K-Band Spectroscopy of Luminous Young Stellar Objects

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THE ASTRONOMICAL JOURNAL, 121: 31913206, 2001 June( 2001. The American Astronomical Society. All rights reserved. Printed in U.S.A.

K-BAND SPECTROSCOPY OF LUMINOUS YOUNG STELLAR OBJECTS

MIKI ISHII,1 TETSUYA NAGATA,1 AND SHUJI SATODepartment of Astrophysics, Nagoya University, Furo-cho, Chikusa-ku, Nagoya 464-8602, Japan

YONGQIANG YAO1 AND ZHIBO JIANG1Purple Mountain Observatory, National Astronomical Observatories, Chinese Academy of Sciences, Nanjing 210008, China

AND

HIDEHIKO NAKAYASubaru Telescope, N ational Astronomical Observatory of Japan, 650 North Aohoku Place, Hilo, HI 96720

Received 2001 January 17; accepted 2001 February 27

ABSTRACT

We present spectroscopy from 2.0 to 2.33 km of 32 luminous young stellar objects (YSOs), which arepresumed to be precursors of Herbig Ae/Be stars. From these stars, Brc, CO, He I, and Fe II wereH

2,

found in emission with detection rates of 97%, 34%, 22%, 9%, and 3%, respectively. We compare thespectral features with those of Herbig Ae/Be stars in the literature to investigate the spectral behavior ofintermediate- to high-mass YSOs and to search for their relations to the spectral energy distributions(SEDs). emission is detected only in Class I SEDs with particularly large spectral indices. The detec-H

2tion of emission is related to the degree of the dispersal of circumstellar envelopes, where mol-H

2H

2ecules are probably excited by shocks from outows. On the other hand, Brc emission, which is generally

thought to occur in stellar winds close to the stars, does not depend on the SEDs. This indicates thatstellar wind from luminous YSOs does not change much from the embedded phase to the opticallyvisible phase. CO emission is also independent of the SEDs, but the detection rate is much lower thanthat of Brc emission. Probably, more specic physical conditions regarding circumstellar disks and stellarradiation are necessary for CO emission to take place.

Key words: circumstellar matter dust, extinction stars: premain-sequence

1. INTRODUCTION

Near-infrared (NIR) moderate-resolution spectroscopyhas been used in recent years to investigate the nature ofyoung stellar objects (YSOs). For low-mass YSOs a numberof objects have been observed with this technique (Greene

& Lada 1996 and references therein). Greene & Lada (1996)statistically studied the spectral behavior of nearly ahundred low-mass YSOs through a moderate-resolution(RD 500) spectroscopic survey between 1.15 and 2.42 km.They found that many sources show absorption featuressimilar to those of late-type main-sequence stars, such as CaI, Na I, and CO, and the strengths of the absorption featuresare correlated with spectral energy distributions (SEDs) : theabsorption-line strengths decrease from the Class IIIsources through Class II sources to Class I sources. This canbe explained by a systematic increase in the veiling of anunderlying stellar photosphere from Class III to Class Isources, because infrared excess increase from Class III toClass I. Moreover, the detection rate of the emissionH

2lines is higher for Class I objects, which indicates that theNIR spectral lines change with their circumstellar environ-ments. Therefore the NIR lines can be used to determine theevolutionary stages of YSOs.

For higher mass YSOs, the relations between the NIRspectral lines and their evolutionary stages are not known.The initial mass function predicts very few high-mass starscompared with low-mass stars. In addition, higher massstars evolve faster. Thus the number of high-mass YSOs in a

1 Visiting Astronomer, Okayama Astrophysical Observatory of the

National Astronomical Observatory, Japan.

single star-forming region is limited, resulting in a scarcityof statistical studies.

Among higher mass YSOs, Herbig Ae/Be stars(HAeBes)intermediate-mass (210 YSOs with the TM

_)

Taurilike SEDshave been relatively well studied.Moderate-resolution spectra in NIR wavelengths have beentaken for about 20 HAeBes (Porter, Drew, & Lumsden1998; Harvey 1984; Rodgers & Wooden 1998). In thoseobservations, NIR spectra of HAeBes are characterized byH I emission (sometimes absorption) lines. How about theNIR spectral properties of luminous YSOs with Class IlikeSEDs, which may be regarded as embedded HAeBes? Spec-troscopic observations of this type have been made byseveral authors (Thompson, Thompson, & Campbell 1981;Walther et al. 1991; Thronson et al. 1980; Thronson &Thompson 1982; Scoville et al. 1983; Porter et al. 1998).

In those observations, H I emission lines are detectedtoward all the objects, and He I, Fe II, Na I, CO emis-H

2,

sion lines are detected toward some of them. However, inmost cases each observation was limited to one or twowell-known YSOs. In this paper, we present K-band spectraof 32 luminous YSOs, which, in the unbiased IRAS survey,show large infrared excesses similar to those of low-massClass I objects. What kind of spectral features do theyhave? Do the features show any dierence from those ofHAeBes? Do the features have any relation to the SEDs?By pursuing these questions, we aim to understand theK-band spectral features of intermediate- to high-massYSOs and to search for links to their evolutionary stages.

1.1. Sample

The target objects are IRAS sources identied as lumi-nous YSOs by Campbell, Persson, & Matthews (1989).

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Campbell et al. (1989) selected bright IRAS point sourceswithin ^5 of the Galactic plane, with mid- to far-infraredenergy distributions similar to YSOs. Using NIR photo-metry, they further selected those with NIR color excess,which is interpreted as thermal emission from circumstellardust heated to D1000 K, and listed them as YSO candi-dates. We refer to the sources as the CPM objects in thefollowing.

CO (J \ 10) emission has been detected from theregions around more than 90% of the CPM objects, indi-cating their association with molecular clouds (e.g., Wouter-loot & Brand 1989). Luminosities are estimated from thekinematic distances. In Table 1 we list the luminosities, dis-tances, and references. Kinematic distances are adopted in

most cases, and photometric distances of associated sourcesare adopted for four objects that are located in the anti-galactic direction (objects 15, 16, 18, 21). The luminositiesare computed from the IRAS ux densities using the equa-tion given by Casoli et al. (1989), which attempts to estimatethe total emission from dust between 5 km and 1 mm. TheIRAS luminosity is a good approximation of the bolometricluminosity for the infrared sources.2

2 For comparison, we also calculated the uxes between D1 and 100

km using the data from Campbell et al. (1989) and IRAS. The results agreeswith the IRAS uxes derived from the equation by Casoli et al. (1989)within ^30%, except for objects 07 and 40, whose NIR-to-IRAS uxes areabout 2 times brighter than the IRAS uxes.

TABLE 1

OBSERVED SOURCES

No.a Sourceb Datec Exp.d (s) Standarde df(kpc) Ref.g LIRAS

h (Lsun

) ai Dustj(1) (2) (3) (4) (5) (6) (7) (8) (9) (10)

01 . . . . . . 00211]6549 1999 Nov 14 960 BS7917 7.1 1 2.7E]04 0.4 A02 . . . . . . 00361]5911 1999 Nov 10 1200 BS1482 3.9 1 1.1E]03 0.1 E

04 . . . . . . 01145]6411 1999 Nov 13 1620 BS8607 6.7 1 9.2E]03 0.8 A/E05 . . . . . . 02445]6042 1999 Nov 13 1260 BS8607 3.4 1 1.8E]03 0.3 A06 . . . . . . 02497]6217 1999 Nov 13 1620 BS2383 5.0 1 1.4E]03 0.6 A/E07 . . . . . . 03134]5958 1999 Nov 10 1200 BS8463 1.0 2 5.2E]01 [0.7 N12 . . . . . . 04064]5053 1999 Nov 16 1440 BS2647 8.4 3 3.6E]04 0.3 N13 . . . . . . 04579]4703 1999 Nov 13 1500 BS8607 2.7 1 3.9E]03 1.6 A15 . . . . . . 05137]3919 1999 Nov 12 1260 BS2025 4.3 4 7.9E]03 1.1 A16 . . . . . . 05198]3325 1999 Nov 10 1800 BS8463 4.4 5 8.1E]03 0.6 E18 . . . . . . 05355]3039 1999 Nov 13 1260 BS8607 1.8 6 2.6E]03 0.5 E21 . . . . . . 05439]3035 1999 Nov 18 1800 BS2110 4 7 5.2E]03 1.4 E23 . . . . . . 05568]3206 1999 Nov 13 1620 BS2383 1.1 8 3.0E]02 0.5 A25 . . . . . . 06210]1432 1999 Nov 16 1800 BS8463 4.7 1 1.5E]03 0.2 E26 . . . . . . 06294]0352 1999 Nov 12 1800 BS2629 1.8 1 2.8E]02 0.0 E27 . . . . . . 06335]1057 1999 Nov 18 1620 BS2647 5.0 1 1.5E]03 0.0 E28 . . . . . . 06351[0055 1999 Nov 10 840 BS2714 6.6 1 4.8E]03 0.5 N

