Disks and Outflows in CO Rovibrational Emission from Embedded, Low-Mass Young Stellar Objects
Gregory J. Herczeg, Joanna M. Brown, Ewine F. van Dishoeck, Klaus M. Pontoppidan
aa r X i v : . [ a s t r o - ph . S R ] J un Astronomy&Astrophysicsmanuscript no. ms c (cid:13)
ESO 2018September 28, 2018
Disks and Outflows in CO Rovibrational Emissionfrom Embedded, Low-Mass Young Stellar Objects ⋆ Gregory J. Herczeg , Joanna M. Brown , Ewine F. van Dishoeck , , & Klaus M. Pontoppidan Max-Planck-Institut f¨ur extraterrestriche Physik, Postfach 1312, 85741 Garching, Germany; [email protected]; Sterrewacht Leiden, Leiden University, P.O. Box 9513, 2300 RA Leiden, The Netherlands; Space Telescope Science Institute, 3700 San Martin Drive, Baltimore, MD 21218, USAReceived December 1, 2010; accepted April 2011
ABSTRACT
Young circumstellar disks that are still embedded in dense molecular envelopes may di ff er from their older counterparts, but arehistorically di ffi cult to study because emission from a disk can be confused with envelope or outflow emission. CO fundamentalemission is a potentially powerful probe of the disk / wind structure within a few AU of young protostars. In this paper, we presenthigh spectral ( R = , ∼ . ′′
3) resolution
VLT / CRIRES M-band spectra of 18 low-mass young stellar objects (YSOs)with dense envelopes in nearby star-froming regions to explore the utility of CO fundamental ( ∆ v =
1) 4.6 µ m emission as a probeof very young disks. CO fundamental emission is detected from 14 of the YSOs in our sample. The emission line profiles show arange of strengths and shapes, but can generally be classified into a broad, warm component and a narrow, cool component. The broadCO emission is detected more frequently from YSOs with bolometric luminosities of < L ⊙ than those with > L ⊙ . The broademission shares many of the same properties as CO fundamental emission seen from more mature disks around classical T Tauri stars(CTTSs) and is similarly attributed to the warm ( ∼ ∼
320 K), narrowlines in CO and in rarer isotopologues. From some objects, the narrow lines are blueshifted by up to ∼
10 km s − , indicating a slowwind or outflow origin. For other sources the lines are located at the systemic velocity of the star and likely arise in the disk. For afew YSOs, spatially-extended CO and H S(9) emission is detected up to ∼ ′′ from the central source and is attributed to interactionsbetween the wind and surrounding molecular material. Warm CO absorption is detected in the wind of six objects with velocities upto 100 km s − , often in discrete velocity components. That the wind is partially molecular where it is launched favors ejection in adisk wind rather than a coronal or chromospheric wind. Key words.
