc from us. The H I emission of this cloud is observed to be extended along the northeast–southwest direction, while the molecular clumps are located almost perpendicularly to the H I emission and its shape may suggest a shock origin of the molecular gas. We identify 3 dense molecular cores of M | 0.2–0.5 M ( , and sizes of | 0.2 pc in diameter. toward the peak position of the core C. The mass of the molecular cores appears to be too small to collapse gravitationally to form a new star. A simplified version of the virial theorem indicates that MBM 40 is not a site for new star formations, at least, for the three cores in the central part of the molecular cloud component as was suggested by Magnani et al. (1996). We found a very poor correlation between the H I and the dust emissions, which may have resulted from the early evolutionary phase of shocked HLCs. The CO (1–0) and the 100 mm emissions show a better relation, but there still exist large uncertainties and we derive the relation WCO /I100 5 160.2 K km s 21 (MJy sr 21 )21 . This ratio is larger by a factor of $ 5 than in dense dark clouds, which may indicate that the CO is less depleted in MBM 40 than in dense dark clouds.

Acknowledgements We are very grateful for supports by Canadian DRAO staffs in making observations and reducing data. S.J.K. and Y.C.M. acknowledge a partial support from KOSEF (R01-1999-000-00022-0). This work was also supported by Strategic National R&D

Y.C. Chol Minh et al. / New Astronomy 8 (2003) 795–803

Progam (M1-0222-00-0005) from Ministry of Science and Technology, Republic of Korea.

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