31 . . . . . . 06535]0037 1999 Nov 18 1800 BS2714 5.4 1 1.4E]03 0.4 N35 . . . . . . 20216]4107 1999 Nov 13 1620 BS7784 3.3 9 1.1E]04 1.4 A36 . . . . . . 21558]5907 1999 Nov 13 1200 BS2383 10.4 3 9.0E]04 0.9 N38 . . . . . . 22539]5758 1999 Nov 13 1800 BS2383 5.8 1 2.9E]04 0.5 E40 . . . . . . 23390]6524 1999 Nov 10 1080 BS1482 . . . . . . [0.3 N02@ . . . . . . 06041]3012 1999 Nov 13 1800 BS2383 4.7 8 2.2E]03 0.1 N03@ . . . . . . 06040]2958 1999 Nov 13 1800 BS2383 4.7 8 2.9E]03 0.2 A04@ . . . . . . 06562[0337 1999 Nov 12 1260 BS2838 7 10 1.7E]04 1.3 E09@ . . . . . . 19442]2427 1999 Nov 18 1800 BS7677 2.1 11 4.0E]04 2.4 A/E11@ . . . . . . 20236]4058 1999 Nov 14 1800 BS7784 1 12 2.6E]02 0.6 E12@ . . . . . . 20266]3544 1999 Nov 14 1620 BS7917 . . . . . . 0.3 E13@ . . . . . . 20272]4021 1999 Nov 14 1440 BS309 2 13 2.1E]03 0.9 E15@ . . . . . . 21300]5102 1999 Nov 10 1200 BS1482 5.4 1 9.3E]03 1.1 A16@ . . . . . . 21336]5333 1999 Nov 16 1800 BS8463 8.5 1 4.2E]04 1.4 A/E17@ . . . . . . 21413]5442 1999 Nov 16 1980 BS8607 7.7 1 1.9E]05 2.2 A

a Reference number of sources listed in Table 1 and Table 3 of Campbell et al. (1989). A prime follows the number for the sources inTable 3.

b IRAS point source name.c Date of the observations.d Exposure time.e Standard star.f Distance.g Reference to the distance.h IRAS luminosity.i Spectral index computed over 2.2 to 25 km.j 3 km dust features observed by Ishii et al. (1998). A, E, A/E, N represent those with 3.1 km ice absorption, those with 3.3 kmH

2O

UIB emission, those with both of the two features, and those with neither of the features, respectively.REFERENCES.(1) Wouterloot & Brand 1989; (2) Persi, Palagi, & Felli 1994; (3) Wouterloot, Brand, & Fiegle 1993; (4) Casoli et al.

1986; (5) Massey, Johnson, & Degioia-Eastwood 1995; (6) Snell et al. 1988; (7) Tapia et al. 1997; (8) Kawamura et al. 1998; (9)MacCutcheon et al. 1991; (10) Bachiller, Perez, & Garcia-Lario 1998; (11) Barsony 1989; (12) Odenwald & Schwartz 1989; (13)Odenwald, & Schwartz 1993.


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No. 6, 2001 SPECTROSCOPY OF LUMINOUS YSOS 3193

Since the luminosities of the CPM objects are D102104they are mostly early-A to B stars, and the luminosityL

_,

range overlaps with that of HAeBes (D10105 e.g.,L_

;Berrilli et al. 1992). We will examine the possibility that thesample includes more massive stars that could develop H IIregions. It is likely that few of the CPM objects are thosedeveloping compact H II regions because (1) they were selec-ted by Campbell et al. (1989) on the basis of not beingidentied with compact H II regions (see also Persson &

Campbell 1987), and (2) their typical infrared SEDs aredierent from those of the compact H II regions (Henning,Pfau, & Altenho 1990). However it is possible that a fewCPM objects are associated with ultra compact H II(UCH II) regions. Radio continuum observations were per-formed for eight objects, and nondetections were reportedfor ve of themobjects 01 by Wilking et al. (1989), 13 and15 by Molinari et al. (1998), 35 by McCutcheon et al. (1995),and 04@ by Garcia-Lario et al. (1993). For the remainingthree (objects 38, 09@, and 17@) UCH II regions were detected.For object 38, 3.6 cm radio continuum was detected byJenness, Scott, & Padman (1995) with ux of 2.6 mJy. If weassume an optically thin thermal bremsstrahlung emissionwith an electron temperature of 104 K as the origin of theradio emission, the photon ux of the Lyman continuum isestimated to be 1045.91 s~1 (Brown 1987), which corre-sponds to a B1B0.5 ZAMS star (Panagia 1973). For object09@, the UCH II region G60.884[0.128 was found away6A.1from the position of the NIR source, and the observed radioux corresponds to the spectral type of earlier than B0.5(Kurtz et al. 1999). However we suspect that the NIR sourceis not directly connected with the UCH II region, because (1)the separation between them is large, and (2) we detectstrong Brc emission apart from the NIR source, nearer tothe UCH II region (see the Appendix). For object 17@, Miral-les, Rodriguez, & Scalise (1994) found the UCH II region(21413]5442S) at the position of the NIR source (withinthe separation of 1.4A) and the spectral type of the ionizing

source is estimated to be O7.5 by Shepherd & Churchwell(1996). From the above, except for object 17@, the CPMobjects are probably intermediate-mass stars, and someobjects with luminosities of more than a few times 104 L

_border on the class of OB stars.

In order to quantify their SED class, we calculated the2.225 km spectral indices a of the CPM objects from therelation

a\d log jFj

d log j. (1)

The K-band (and L-, M-, N-band when available) photom-etry by Campbell et al. (1989) and IRAS 12 and 25 kmuxes were used for the calculation. IRAS uxes were notused for two objects (16 and 38), which are classied astype D SEDs with mismatched short- and long-wavelength uxes by Campbell et al. (1989). Since the twosources show bright nebulosity in NIR (the J-, H-, K-bandimages; Ishii et al. 2001) IRAS emission must be extended.The N-band data for object 05 were not used because of the10 km absorption (see Fig. 7 in Campbell et al. 1989). Theresult, listed in Table 1, is that all but object 7 have nearlyzero to positive values of a. Thus the CPM objects showSEDs similar to those of low-mass Class I objects, exceptfor object 07 showing a Class II SED. The CPM objectsmay be precursors of HAeBes, many of which have SEDs

similar to those of low-mass Class II objects (Hillenbrand etal. 1992). The evolutionary link between the CPM objectsand HAeBes was discussed in Ishii et al. (1998) by compar-ing their infrared SEDs.

Thus, the CPM objects are regarded as intermediate- tohigh-mass embedded YSOs from the facts of (1) emissionfrom circumstellar dust (near- to mid-infrared excess), whichis characterized as Class I SEDs, (2) association with molec-ular clouds (CO J \ 10 line), and (3) their luminosity(IRAS uxes).

The CPM objects have been also classied into four typesby Ishii et al. (1998) according to dust features in the 3 kmbandthe 3.1 km ice absorption and the 3.3 kmH

2O

unidentied infrared-band (UIB) emission. The four typesare objects with the ice absorption (marked as A in col.[10] of Table 1), those with the UIB emission (marked as E ), those with both the ice absorption and the UIB emis-sion (marked as A/E ), and those with neither of the twofeatures (marked as N ). Each of the four types shows adistinct SED. In NIR wavelengths SEDs of the objects withthe 3.1 km ice absorption are redder than those withH

2O

the 3.3 km UIB emission; the color temperature ofTK~L

D1000 K seems to be the dividing point. In mid-infrared

wavelengths, SEDs of the objects with the ice absorp-H2Otion are similar to those with the UIB emission, and theyare redder than the SEDs of the objects with neither the iceabsorption nor the UIB emission. Ishii et al. (1998) inter-preted these variations in SED as dierent evolutionaryphases: they regarded the objects with the ice absorption asthe youngest type, and they evolve into those with the UIBemission and into those without either of the dust features.

2. OBSERVATIONS

The observations were carried out on 1999 November1018, at the 1.88 m telescope in Okayama AstrophysicalObservatory, using the Okayama astrophysical system forinfrared imaging and spectroscopy (OASIS, Yamashita et

al. 1995; Okumura et al. 2000). OASIS is equipped with aNICMOS3, 256] 256 HgCdTe array, with each pixel cor-responding to The spectral resolution was j/*jD 5000A.96.with a 300 mm~1 grating and a wide slit. The wave-2A.4length coverage was set to approximately 2.0 to 2.33 km.The slit length was about 230A and was aligned along theeast-west direction.

We observed 32 CPM objects. Since most of them werenot seen on a slit viewer with image intensier CCDcamera, the objects were rst searched for at K with theOASIS imaging mode before moving to the spectroscopicmode. In the case of a multiple source, the slit was centeredon the K-band peak, as noted individually in the Appendix.Each object was observed about 10 times to reduce theeects of bad pixels with the telescope dithered along the230A long slit. The exposure time was 2 to 3 minutes,depending on source brightness. The telescope was guidedduring the exposure time by monitoring the nearby opti-cally visible stars (or objects themselves when they werevisible) with the slit-viewing camera. We also observed stan-dard stars (A V and G to K) for atmospheric correction, aswell as for ux calibration. The dierence in air massbetween the CPM objects and the standard stars was lessthan 0.1. Weather conditions were not stable on November10 and 16. Typical seeing conditions were D2A FWHM.The seeing is good and stable on November 18 (D1A.6FWHM), and slightly poor on November 16 (D3A


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FWHM). We present the date of the observations, the totalintegration time, and the standard star for each object inTable 1.

3. DATA REDUCTION AND RESULTS

All data were reduced with IRAF.3 Each frame was dark-subtracted and at-elded. The at eld was constructed bya set of two dome at frames taken with a illuminating lampon and o. Wavelength calibration was done for each frame,

because the wavelength sometimes shifts on the arrayduring the observations. OH air-glow lines were used toobtain a wavelength solution; we referred to the UKIRTWeb site4 for the wavelengths of the OH lines. The accuracyof the wavelength determination was estimated to beD0.001 km. Background emission was subtracted by ttinga third-degree polynomial function to the o-source regionalong the long slit (5A50A away from the object) to removethe OH lines, thermal background, and nebular emission.