Proplanetary disks; Stars: protostars: profiles; infrared: stars; Techniques: spectroscopic
1. INTRODUCTION
Circumstellar disks play a central role in the growth of youngstars and in the formation of planetary systems (e.g. Shu et al.,1994; Lissauer, 1993). Disks channel material onto the star andprovide an environment for grains to grow into planets. Whilethe star + disk system is still enshrouded in the envelope, thestar builds up most of its mass, likely accreted from the disk.Meanwhile, the disk mass can be replenished by the surroundingenvelope. Eventually the envelope and later the disk disappear,leaving only the star, any planets or planetesimals that may haveformed, and remnant debris.During the Stage 0 phase of young stellar object (YSO) evo-lution, the dense envelope surrounds the star and any disk thatmay be present. Most of the luminosity from the central pro-tostar is reprocessed by cold dust and escapes as far-IR emis-sion (e.g. Andr´e et al., 1993; Robitaille et al., 2006; Crapsi et al.,2008). At the Stage 1 phase of YSO evolution, the envelopeand disk masses are similar. Any disk likely forms during theStage 0 phase. Evans et al. (2009) measured population fractionsand converted them to timescales to estimate that the Stage 0stage typically lasts 0.1-0.16 Myr and that Stage I stage typi-cally lasts ∼ . ⋆ This work is based on observations collected at the EuropeanSouthern Observatory Very Large Telescope under program ID 179.C-0151. support in hydrodynamic models of collapsing cold cores (Bate,2010; Machida & Matsumoto, 2010).When the protostar is still embedded in an optically-thickenvelope, the disk is di ffi cult to study because emission fromthe disk, if present, can be confused with emission from theenvelope. Observations with high spatial resolution can breakthis degeneracy by spatially discriminating between the loca-tion of compact disk emission and extended envelope emission(e.g. Padgett et al., 1999; Eisner et al., 2005; Beckford et al.,2008; Tobin et al., 2010). For example, sub-mm interferometryhas been used to separate a compact emission source, as ex-pected for a disk, and an extended emission source, as expectedfor an envelope (e.g. Keene & Masson, 1990; Wilner et al.,1996; Looney et al., 2000; Brown et al., 2000; Jørgensen et al.,2005, 2007; Lommen et al., 2008; Jørgensen et al., 2009). Awide range of disk-to-envelope mass ratios have been mea-sured for low-mass young stellar objects, as would be expectedfor a sample that includes systems at a range of evolutionaryphases (Lommen et al., 2008; Jørgensen et al., 2009). In mostcases the disk + envelope mass is much smaller than the stel-lar mass, indicating that the main accretion phase has ended.Disk masses are similar for objects in Stage 0 and Stage Istages (Jørgensen et al., 2009), which suggests that once thedisk grows to a certain size, the mass that the disk loses tothe star is replaced by accreting mass from the envelope. Thepresence of disks around Stage I objects is supported by obser- vations of Keplerian emission profiles in cold rotational linesof HCO + (e.g. Hogerheijde & Sandell, 2000; Saito et al., 2001;Brinch et al., 2007; Lommen et al., 2008).These (sub)-mm observations trace molecular emission fromthe disk at large radii, but little is known about the inner AUof young disks that are still embedded in envelopes. If theregion is similar to the inner disks around classical T Tauristars (CTTSs), the chemical and temperature structure can beprobed by the same line diagnostics used to study CTTS in-ner disks (see reviews by Najita et al., 2007; Dullemond et al.,2007). Unfortunately, many of these techniques are di ffi cult toapply to embedded objects because emission from the disk, en-velope, outflow, and cloud can be confused and because embed-ded objects are highly reddened. For example, in Spitzer
IRSspectral surveys of young stars, emission in mid-IR gas linesis detected more frequently and is more luminous from embed-ded objects than from classical T Tauri stars (Lahuis et al., 2007;Flaccomio et al., 2009; Lahuis et al., 2010). However, follow-upof potential disk tracers, including [O I], [Ne II] and H emissionat high spectral resolution (Hartigan et al., 1995; Herczeg et al.,2006; van Boekel et al., 2009; Pascucci & Sterzik, 2009), athigh spatial resolution (Walter et al., 2003; Saucedo et al., 2003;Beck et al., 2008), or in large samples (G¨udel et al., 2010), in-dicate that even for CTTSs, outflow emission often dominatesover any disk emission. Emission from disks and outflows canbe even more confused for Stage 1 objects because they typi-cally drive more powerful outflows than CTTSs. Indeed, emis-sion in H O emission in the near-IR and sub-mm has also beenattributed to disks around embedded objects (Carr et al., 2004;Jørgensen & van Dishoeck, 2010).Fundamental-band ( ∆ v =
1) CO emission in the M-bandhas been demonstrated to be a powerful probe of warm ( ∼ ∼ . − . v =
7, likely indicative of strongUV emission from the stellar photosphere (Brittain et al., 2009;van der Plas et al., 2009)Several studies have shown that warm CO is also presentaround embedded objects. In one case, GSS 30, CO fundamen-tal emission is seen up to 2 ′′ (240 AU projected distance) fromthe central source. In a larger sample, Pontoppidan et al. (2003)detected bright CO emission from 18 of 44 objects observed with VLT / ISAAC at moderate resolution ( R = λ/δλ = ), but didnot analyze this emission because, when spectrally unresolved, multiple absorption and emission components can introducespurious velocity shifts and cause misleading non-detections(see also Pontoppidan et al., 2005a). Brittain et al. (2005) at-tributed CO emission from the embedded object HL Tau to thedisk, based on similarity of the emission profile seen at high res-olution to the profiles from CTTSs that lack envelopes. Earlystudies by Chandler et al. (1993) and Najita et al. (1996) foundvibrationally-excited CO overtone ( v = −
0) emission from twoembedded objects, WL 16 and L1551 IRS 5.Now that CO v = − ffi cult to detect withother gas tracers. Embedded objects should have larger accretionrates than classical T Tauri stars (CTTSs), which may introducedi ff erences between young disks around embedded objects andmore mature disks around CTTSs. Any warm (500 − R = λ/ ∆ λ ∼ , Very Large Telescope .CRIRES is uniquely capable of high-resolution M-band spec-troscopy with relatively broad spectral coverage at high spatialresolution. The data presented here were obtained as part ofa large program to survey CO emission from YSOs at di ff er-ent stages of pre-main sequence evolution (Pontoppidan et al.,2011b). Most observations in this paper were obtained under ex-ceptional weather conditions. In §
2, we describe the observa-tions and the selected sample. In §
3, we describe and analyzethe observed line profiles, including separating the emission intobroad and narrow components, analyzing spatially-extended COand H emission, and finding CO absorption in winds of six ob-jects. In §
4, we attribute the components to di ff erent regions inthe YSO and discuss the implications of our results on disk struc-ture and wind launch regions.
2. OBSERVATIONS
The selected targets are Stage I YSOs that are still embeddedin circumstellar envelopes and that are members of nearby star-forming regions. Most of the targets in our sample were cho-sen based on fundamental CO emission that was previously de-tected in R ∼ ,
000 ISAAC spectra (Pontoppidan et al., 2003,2005a). Most targets are low-mass YSOs based on their lumi-nosity. One object, SVS 20 S, may have a luminosity consistentwith intermediate-mass YSOs. A few other objects may evolveinto Herbig AeBe stars.Because SEDs of heavily-extincted stars can be similar toSEDs of stars embedded in circumstellar envelopes, the presence For simplicity, we use the physical terminology “Stage” rather thanthe observational terminology “Class” throughout the paper, althoughin some contexts “Class” would be the more appropriate term.2erczeg et al.: CO Emission from Embedded Objects
Table 1.