3 IRAF is distributed by the National Optical Astronomy Observa-

tories, which are operated by the Association of Universities for Researchin Astronomy, Inc., under cooperative agreement with the National

Science Foundation.4 See http://www.jach.hawaii.edu/JACpublic/UKIRT/astronomy/calib/oh.html.

Then individual spectra were extracted along 4 to 6 pixelwide aperture using the APALL task. The extracted spectrawere leveled with each other and median combined toproduce a nal spectrum. Spectra of objects were furtherdivided by A V standard stars reduced in the same mannerto remove atmospheric absorption features. Brc absorptionin the A V standards had been removed with the SPLOTtask after being divided by the late-type standards (G8 toK7), which show little Brc absorption (Wallace & Hinkle

1997). Finally the spectra were multiplied by Planck func-tions at blackbody temperatures corresponding to the A Vstandards. The blackbody functions were normalized to theux at K. The K magnitudes were taken from the infraredcatalog of Gezari et al. (1993), and those that are not in thecatalog were derived from the V magnitudes, assuming theV[K colors of the spectral type (Tokunaga 2000). Thereduced spectra are shown in Figure 1. The accuracy of uxlevel was estimated to beD20% from the change of signalsobtained at dierent slit positions. For ve objects (12, 28,31, 12@, 16@), however, the accuracy is worse (30% to 50%).This is probably because of unstable weather or unsuc-cessful guiding during the exposures. The use of V[Kcolors or cataloged magnitudes to derive the K magnitudes

of standard stars also causes the ux level to be uncertain,although the uncertainty may be smaller than that because

FIG. 1.Medium-resolution (j/*j D 500) spectra of the observed sources. Fluxes are shifted vertically for clarity.


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FIG. 1.Continued

of the spill of stellar light from the slit during the obser-vations. In addition, it is very likely that some of theobserved YSOs are variable, which also causes the uncer-tainty of ux level. However we did not include these factorsto estimate the uncertainty of the line ux in Table 2. It wasestimated from the change of signals obtained at dierentslit positions (D20%50%, as noted above) and from therms noise of the local continuum.

We compared the reduced spectra with the low-resolution spectra (j/*jD 40) taken with the PASP2(Prism Array Spectrophoto/Polarimeter 2, Ishii et al. 1998)to check the consistency between the two data. While SEDsshow the same trend between the OASIS and the PASP2data, ux levels with OASIS are generally lower than thosewith PASP2. This is most likely because smaller aperture

wide slit) was used for OASIS than for PASPS2 (5A or(2A.49A diameter), and most of the CPM objects are accompa-nied with nebulosity at K (Ishii et al. 2001). For two objects(38 and 09@), however, uxes with OASIS are brighter thanthose with PASP2, even if the errors are considered (weassume 20% for OASIS data and 15% for PASP2 data).The uxes with OASIS are 1.4 times and 3 times brighterthan those with PASP2 for objects 38 and 09@, respectively.On the other hand, photometries by Campbell et al. (1989)measured with a 8A aperture are 1.3 times and 3 times

brighter than the OASIS data for objects 38 and 09@, respec-tively. It is possible that the dierent positions wereobserved between the OASIS and the PASP2, since the twoobjects are accompanied by bright nebulae (Ishii et al.2001). In particular the PASP2 observation for object 09@may be o the K-band peak, because it is too faint com-pared with the OASIS and Campbell et al.s observations.

We searched for the molecular and atomic features thatYSOs and/or normal stars show in the K-band spectraH I (2.166 km), (2.122, 2.223, 2.248 km), He I (2.059, 2.113H

2km), He II (2.189 km), Fe II (2.089 km), Mg I (2.281 km),Mg II (2.137/2.144 km), Na I (2.206/2.209 km), Ca I (2.261/2.263/2.266 km), and CO (2.294, 2.323 km). We regard thefeatures as detected if they are more than 3 p abovethelocal continuum. Since theatmosphericabsorption at thewavelengths around He I 2.059 km is deep and changes verysteeply with wavelength, we were cautious of the detectionof the line; we excluded the lines that could be due to theatmospheric features, and those that are less than 5 p of thelocal continuum, from the detection. The detected featuresare Brc emission (31 objects), emissions (11 objects), COH

2bandhead emissions (seven objects), He I emission (threeobjects), and Fe II emission (one object). No absorptionfeature is found in the spectra of the CPM objects. The linewidths are not resolved with our observations, so their


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TABLE 2

EQUVALENT WIDTHS AND FLUXES OF THE OBSERVED OBJECTS

Brc H2

V \ 10 S(1) CO V \ 20

NO.a Flux (W m~2) EW (A )b Flux (W m~2) EW (A ) Flux (W m~2) EW (A )

01 . . . . . . 2.0E-16 ^ 2.3E-17 4.7 ^ 0.6 \7.8E-18 \1.6 \4.4E-18 \4.102 . . . . . . 4.1E-17 ^ 3.8E-18 3.2 ^ 0.3 \1.9E-18 \1.2 \6.1E-18 \1.904 . . . . . . 1.8E-16 ^ 5.3E-17 20.4 ^ 0.3 \7.7E-18 \1.3 4.1E-17 ^ 1.2E-17 4.3 ^ 0.7

05 . . . . . . 6.0E-17 ^ 4.8E-18 8.5 ^ 0.2 \1.1E-18 \0.8 \2.3E-18 \2.806 . . . . . . 6.2E-17 ^ 7.1E-18 10.3 ^ 0.4 1.6E-17 ^ 1.8E-18 2.7 ^ 0.5 \6.4E-18 \6.307 . . . . . . 7.1E-17 ^ 6.4E-18 2.1 ^ 0.3 \5.1E-19 \0.8 \1.5E-17 \6.812 . . . . . . 5.0E-16 ^ 1.8E-16 3.4 ^ 0.5 \1.0E-16 \1.3 9.7E-16 ^ 3.6E-16 5.9 ^ 0.413 . . . . . . 1.0E-17 ^ 1.8E-18 12.6 ^ 2.3 4.8E-18 ^ 8.9E-19 7.4 ^ 2.4 \1.4E-19 \6.215 . . . . . . 2.3E-17 ^ 1.4E-18 8.1 ^ 0.6 1.5E-17 ^ 9.4E-19 5.7 ^ 1.2 3.1E-17 ^ 2.0E-18 9.0 ^ 1.316 . . . . . . 6.3E-17 ^ 5.5E-18 30.4 ^ 1.1 1.4E-17 ^ 1.2E-18 6.3 ^ 0.9 \2.8E-18 \8.718 . . . . . . 1.1E-16 ^ 1.8E-17 4.6 ^ 0.3 \5.2E-18 \0.9 1.3E-16 ^ 2.1E-17 4.9 ^ 0.421 . . . . . . 3.8E-16 ^ 4.1E-17 78.8 ^ 0.7 1.7E-17 ^ 1.8E-18 3.5 ^ 0.5 \3.4E-18 \5.623 . . . . . . 4.8E-16 ^ 9.9E-17 10.9 ^ 0.3 \2.3E-18 \0.8 \9.2E-17 \2.725 . . . . . . 1.1E-16 ^ 1.5E-17 10.6 ^ 0.4 \4.2E-18 \1.2 \5.4E-18 \2.926 . . . . . . 6.3E-17 ^ 1.3E-17 4.3 ^ 0.3 \9.7E-18 \0.5 5.3E-17 ^ 1.1E-17 3.4 ^ 0.727 . . . . . . 6.6E-17 ^ 2.1E-17 8.3 ^ 0.3 \1.3E-18 \1.3 \4.7E-18 \3.028 . . . . . . 1.0E-16 ^ 4.3E-17 8.4 ^ 0.8 \2.8E-17 \1.3 \1.3E-17 \2.831 . . . . . . 1.2E-17 ^ 6.0E-18 7.7 ^ 0.8 \6.3E-18 \2.9 \1.3E-17 \5.3

35 . . . . . . 1.0E-17 ^ 2.2E-18 3.2 ^ 0.8 1.6E-17 ^ 3.2E-18 5.2 ^ 1.0 \3.4E-18 \4.636 . . . . . . 1.9E-16 ^ 5.8E-17 3.5 ^ 0.5 \1.4E-18 \1.3 \8.2E-17 \4.138 . . . . . . 2.3E-17 ^ 3.6E-18 5.8 ^ 0.6 2.0E-17 ^ 3.2E-18 5.2 ^ 0.6 \3.5E-19 \3.240 . . . . . . 4.1E-16 ^ 4.9E-17 6.4 ^ 0.6 \1.2E-18 \1.0 3.4E-16 ^ 4.0E-17 5.1 ^ 1.202@ . . . . . . 4.6E-16 ^ 7.8E-17 17.8 ^ 0.4 \1.2E-17 \1.1 \1.3E-18 \3.103@ . . . . . . 6.7E-17 ^ 8.2E-18 2.9 ^ 0.4 \4.1E-18 \0.9 2.4E-16 ^ 3.0E-17 10.3 ^ 0.604@ . . . . . . 4.9E-16 ^ 9.7E-17 93.4 ^ 1.5 3.6E-17 ^ 7.0E-18 6.7 ^ 0.4 \7.7E-18 \7.709@ . . . . . . \7.3E-19 \3.8 2.4E-17 ^ 5.0E-18 10.5 ^ 1.1 \6.8E-19 \6.611@ . . . . . . 1.7E-16 ^ 4.0E-17 12.9 ^ 0.4 \9.3E-18 \1.0 \2.2E-17 \1.512@ . . . . . . 1.4E-16 ^ 5.5E-17 28.6 ^ 1.0 \9.5E-18 \2.0 \1.2E-17 \9.213@ . . . . . . 5.9E-17 ^ 7.9E-18 7.9 ^ 0.6 \6.7E-18 \1.6 \5.9E-19 \4.915@ . . . . . . 5.8E-17 ^ 6.9E-18 14.5 ^ 1.2 \2.3E-18 \1.9 \1.1E-17 \5.116@ . . . . . . 6.8E-17 ^ 2.3E-17 30.5 ^ 1.4 3.1E-17 ^ 1.1E-17 14.0 ^ 1.9 \2.7E-18 \3.817@ . . . . . . 2.7E-16 ^ 4.1E-17 69.6 ^ 1.1 1.5E-17 ^ 2.3E-18 4.5 ^ 0.3 \3.6E-18 \4.2

a Same as col. (1) of Table 1.

b A positive EW corresponds to emission.

FWHM is less than 600 km s~1, with the exception of theline width of V \ 10 S(0) for object 04@, which hasH

2FWHM ofD700 km s~1.

Equivalent widths were measured over a 0.01 km wideinterval centered at the detected lines, except for CO V \ 20 whose measurement was from 2.289 to 2.313 km. Thecontinuum for the equivalent width measurements wasdetermined by tting a straight line to the pixels shortwardand longward the features with a width ofD0.008 km. Theuncertainty of the equivalent widths was estimated follow-ing Ali et al. (1995), which is given as

J2]Npixels*jp

contcont , (2)

where *j, cont, and are the number of pixelsNpixels

, pcont

within the full extent of the line (7 to 8 pixels), dispersion (13pixel~1), continuum, and rms noise of the continuum,A

respectively. Since the signal-to-noise ratios, arecont/pcont

D100, 1 p of the equivalent widths is typically D0.5 A .In Table 2, the equivalent widths and the uxes of the Brc,

V \ 10 S(1), CO V \ 20 are presented. When a line isH2

not detected, 3 p is presented as an upper limit.We also searched for emission from nebula or nearby

stars along the long slit. When such emission was found,the spectrum was extracted with the same manner as for the

CPM objects. The results are noted individually in theAppendix.

4. DISCUSSION

4.1. Detection Rate

4.1.1. Brc, andCOH2

,

Since Brc, and CO emissions are frequently seen inH2

,the spectra of the CPM objects, we compare the detectionrates of these lines with those in other types of YSOs in theliterature (Table 3). The detection rates in low-mass YSOsare taken from Greene & Lada (1996), which are derivedfrom the systematic, ux-limited survey within the o Ophmolecular cloud. Since equivalent widths of more thanD0.5 are measured in Greene & Ladas observations,Atheir detection limit is similar to or a little better than ourobservations. HAeBes are selected from those listed by The ,De Winter, & Perez (1994), except for WL 16 and 1548 C27,which are regarded as HAeBes by Najita et al. (1996). Thedata for HAeBes are taken from various observations asnoted in column (5) of Table 3. Since the observationalmethod and the detection limit are dierent for those obser-vations, the derived detection rates should be taken withcaution.

For low-mass stars, Brc emission is found with similardetection rates from Class I to Class II objects, and not


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TABLE 3

DETECTION RATES OF Brc, CO, AND LINES IN YSOSH2

Type Brc (%) CO V\ 20 (%) H2

V\ 10 S(1) (%) Reference(1) (2) (3) (4) (5)

Low-mass YSOs:

Class Ia . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 60 10 30 1Flatb . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 60 0 10 1Class IIc . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 50 0 4 1

Class IIId . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 0 0 0 1Intermediate- to high-mass YSOs:

Herbig Ae/Be starse . . . . . . . . . . . . . . . . . . D93 D15 D4 2CPM objects . . . . . . . . . . . . . . . . . . . . . . . . . 97 22 34 3

a 10 YSOs in oOph.b 10 YSOs in oOph.c 24 YSOs in oOph.d Nine YSOs in oOph.e 28 objects for Brc, 26 for CO, and 25 for H

2.

REFERENCES.(1) Greene & Lada 1996; (2) Harvey 1984; Carr 1990; Rodgers & Wooden 1998; Porter et al.1998; Nisini et al. 1995; Geballe et al. 1987; Carr 1989; Simon & Cassar 1984; Hamann & Simon 1986; Clark& Steele 2000; Thompson & Reed 1976; McGregor, Hyland, & Hillier 1988; (3) this paper.

found in Class III objects. Even the highest detection rate inthe low-mass stars60% in Class I objectsis lower than

the rate in the CPM objects (97%). CO emission is rarelyfound in low-mass stars. In the o Oph sample, it is detectedin only one object (WL 16), resulting in the detection rate of10% in Class I objects. Considering that WL 16 is some-times regarded as a Herbig Ae star (e.g., Najita et al. 1996),the detection rate may be nearly 0%. The situation does notchange even if we consider all the objects observed byGreene & Lada (1996); CO emission is detected for onlytwo (WL 16 and IRAS 04239]2436) out of 96 objects withluminosities of 0.1 to 50 Thus we consider that theL

_.

detection rate of CO emission in low-mass stars to be0%2%, which is much lower than the rate in the CPMobjects (22%). The rate in the CPM objects is comparableto the rate found in higher mass stars with luminosities of

1104 (25%, Carr 1989). emission is more frequentlyL_ H2found in less evolved SEDs of low-mass stars. Class Iobjects show the highest detection rate (30%) among thelow-mass sample; the rate is comparable to that for theCPM objects.

The comparison of the detection rates between the CPMobjects and the low-mass stars indicates that (1) the rates ofBrc and CO emission are higher in the CPM objects than inthe low-mass stars, and (2) the rates of in the CPMH

2objects is similar to that in the Class I objects that showsimilar SEDs to the CPM objects. The detection rates ofBrc and CO emissions seem to relate to mass of the stars,while the rate of emission seems to relate to SEDs.H

2For HAeBes the data are taken from many investigations

in the literature, which makes the comparison with theCPM objects complicated. When a feature is not detected inHAeBes we checked whether the detection limit is compara-ble to or lower than ours.

The detection rate of Brc emission is 93% (26/28) inHAeBes, similar to that in the CPM objects. Since HAeBeswithout the emission (V1686 Cyg and WW Vul) are takenfrom Rodgers & Wooden (1998), where EWs of more than 1

are measured, the comparison between the CPM objectsAand HAeBes is not aected by the dierence in the detectionlimits.

Observations of CO emission for HAeBes are taken fromCarr (1989), Geballe & Persson (1987), Porter et al. (1998),

Rodgers & Wooden (1998), Greene & Lada (1996), Harvey(1984), and McGregor, Hyland, & Hillier (1988). The upper

limits of the EWs of the nondetected sources are estimatedto be which is similar to or better than our obser-[4 A ,5vations. The detection rate of CO emission in HAeBes(15%, 4/26) is similar to that in the CPM objects (22%).However, the rate in HAeBes depends on how we dene thesample. Three out of the four HAeBes with CO emission,MWC 349, WL 16, and 1548 C27, do not fulll the originalcriteria as HAeBes by Herbig (1960). When we exclude suchobjects and only select those listed in Table 1 of et al.The(1994), which were originally recognized as true members orpotential candidates of HAeBes, only one (V645 Cyg) out of21 objects shows CO emission, resulting in a detection rateof 5%. On the other hand, there are two other objects withCO emission, CPD [572824 and GG Car (McGregor etal. 1988), which are listed as candidates of HAeBes by Theet al. (1994) but not included in our sample because of theirunknown evolutionary states. Including the two sources tothe HAeBe sample results in the detection rate of 21%(6/28). Thus the detection rate of CO emission in HAeBesranges from 5% to 21%.

Observations of emission for HAeBes are taken fromH2

Rodgers & Wooden (1998), Porter et al. (1998), Clark &Steele (2000), Greene & Lada (1996), Carr (1990), Harvey(1984), Simon & Cassar (1984), Hamann & Simon (1986),and McGregor et al. (1988). The upper limits of the EWs ofthe nondetected sources are estimated to be 12 exceptAfor D5 for HD 259431 (Harvey 1984) and CDA[4211721 (McGregor et al. 1988).6 Thus the most HAeBes

5 We estimate the upper limit of the EWs of the nondetected sources to

be (1)D1 for Rodgers & Wooden (1998) and Porter et al. (1998) from theAminimum of the measured EWs; (2) for Harvey (1984) from the[3 Acontinuum levels in the gures, assuming the lower limit for line uxes ofD4] 10~16 (W m~2); (3)D4 and 23 for McGregor et al. (1988) andA ACarr (1989), respectively, from the upper limits in the tables.

6 We estimate the upper limit of the EWs from (1) the minimum of themeasured EWs for Rodgers & Wooden (1998), Porter et al. (1998), Clark &Steele (2000), and Greene & Lada (1996); (2) the continuum levels in thegures or K magnitudes, assuming the lower limit of line uxes for Harvey(1984), Simon & Cassar (1984), Hamann & Simon (1986), and McGregor etal. (1988); (3) the continuum levels in the gures or K magnitudes and theupper limit of line uxes in the tables in Carr (1990).