Evidence for envelopes within our sample
Source mm-int. a Compact b Outflow c Ref. d TMC 1A Y Y Y 1,2,3,7GSS 30 Y Y n 1,2,4WL 12 Y Y s 1,2,4Elias 29 Y Y Y 1,2,4,8IRS 43 S + N Y Y s 1,2,4IRS 44 E + W – Y Y 1,4IRS 63 Y Y Y 1,2,4L1551 IRS 5 AB e Y Y Y 3,7,13HH 100 IRS – Y Y 5,9WL 6 – Y n 1,4Elias 32 – Y – 4CrA IRS 2 – – – –HL Tau Y Y Y 3,10,12SVS 20 S + N – Y – 7RNO 91 – Y Y 6,11 a Envelope detected in mm-interferometry. b Presence of cold, compact gas centered at object position. c Refers to cold molecular outflows detected in the (sub)-mm.(s): Molecular outflow is detected in only a single lobe.(n): No molecular outflow detected. d The references are examples of each phenonemon andare not intended to be complete. e Unresolved and hereafter referred to as L1551 IRS 51: Bontemps et al. (1996), 2: Jørgensen et al. (2009)3: Saito et al. (2001), 4: van Kempen et al. (2009)5: van Kempen et al. (2009), 6: Arce et al. (2006)7: Gregersen et al. (2000), 8: Lommen et al. (2008)9: Groppi et al. (2007), 10: Wilner et al. (1996)11: Chen et al. (2009), 12: Cabrit et al. (1996)13: Wu et al. (2009) of an envelope for objects in our sample is typically confirmedby the presence of at least one of the following three criteria: (1)a clearly extended component in the visibility curves of (sub)-mm emission (Lommen et al., 2008; Jørgensen et al., 2009), (2)a compact structure of cold dense gas, such as HCO + , at the sameposition as the target (Saito et al., 2001; Groppi et al., 2007;Lommen et al., 2008; van Kempen et al., 2009), or (3) a coldmolecular outflow seen close to the YSO in both a red- andblue-shifted lobe (e.g. Bontemps et al., 1996; Arce et al., 2006).Table 1 summarizes the evidence for our classification. One ob-ject, CrA IRS 2, is less-well studied and not confirmed as a Stage1 object by these criteria, but is likely embedded in an envelope(Nutter et al., 2005). For the purposes of this paper, both ob-jects in spatially-resolved multiple systems with envelopes areassumed to be embedded (see Appendix A for more details).The classification of these targets as embedded is consistentwith most previous classifications, but can di ff er from classifica-tions that are based on SEDs alone. For example, McClure et al.(2010) used 12-5 µ m colors to classify WL 12, IRS 43, and IRS44 as YSOs with envelopes and WL 6, Elias 29, and IRS 63as YSOs with disks but no significant envelope. When classi-fying the latter three objects as disks, McClure et al. (2010) ex-plain the discrepancy with previous work by suggesting that theenvelopes are tenuous and close to disappearing. This interpre-tion may apply to WL 6 and IRS 63 (van Kempen et al., 2009).However, the sub-mm visibility of Elias 29 shows no indica-tion of any compact structure, which means that the envelopemass must be much larger than the mass of the undetected disk(Lommen et al., 2008). The disks for these three sources maydominate their near-IR spectra, but the sub-mm interferometryrequires the presence of an envelope. Table 2 lists selected properties of our targets from theliterature. The stellar properties, including the e ff ective tem-perature and luminosity of the photospheric emission andthe extinction, are uncertain because of the di ffi culty in de-tecting photospheric features (e.g. White & Hillenbrand, 2004;Doppmann et al., 2005) and in accurately measuring the extinc-tion. In some cases, the A V and L bol are obtained from di ff erentsources. The bolometric luminosities that are measured from thefar-IR / sub-mm SEDs are not particularly sensitive to either ex-tinction or spectral type. All luminosities are corrected for themost recent distance measurement to the parent cloud. In a fewcases, significant discrepancies in derived stellar parameters ex-ist in the literature .The sample likely includes a range of masses for the centralobject, with a general trend that more luminous objects are likelyalso more massive. Unfortunately, the lack of accurate stellarproperties prevents a reliable determination of mass from pre-main sequence evolutionary tracks.The systemic velocity listed in Table 2 is obtained from liter-ature measurements of (sub)-mm lines of the molecular cloud orof the envelope. All CO and H velocities listed in the paper arerelative to the measured velocity of CO and C O absorptionin our M-band spectra.
We used CRIRES (K¨aufl et al., 2008) on the VLT-UT1 to obtainhigh-resolution echelle spectra from 4.6–4.9 µ m of 15 Stage Iobjects, three of which are binaries. Table 3 lists our observationlog .CRIRES has four 1024 x 512 pixel InSb detectors that eachcovers ∼ . µ m in a single integration. Every object in oursample was observed at multiple wavelength settings to coverspectral gaps between chips and to observe lines with a widerange of rotational levels. Table 3 includes the wavelength set-ting used for each object. Online Table A.1 lists the wavelengthranges covered for each setting. Each pixel covers ∼ . ∼ .