8/16


without emission were observed with detection limitsH2

similar to our observation (12 For HAeBes, emis-A ). H2

sion is found in only one object (PV Cep), resulting in adetection rate of 4%. The low detection rate is similar tothat in low-mass Class II objects.

Therefore, the comparison of the detection rate betweenthe CPM objects and HAeBes indicates that (1) Brc emis-sion is found with similar rates in regard to both, (2) it is notclear if CO emission is found in HAeBes with a similar rate

as that for the CPM objects, and (3) the emission is moreH2frequently found in the CPM objects than HAeBes. Trends1 and 3 can be regarded as a relation between the Class Iand Class II SEDs both in the low-mass and higher massstars: while Brc emission is found with similar detectionrate between the two SED classes, emission is moreH

2frequently seen in Class I SEDs than in Class II SEDs.

4.1.2. He I andFe II

He I emission and Fe II emission were found in three andone of our samples, corresponding to the detection rates of9% and 3%, respectively.

Both emissions are found in early-type emission-linestars. For classical Be stars, He I emission and Fe II emission

are occasionally found in stars of spectral types earlier thanB3, and in stars of B7B1, respectively (Clark & Steele2000). The emissions are also seen in the two Herbig Bestars LkHa 101 and MWC 297 (Simon & Cassar 1984;Porter et al. 1998).

For massive YSOs (BN-type objects and Herbig Be stars)observed by Porter et al. (1998), the detection rates of He Iemission and Fe II emission are 50% and 30%, respectivley.The much higher detection rates by Porter et al. (1998) thanthe rates in our sample is puzzling, because Porter et al.sYSOs seem to belong to similar category to ours, as inferredfrom the fact that CPM objects in the southern hemisphereare included in their sample. For He I emission, our strictlimitation to the detection (3) lowers the detection rate. If

we automatically pick up the emission above 3 p of the localcontinuum, in the same way as for other lines, the detectionrate of He I emission would come up to 25%, but it is stilllower than Porter et al.s rate. The low detection rates ofHe I and Fe II emissions in our sample may be because oursample includes smaller number of early-B stars than Porteret al.s sample, supporting our inference that most of ourobjects are intermediate-mass stars (1.1).

4.2. Brc Emission

In Figure 2, the Brc equivalent widths and line lumi-nosities of the CPM objects and HAeBes in the literatureare plotted; The CPM objects and HAeBes are representedby lled circles and triangles, respectively. In the case ofnondetections, upper limits are presented as crosses andslashes for the CPM objects and HAeBes, respectively.

In Figures 2a and 2b, Brc emission is shown againstthe spectral indices ranging from 225 km. Note thata

2h25,

the CPM objects occupy the region of positive indices, indi-cating their Class I SEDs. On the other hand, HAeBes tendto occupy the region of negative indices: most of them showthe indices of [1.6 to 0, which approximately correspondsto Class II SEDs. In Figure 2a, most of the CPM objectsshow equivalent widths of a few to 30 and only three ofA ,them show EWs of more than 70 Except for the threeA .objects, the equivalent widths of the CPM objects aresimilar to those of HAeBes. The similarity between the

FIG. 2.Brc emission: (a) equivalent width vs. spectral index; (b) lineluminosity vs. spectral index; (c) equivalent width vs. source luminosity; (d)line luminosity vs. source luminosity. The CPM objects with Brc emissionare marked as lled circles. For those without a detection, 3 p upper limitsare marked as crosses. We also plot Herbig Ae/Be stars as triangles in thegures; equivalent widths are taken from Porter et al. (1998), Rodgers &Wooden (1998), Greene & Lada (1996), and Clark & Steele (2000); uxesare taken from Nisini et al. (1995), McGregor et al. (1988), and Carr (1990);photometry from 2 to 25 km is taken from Hillenbrand et al. (1992) or theinfrared catalog of Gezari et al. (1993) to compute the spectral index.

CPM objects and HAeBes can also be seen in the line lumi-nosities in Figure 2b. Thus it seems that Brc emission doesnot change from Class I SEDs to Class II SEDs, consistentwith its similar detection rates in Class I and Class II SEDs(4.1.1).

In Figures 2c and 2d, Brc emission is shown against thesource luminosities. In Figure 2c, lower luminosity sourcestend to show small EWs. For example, the sources withluminosities less than 100 show EWs less than 10L

_A ,

while those with luminosities more than 1000 showL_

EWs up to D90 The tendency can also be applied toA .low-mass stars. For example, only three out of 50 low-massstars with Brc emission observed by Greene & Lada (1996)have EWs of more than 10 Thus lower mass stars showA .small EWs, while some luminous stars show large EWs.This conrms the larger detection rate of Brc emission inhigher mass YSOs (4.1.1). It has been shown that the Brcluminosity of HAeBes correlates with the source luminosity


9/16


(Persson et al. 1984). We found similar correlation for theCPM objects in Figure 2d, where the line luminosities areplotted against the source luminosities. Brc emission seemsto be correlated with the sources luminosities rather thanthe SEDs.

Brc emission of the CPM object resembles that ofHAeBes in the (1) equivalent widths, (2) line luminosities,and (3) correlation to the source luminosities. Thus we inferthat the Brc emission of the CPM objects arises with the

same mechanism as that of HAeBes. H I lines of HAeBesand other massive YSOs are thought to arise from stellarwind. The line width indicates the typical wind velocities of100200 km s~1 (Persson et al. 1984). Applying the windmodel to the H I lines of HAeBes derived the mass-loss rateof 10~610~8 yr~1 (Nisini et al. 1995; Benedettini et al.M

_1998). From the model tting, Benedettini et al. (1998)showed that the H I lines of HAeBes occur in an ionizedregion withinD10 R

*.

How could the correlation between Brc and source lumi-nosities be explained? Although previous observations ofembedded YSOs and HAeBes have also shown similarcorrelation, the interpretation varies among the authors(Persson et al. 1984; Nisini et al. 1995; Corcoran & Ray

1998). In the NIR H I line observations of HAeBes, Nisini etal. (1995) interpret the correlation between the stellar lumi-nosity and the line luminosity as the correlation betweenthe stellar wind and the stellar mass (stellar-driven wind).On the other hand, in the observations of [O I] and Ha linesof HAeBes, Corcoran & Ray (1998) interpret the correlationas reecting the correlation between the accretion lumi-nosity and the line luminosity, and they conclude that windsfrom HAeBes arise in the same manner as those from TTauri stars (accretion-driven wind).

4.3. CO Emission

In Figures 3 and 4, the CO equivalent widths and lineluminosities of the CPM objects and HAeBes in the liter-

ature are plotted; the symbols are the same as those inFigure 2.

In Figures 3a and 3b, CO emission is not correlated withthe SEDs over 225 km. Although no CO emission isdetected in redder SEDs objects the tendency(a

2~25Z1),

does not seem to be meaningful, because very embeddedYSOs such as BN, NGC 1333 SSV 13, and S106-IR showthe CO emission (Geballe & Persson 1987; Carr 1989). Theequivalent widths and the line luminosities are similarbetween the CPM objects and HAeBes: i.e., the CO emis-sion does not change from Class I SEDs to Class II SEDs.In Figure 3c, the CO equivalent widths of the CPM objectsare not correlated with the source luminosities. In the rangeof luminosities of 102105 there is no tendency for theL

_,

more luminous sources to have either larger EWs or thehigher detection rate. On the other hand, CO luminosity iscorrelated with the source luminosity in Figure 3d, as waspointed out by Carr (1989).

CO emission and Brc emission are similar in the corre-lation with the stellar luminosity. We investigate the rela-tion between the CO emission and Brc emission. In Figure4a, CO line uxes are plotted against the Brc uxes. Theuxes are converted to luminosities in Figure 4b. In thosegures, CO emission is correlated with the Brc emission forboth the CPM objects and HAeBes. As Carr (1989) pointedout, however, since most of the objects with Brc emission donot show CO emission, the relation between CO and Brc

FIG. 3.CO emission presented in the same way as Fig. 1. Filledcircles, crosses, triangles, and slashes represent the CPM objects with COemission, those without a detection, Herbig Ae/Be stars with the emission,and those without the emission, respectively. The data of Herbig Ae/Bestars are taken from Carr (1989) and Geballe & Persson (1987).

could not be the direct one, although both of them may berelated to stellar luminosity.

The CO emission occurs in a region with a density ofcm~3 and a temperature of 30004000 K (Scoville etZ1010

al. 1983; Carr 1989). Such a dense, hot, and neutral condi-tion could be satised in a circumstellar disk or neutralstellar wind. While the both situations have been modeledby several authors to t the CO observations, the accretiondisk models with mass accretion rates of 10~810~7 M

_yr~1 more successfully reproduce the observed features ofthe CO emission than the wind models where unreasonablyhigh mass-loss rates are needed (Chandler, Carlstrom, &Scoville 1995; Carr 1989; Calvet et al. 1991; Najita et al.1996). In addition, all of the four HAeBes with CO emissionhave some features indicating the presence of disks: asym-metric nebular morphology, the presence of outow, a jet,or an H-H object for V 645 Cyg, MWC 349, and 1548 C27(Goodrich 1986; Schulz et al. 1989; White & Becker 1985;Leinert 1986; Dent & Aspin 1992; Mundt et al. 1984), and aprole of the CO bandhead emission for WL 16 (Carr et al.1993; Najita et al. 1996). Therefore CO emission presum-ably arises in circumstellar disks.


10/16


FIG. 4.Relation between Brc emission and CO emission. (a) CO uxvs. Brc ux. Filled circles, crosses, triangles, and slashes represent the CPMobjects with CO emission, those without the emission, Herbig Ae/Be starswith the emission, and those without the emission, respectively. ( b) CO lineluminosity vs. Brc luminosity. Each symbol is the same as in (a).

In the disk models, disks are heated by the stellar radi-ation in addition to the viscous heating due to accretion.The correlation between the CO luminosity and the stellarluminosity can be explained in terms of the disk irradiationby the central star (e.g., Carr 1989, Fig. 4).

Dierent origins between Brc and COstellar wind andcircumstellar diskwould explain why most of the objectswith Brc emission lack CO emission. The lower detectionrate of CO emission than that of Brc emission implies amore strict condition for CO emission relating to diskproperties such as mass accretion rate or inclination to theline of sight. Calvet et al. (1991) modeled an optically thickaccretion disk with a temperature inversion in the upperdisk atmosphere due to the stellar irradiation. In theirmodel, whether CO emission occurs depends on the massaccretion rate relative to the stellar eective temperature:an increase of the accretion rate for a xed stellar tem-perature cancels the eect of stellar irradiation, resulting inthe CO bands in absorption. For example, at a stellar tem-perature ofD10,000 K, an accretion rate of more than

yr~1 makes CO bands disappear or turn intoZ10~6 M_

absorption lines (Fig. 7 of Calvet et al. 1991). It can be alsopossible that some objects lack the disks. Thus the rarity ofCO emission seems to result from a subtle combination ofthe stellar and disk properties.

Emission4.4. H2

In Figures 5 and 6, the equivalent widths and lineH2

luminosities of the CPM objects and HAeBes in the liter-

ature are plotted; the symbols are same as in Figure 2.In Figure 5a, the CPM objects with emission haveH2

of 0.5, while those without the emission tend toa2h25

Zshow bluer SEDs. In Figure 5b, seven HAeBes taken fromCarr (1990) are plotted together with the CPM objects. InHAeBes whose SEDs are Class IIlike we nd only onedetection of in the literature (4.1.1). The HAeBes withH

2emission (represented as a triangle) also shows relativelyH

2large Thus emission clearly depends on the SEDsa

2h25. H

2from Class I to Class II, and it is detected only in Class ISEDs with particularly large spectral indices.

In Figures 5c and 5d, emission is shown against theH2

source luminosities. In Figure 5c there is no clear corre-

V \ 20 S(1) emission presented in the same way as Fig. 1.FIG. 5.H2

Filled circles, crosses, triangles, and slashes represent the CPM objectswith emission, those without the detection, Herbig Ae/Be stars with theH

2emission, and those without the emission, respectively. The data of HerbigAe/Be stars are taken from Carr (1990).


11/16


FIG. 6.Relation between Brc emission and emission. (a) ux vs.H2

H2

Brc ux. Filled circles, crosses, triangles, and slashes represent the CPMobjects with emission, those without the emission, Herbig Ae/Be starsH

2with the emission, and those without the emission, respectively. (b) lineH

2luminosity vs. Brc luminosity. Symbols are the same as in (a).

lation between the equivalent widths and the source lumi-nosities. In Fig. 5d, the uxes seem to be correlated withH

2the source luminosities as in the case of Brc in Figure 2d.We investigate the relation between emission and BrcH

2emission. In Figure 6a, uxes are relatively constantH

2against Brc uxes. On the other hand, and Brc lumi-H

2nosities seem to be correlated with each other in Figure 6b.We suspect the correlation is only supercial, which comesfrom the scale expansion by the distance squared eect.

To infer the emission mechanism, we calculate the ratioof the 21 S(1)/10 S(1) lines of of the CPM objects.H

2H

221 S(1) is detected in only one object (04@), and the ratio is0.3 ^ 0.1. For the other 10 objects with 10 S(1), upperH

2limits of the ratio are given; 0.3 for object 35 and D0.1 forthe others. Generally, the line ratios ofD0.1 are observed inshocked regions, and those ofD0.5 are observed in photo-dissociation regions (PDRs) (Martini, Sellgren, & DePoy1999; Shull & Beckwith 1982, and references therein). Thesmall ratios for the CPM objects could support the excita-tion by shock. There is another possibility, however, that

the emission occurs in dense PDRs with densities ofH2cm~3 where the 21 S(1)/10 S(1) ratios could beZ105

similar to those in shocked regions (Luhman, Engelbracht,& Luhman 1998; Draine & Bertoldi 1996). For sevenobjects that are marked as E or A/E in Table 1 out ofthe 11 CPM objects with the 3.3 km UIB emission haveH

2,

been observed by Ishii et al. (1998). Since the UIB emissionis closely associated with UV photons it is possible that H

2emission of the seven objects occur in PDRs. Whether the

emission of the CPM objects occurs in shocks or denseH2PDRs cannot be distinguished from the 21 S(1)/10 S(1)ratio alone. It will be necessary to observe several other H

2lines to distinguish the two excitation mechanisms.

On the other hand, there are some observational signa-tures indicating outows or shocks for most of the CPMobjects with Two objects, 09@ and 17@, are known asH

2.

molecular outow sources (Barsony 1989 ; Shepherd &Churchwell 1996), and an additional four objects (06, 15, 35,04@) are candidate outow sources because of their wing(s)or wide velocity widths of their CO (J \ 10) line (Bachiller,Perez, & Garcia-Lario 1998; Wouterloot & Brand 1989;Wilking et al. 1989; McCutcheon et al. 1991). Water maseremission at 22 GHz may also indicate the presence of

shocks because it is closely connected with molecular out-ows (e.g., Codella, Felli, & Natale 1996; Henning et al.1992). The maser was detected for ve objects13, 15, 38,09@, and 17@ (Wouterloot et al. 1993; Henkel et al. 1986;Palla et al. 1991; Henning et al. 1992). Finally, object 38shows bipolar morphology in optical and infrared (Cohenet al. 1989; Ishii et al. 2001). Therefore, the presence ofoutows or shocks is suggested for eight out of the 11 CPMobjects with by the CO (J \ 10) proles, water masers,H

2or a nebular morphology.

The fact that emission is found for those with Class IH2

SEDs (Fig. 5a) indicates that the emission is closely associ-ated with the circumstellar envelopes. In Figure 7, the IRAS(1225)-(2560) two-color diagram, the CPM objects with

( lled circles) show larger (1225) than the CPM objectsH2without (open circles). The redder (1225) color indicatesH

2that emission from circumstellar dust heated to D100300K is stronger in the CPM objects with than in thoseH

2without HAeBes, which are represented as crosses inH

2.

the gure, show bluer (1225) than the CPM objects withemission. Thus the CPM objects with are dierentH

2H

2from the objects without the in the presence of emissionH

2from warm dust. There is another point suggesting theassociation with the extended material: all the CPM objectswith emission are nebulous in NIR (Ishii et al. 2001),H

2clearly indicating the presence of extended circumstellarmatter. We examine the possibility of emission arising inH

2regions away from the stars. The emissions of objects 16H

2and 09@ (and probably 38) are more extended than thestellar continuum, and the emissions separated from theH

2objects are also detected for objects 38 and 09@ (Appendix).For other objects the emissions cannot be distinguishedH

2from the extent of stellar continuum, but the slit width inour observation covers extended regions because the dis-tances of the CPM objects with range from 2 to 8 kpc.H

2Thus it is possible that the emission of the CPM objectsH

2extends to 0.030.1 pc within the slit.

From the association with molecular outows and withthe extended circumstellar material, we infer that emis-H

2sion of the CPM objects arises from the shocks betweenoutows and envelopes.


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FIG. 7.Relation between the emission and IRAS colors. In theH2

IRAS (1225) vs. (2560) diagram, CPM objects with and without H2

emission are plotted as lled circles and open circles, respectively. HerbigAe/Be stars with and without the emission are also plotted as plusesH

2and a triangle, respectively.

It is striking that only one of the CPM objects with H2

shows CO bandhead emission. The rare coincidence of H2

and CO was also noted by Carr (1989), who ascribes it tothe dierent emitting regions. The following points alsopoint to the dierence between and CO. (1) WhileH

2H

2emission is correlated with the IRAS (1225) color (Fig. 7),no correlation with the IRAS color is seen in the case of CO

emission (Fig. 8). (2) While all the CPM objects with areH2nebulous in NIR, about half of the CPM objects with COdo not show nebulosity (Ishii et al. 2001). (3) CO emission isnot detected for objects with intense Brc emission, while thedetection of emission does not depend on Brc emission:H

2the CO emissions are only found in those objects whose BrcEWs are less than 20 while those with emission showsA , H

2the Brc EWs ofD093 (Table 2). The rst and secondApoints may indicate that CO emission is not associated withthe extended circumstellar material. The third point mayindicate that CO emission arises from outer layers of theinner disk, adjacent to the region of H I recombination lineemission.

4.5. Comparison with the 3 km Dust FeaturesIshii et al. (1998) classied the CPM objects into four

types on the basis of the dust features in the 3 km band, andthey found that SEDs of the four types are distinct fromeach other (1.1). In this section, we investigate the K-bandspectral features in their relation to the 3 km dust featuresand to the continuum emission in NIR.

In Table 4, we compare the detection rates of the K-bandspectral features for each type of the 3 km feature. Brcemission and CO emission do not depend on the 3 km dustfeatures. emission is not detected for the objects withoutH

2the dust features. This is mostly related to the fact that theobjects without the dust features are distributed in bluer

FIG. 8.Relation between the CO emission and IRAS colors. In theIRAS (1225) vs. (2560) diagram, CPM objects with and without COemission are plotted as lled circles and open circles, respectively. HerbigAe/Be stars with and without the CO emission are also plotted as crossesand a triangle, respectively.

part of the IRAS (1225) versus (2560) diagram (Ishii et al.1998), which overlaps with the region of the objects without

emission in Figure 7.H2In Figure 9, equivalent widths of (1) Brc emission, (2) CO

emission, and (3) emission are plotted against the spec-H2

tral indices over 1.34.2 km. The CPM objects are markedaccording to the dust features. Note that the objects withthe ice absorption (open circles and crosses) are distrib-H

2O

uted in the redder (more rightmost) part of the gures thanthe objects with the UIB emission ( lled circles and slashes),as we mentioned previously. The gures show that theequivalent widths of the K-band spectral features arerelated to neither the classication using the 3 km dustfeatures nor the SEDs over the H, K, and Lbands.

NIR SEDs of the CPM objects are characterized by emis-sion from circumstellar dust with temperatures of 6001500K (Ishii et al. 1998). Thus the NIR continuum occurs in

TABLE 4

DETECTION RATES OF BRc, CO, AND LINES ACCORDING TO THEH2

3 kM FEATURESa

ICEb ICE and UIBc UIBd No FeatureseType (percent) (percent) (percent) (percent)

Brc . . . . . . 100 (9/9) 75 (3/4) 100 (12/12) 100 (7/7)

CO . . .. . . 22 (2/9) 25 (1/4) 17 (2/12) 29 (2/7)

H2

. . . . . . 44 (4/9) 75 (3/4) 33 (4/12) 0 (0/7)

a The number of detections over the number of samples are in parenth-esis.

b Objects with 3.1 km ice absorption.H2

Oc Objects with both 3.1 km ice absorption and 3.3 km UIB emis-H

2O

sion.d Objects with the 3.3 km UIB emission.e Objects with neither of the 3 km features.


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FIG. 9.Equivalent widths of (a) Brc emission, (b) CO emission, (c) H2

emission are plotted vs. spectral indices computed over 1.34.2 km. Theindices are calculated from the H-, K-, L-band photometry by Ishii et al.(1998). The CPM objects are classied as in Ishii et al. (1998): open circles,lled circles, half-lled circles, and triangles represent the CPM objectswith the 3.1 km ice absorption, those with the 3.3 km UIB emission,H

2O

those with both the 3 km features, and those without the 3 km features,respectively. When the K-band spectral features are not detected, upperlimits are plotted separately according to the 3 km features: crosses,slashes, back slashes, and pluses represent the CPM objects with the 3.1km ice absorption, those with the 3.3 km UIB emission, those withH

2O

both the 3 km features, and those without the 3 km features, respectively.

dierent regions from those where the K-band lines occur:the continuum-emitting region is cooler than the regions ofBrc and CO emission because dust evaporates at tem-peratures higher thanD1500 K, and hotter than the regionswhere emission occursthe circumstellar envelope.H

2Therefore it is not surprising that the K-band lines are

unrelated to the NIR continuum.Since Brc and CO do not change within the Class I to

Class II SEDs (4.2, 4.3), it is natural that the emissions arenot related to the classication using the 3 km dust

ice absorption and UIB emission. On thefeaturesH2

Oother hand, emission is not detected for the objects withH

2neither of the two features, while the emission is detected inboth the objects with ice absorption and those withH

2O

UIB emission with similar detection rates, line uxes, andequivalent widths. The nondetection of emission for theH

2objects without the dust features is consistent to their loweramounts of cold circumstellar dust and their being in amore evolved phase than the objects with the 3 km features(Ishii et al. 1998).

4.6. Relation of the Spectral L ines and the SEDs

It is well conrmed that the evolutionary stages of low-mass young stars are represented by their infrared SEDs.Class I objects, which are surrounded by massive envelopes,

evolve into Class II objects, which are surrounded by cir-cumstellar disks but lack envelopes, while Class III objects,which lack both envelopes and disks, are the most evolved(e.g., Adams, Lada, & Shu 1987). In this context, the behav-ior of the infrared spectral lines is studied in relation to theinfrared SEDs (Greene & Lada 1996). We compare oursample of high-mass YSOs with low-mass YSOs in the rela-tion between the K-band spectral lines and the SEDs.

For low-mass stars, Brc emission is detected both for the

Class I and Class II objects with a similar rate, but it is notdetected for Class III objects (Table 3). In the case of oursample of high-mass YSOs, Brc emission does not dierbetween the CPM objects and HAeBes: i.e., Brc emissiondoes not dier between Class I SEDs and Class II SEDs asin the case of low-mass stars. On the other hand, twoHAeBes with Class III SEDs in our sample (HD 37490 &HD 53367) also show Brc emission with intensities similarto other HAeBes (Fig. 2b). In addition, the detection rate ofBrc emission is much higher for our sample than for low-mass stars. The dependence of emission on Class I toH

2Class III SEDs does not dier between low-mass and high-mass YSOs. The emission is most frequently detected inClass I SEDs, and it is not detected in Class III SEDs. CO

emission is detected in a certain fraction of our sample andseems not to be related to SED, but the emission is rarelydetected in low-mass stars.

In summary, emission depends on the SEDs of bothH2

low-mass and high-mass YSOs, while Brc and CO emissionare more aected by the stellar luminosity than by theSEDs. For high-mass YSOs, evolutionary sequence accord-ing to infrared SEDs is not well established (e.g., Lada &Adams 1992; Hillenbrand et al. 1992; Fuente et al 1998).Even so, the spectral behavior of our sample of high-massYSOs could be explained in the following way, assumingevolution from Class I SEDs (CPM objects) to Class IISEDs (HAeBes). Brc emission arises in the ionized windclose to stars The property of the stellar wind([10 R

*

).

does not change from Class I to Class II SEDs, but ratherdepends on the stellar luminosities. emission occurs inH

2shocks between outows and envelopes. The emission isdetected in those with Class I SEDs surrounded byenvelopes, and it comes to be detected less frequently inthose with Class II SEDs, where outows have already dis-tracted the envelope material. CO emission occurs in dense,hot, and neutral regions, presumably circumstellar disks.The relatively lower detection rate of this emission com-pared with Brc and is probably due to the emissionH

2depending on both the disk and the stellar properties.Although the CO emission does not depend on Class I toClass II SEDs, the rare coincidence with emission couldH

2indicate relation to the evolutionary sequence, or relationbetween the outows and disks.

4.7. He I Emission

He I emission at 2.06 km is detected for three objects (21,04@, and 17@). These objects show some common character-istics. First they have Brc emission with largest equivalentwidths in our sample ([70 In the observations ofA ).massive YSOs by Porter et al. (1998), there are no objectswhose Brc EWs more than Thus Brc emission of the50 A .three objects is especially strong among massive YSOs.Second, the three objects are particularly luminous in oursample (Table 1). They are estimated to be late-O to early-Bstars (Tapia et al. 1997; Bachiller et al. 1998; Shepherd &


14/16


Churchwell 1996) and two of them (objects 21 and 17 @) areassociated with (ultra compact) H II regions (Tapia et al.1997; Miralles et al. 1994). Finally, the three objects show

emission. The rst and second characteristics indicateH2

that He I emission in these objects arises in H II regionsassociated with OB stars. These objects may be the mostmassive in our sample. On the other hand, the presence of

emission probably indicates that the H II regions haveH2

not yet evolved.

4.8. Absence of Absorption Features

We do not detect any absorption features in our sample.For low-mass Class II and Class III objects, absorptions ofNa I (2.21 km) and Ca I (2.26 km), with equivalent widths ofabout [2 and CO (2.29 km), with equivalent widths ofA ,about [5 are detected (Greene & Lada 1996 ; Casali &A ,Matthews 1992). Such absorptions could have beendetected in our observations where uncertainties of equiva-lent widths areD1 For optically visible OB stars, on theA .other hand, He I absorption at 2.11 km is found for late-Oto early-B stars, and He II absorption at 2.19 km is found forO stars (Hanson, Conti, & Rieke 1996). Since the equivalent

widths of these features are about [1 it could have beenA ,possible to detect these features in our observations.The absence of those absorption features may indicate

that our sample includes neither low-mass stars nor OBstars, conrming that our sample is in the same mass rangeas HAeBes. However, it is more likely that the continuumemission from the circumstellar material dilutes the photo-spheric features, as is the case for low-mass Class I objects(Greene & Lada 1996; Greene & Lada 2000). Spectroscopyat higher resolution with high signal-to-noise will be neces-sary to detect the photospheric absorption features in oursample. Observations in the J and H bands, where thermalemission from dust becomes negligible, will also be helpful.

5. CONCLUSIONS

Moderate-resolution spectroscopy from 2.0 to 2.33 kmhas been done for 32 YSOs. These objects are IRAS sourcesidentied as luminous, embedded YSOs by Campbell,Persson, & Matthews (1989), which we refer to as CPMobjects. Since the luminosities of the CPM objects are com-parable to those of Herbig Ae/Be stars, these objects can beregarded as the precursors of optically visible Herbig Ae/Bestars. The main results in this paper are the following.

1. We detected Brc, CO, He I, and Fe II, all in emis-H2

,sion, with detection rates of 97%, 34%, 22%, 9%, and 3%,respectively. No absorption features were found in theK-band spectra.

2. Brc emission was found in most of the CPM objects.The similarity of the detection rate and line intensity tothose of Herbig Ae/Be stars in the literature indicates thatthe wind from luminous YSOs is sustained from the embed-ded phase represented by the CPM objects to the opticallyvisible phase represented by Herbig Ae/Be stars.

emission was found preferentially in the red CPM3. H2

objects with molecular outows, and it was rarely found inthe bluer CPM objects and in Herbig Ae/Be stars. We inferthat the emission arises from the shocked region of out-H

2ows striking envelopes.

4. CO emission was found in the CPM objects, as well asin Herbig Ae/Be stars, independent of the SEDs. The detect-

ability of the emission may depend on disk properties, suchas mass accretion rate or inclination to the line of sight, andon the stellar luminosity.

5. He I and Fe II emissions, indicative of a massive YSO,were detected in only three out of the 32 CPM objects; oursample includes few massive YSOs, precursors of OB stars.

We would like to thank the sta of the Okayama Astro-

physical Observatory for the observations. We also thankD. Kato for the help in the observation. This work has beensupported in part by Grants-in-Aids for Scientic Researchby the Ministry of Education, Science, Sports, and Cultureof Japan. Y. Y. and Z. J. are supported by NSFC grant19803005.

APPENDIX

NOTES ON INDIVIDUAL OBJECTS

02. IRAS 00361]5911 is a binary aligned in north-southdirection with 2A separation. The brighter northern source

]5927@49A)7 was centered on the slit.(0h38m59s.4,05. IRAS 02445]6042 is a binary aligned in southeast-

northwest direction with 3A separation. The brightersoutheast source ]6055@5A) was on the slit.(2h48m25s.7,07. IRAS 03134]5958 shows a Class II SED, which is

the bluest in our sample. In the K-band spectrum, only Brcemission with the smallest equivalent width in our sample isdetected. This object is listed by et al. (1994) as a candi-Thedate HAeBe.15. A bright star located 10A west from IRAS

05137]3919 was also in the slit. This nearby star]3922@19A) is visible on the Palomar Observa-(5h17m12s.9,

tory Sky Survey (POSS) plate. Its K-band spectrum showsBrc in absorption (EW \ 7.2 ^ 1.3 on the blue contin-A )uum.16. IRAS 05198]3325 is a cluster associated with bright

nebulosity (Ishii et al. 2001). In our K-band spectroscopythe brightest object at K ]3328@38A) was cen-(5h23m8s.4,tered on the slit. emission of this object is a little moreH

2extended than the stellar continuum. A spectrum of anearby star ]3328@38A) was also taken. It shows(5h23m7s.6,no line on the blue continuum.21. The kinematic distance derived from the CO (J \ 1

0) radial velocity and a rotation curve of the Galaxy is 16.5kpc (Wouterloot & Brand 1989). However Tapia et al.(1997) concluded from their JHK imaging observations thatthe IRAS source is a cluster which is 34 kpc away from uswith a visual extinction ofD15 mag. We adopt the latterdistance. The Ha image reveals the photoionized regionaround the IRAS source (Tapia et al. 1997). In our obser-vation the K-band peak ]3036@14A), which(5h47m12s.5,corresponds to Gy 2-18 object 11 by Tapia et al. (1997), wascentered on the slit. The spectrum shows 2.06 km He Iemission (EW \ 5.6 ^ 1.2 strong Brc emissionA ),(EW \ 78.8 ^ 0.7 and emissions. In addition, a spec-A ), H

2trum of the nearby star ]3036@12A ; corre-(5h47m12s.3,sponds to Gy 2-18 object 9 by Tapia et al. 1997) was alsotaken. The spectrum is redder than that of the K-band peak,

7 R.A. and decl. (J2000.0) are taken from the K-band images obtained

by Ishii et al. (2001).


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and shows emissions [EW \ 11.3 ^ 1.4 for 10 S(1),H2

A3.7 ^ 0.7 for 10 S(0), and 1.3 ^ 0.9 for 21 S(1)] and Brcemission (EW \ 6.4 ^ 1.2 While the emissions areA ). H

2stronger than those of the peak position, Brc emission isweaker.38. IRAS 22539]5758 is a bipolar nebula. The K-band

spectrum was taken at the peak of the southern nebula]5814@43A), which roughly corresponds to(22h55m59s.8,

the optical peak (position D of Cohen et al. 1989). Cohen et

al. (1989) suggest from their optical spectroscopic obser-vations that the position coincides with an early-to-middleB star seen directly along the axis, which is sufficientlyinclined to our line of sight. The K-band peak also coincideswith the UCH II region (with the separation of detected1A.3)by Jenness, Scott, & Padman (1995). In our observations,

emission is also detected 5A west of the objectH2 ]5814@43A), where no stellar component, but(22h55m59s.2,

a bright nebular knot, is found in the K-band image by Ishiiet al. (2001). The equivalent widths are 23.6 ^ 5.0 for 10AS(1) and 16.3 ^ 5.1 for 10 S(0). The continuum is nearlyAat and a little redder than that of the peak position.04. The K-band spectrum of IRAS 06562[0337 is

unusual among our sample of objects; it shows 2.08 km

Fe II emission (3.5 ^ 1.0 in addition to 2.06 km He IA ),emission (EW \ 12.0 ^ 2.5 emissions, and Brc emis-A ), H

2sion. Although the object had previously been considered aprotoplanetary nebula, CO and CS observations revealedthe YSO nature ; the IRAS luminosity derived from the COvelocity is 2] 104 corresponding to a B0B2 ZAMSL

_,

star (Bachiller et al. 1998). In the 6 cm continuum obser-vation, there was no detection, giving a 3 p upper limit of0.1 mJy (Garcia-Lario et al. 1993). In the optical spectrum,permitted and forbidden line emissions vary dramatically ina timescale of several years (Kerber, Lercher, & Roth 1996;Alves et al. 1998; Garcia-Lario et al. 1993). Kerber et al.(1996) referred to the object as the Ironclad Nebula fromthe presence of a wealth of permitted and forbidden Fe II

lines.09. IRAS 19442]2427 lies in the S87 region. In our

observation the K-band peak ]2435@25A) was(19h46m19s.4,centered on the slit. The spectrum is the only example in oursample that does not show Brc emission and instead showsrelatively strong emission. On the other hand Brc andH

2emissions from the nearby positions were detected: (1)H

2strong Brc emission (EW \ 55.4 ^ 4.0 and emissionsA ) H

2[EW \ 30.8 ^ 2.9 for 10 S(1), 11.9 ^ 2.9 for 10 S(0),Aand 10.2 ^ 2.4 for 21 S(1)] are found 5A east from theK-band peak; (2) weak and Brc emission are found to beH

2extended around the object within D30A east and D10A

west; (3) a spectrum of a nearby star (19h46m20s.4,]2435@25A) was also taken; it shows Brc emission(EW \ 45.7 ^ 5.5 and emission [EW \ 20.1 ^ 2.9A ) H

2A

for 10 S(1), 11.8 ^ 3.0 for 10 S(0), and 5.5 ^ 2.7 for 21S(1)]. The continuum of these nearby emissions is at andbluer than that of the peak position. At the nearby emission(1) ]2435@25A) no stellar component, but a(19h46m19s.8,bright nebular knot, was found in the K-band image byIshii et al. (2001). The position is D4A south from the UCHII region G60.884[0.128, where diuse, extended radiocontinuum emission at 3.6 cm was found by Kurtz et al.(1999).13. The spectrum of a faint nearby star (20h29m4s.3,

]4032@1A), which is located 4A west from object 13@]4032@3A), was also taken, and no lines on the(20h29m4s.6,

at continuum were detected.15. The spectrum of a nearby star (21h31m44s.1,

]5115@37A), which is located 10A west from object 15@]5115@37A), was also taken, and no lines on(21h31m45s.1,

the at continuum were detected.16. IRAS 21336]5333 is a binary aligned in an east-west

direction with 4A separation. While we centered the slit onthe brighter eastern source ]5347@12A), the(21h35m21s.3,western source ]5347@12A) was also on the(21h35m20s.8,slit. Spectrum of the western source shows Brc emission(EW \ 60.6 ^ 2.9 and emission (EW \ 3.9 ^ 1.9A ), H

2A )

on the at continuum which is bluer than that of the easternsource. While the Brc emission of the western source isstronger than that of the eastern one, emission is weaker.H

217. IRAS 21413]5442 is the only object in our sample

that shows He I emission at 2.11 km (EW \ 5 ^ 0.4 inA )addition to 2.06 km He I emission (EW \ 51.7 ^ 0.8 A ), H

2emission, and Brc emission. The IRAS luminosity of2] 105 corresponding to a O6O6.5 ZAMS star, is theL

_,

largest in our sample, which may be related to the strongHe I emissions. The UCH II region (21413]5442S) wasdetected by Miralles et al. (1994) at the K-band peak

]5456@19A) within the separation of and(21h43m1s.5, 1A.4,the spectral type of the ionizing source is estimated to beO7.5 by Shepherd & Churchwell (1996). The radio obser-vation by Miralles et al. (1994) reveals that the spectralindex between 5 and 15 GHz (0.99 ^ 0.01) is dierent fromthat of an optically thin H II region. This may be related tothe bipolar outow found by Shepherd & Churchwell(1996). In our observations a spectrum of a nearby star

]5456@19A), which is located 5A west from(21h43m0s.9,object 17@ was also t aken. I t shows emissionsH

2[EW \ 11.3 ^ 2.4 for 10 S(1), and 10.7 ^ 1.8 for 10A AS(0)] on the nearly at continuum.

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