25 km s − , in the dispersion direction. The 0 . ′′ . ′′
086 pixels in the dispersion direction leads to R ∼ , ′′ nods in an ABBA pat-tern. Each nod consists of 60s integrations. Random dither pat-terns with a width of 1 ′′ were used for each nod to distribute thecounts over di ff erent pixels at each integration. Total integrationtimes for each setting are listed in Table 3.Most sources in our sample were too faint for the MACAOadaptive optics system, which usually feeds CRIRES. The AOloop was closed for HL Tau and RNO 91, though with a lowStrehl ratio and poor correction. For other objects, the M-bandseeing during our observations was exceptional, typically be-tween 0 . ′′ . ′′
45 as measured from the spatial profile of theM-band continuum emission in the cross-dispersion direction for For IRS 63, the photospheric properties measured byDoppmann et al. (2005) are not adopted here because the v lsr = − . − of photospheric absorption features is highly anomalousrelative to the other Stage I objects in the ρ Oph molecular cloud andis inconsistent with the velocity of CO absorption by the envelopeseen in our data. We speculate that Doppmann et al. (2005) measuredoptically-thin absorption in an outflow and attributed the absorption tothe photosphere. Alternately, the slit rotation may have induced a largevelocity o ff set, in which case the derived e ff ective temperature wouldstill be accurate (T. Greene, private communication). Data are available for download athttp: // http: // / ∼ pontoppi 3erczeg et al.: CO Emission from Embedded Objects Table 2.
Sample Properties
Target d PA aout A V T phot L bol v lsr v babs Flux c pc ◦ mag K L ⊙ km s − km s − µ m Method RefHH 100 IRS 130 – 30 4060 15 5.9 5.8 11.3 ISO 2,8CrA IRS 2 130 – 20 4900 12 (5.9) 6.1 5.40 ISO 8,15WL 6 120 ∼
40 9.8 – 2.4 3.5 4.5 1.4 IRAC 1,23WL 12 120 – 9.8 4000 2.4 3.5 7.6 1.1 IRAC 1, 10IRS 63 120 ∼
240 9.8 – 3.0 3.5 2.0 1.3 IRAC 1,24IRS 43 S + N 120 25 9.8 4400 5.5 3.5 4.1 1.6 IRAC 1,10,17IRS 44 E + W 120 ∼ ∼
135 9.8 – 38 3.5 5.5 24.3 ISO 1,21GSS 30 120 ∼
45 9.8 – 13 3.5 7.5 23 IRAC 1,20HL Tau 160 ∼
45 7.4 4350 6.6 7 8.2 5.49 ISO 3,4,9,22RNO 91 120 155 – – 2.3 0.8 0.5 1.8 16 5,12,18SVS20 S 415 – 30 5900 142 1.9 8.3 3.5 IRAC 10,14,7SVS20 N 415 – 4 3270 0.27 1.9 (8.3) 1.1 IRAC 13,7TMC 1A 160 −
10 – – 2.8 5.6 5.3 1.0 IRAC 6, 3,7, 25L1551 IRS 5 AB 160 256 28 3300 23 6.9 4.7 3.94 ISO 6, 3,7,10,11,19 a Approximate position angle of outflow b v lsr of CO and C O absorption in our CRIRES spectra c − erg cm − s − Å − () indicates assumed value from nearby source.Distances: 130 pc to CrA (de Zeeuw et al., 1999), 120 pc to Oph (Loinard et al., 2008), 415 pc to Serpens(Dzib et al., 2010), and 160 pc to Taurus (Torres et al., 2009)1: Evans et al. (2009), 2: van Kempen et al. (2009) 3:Furlan et al. (2008) 4: Hayashi et al. (1993)5: L¨ohr et al. (2007); 6: Yang et al. (2002) 7: Gregersen et al. (2000) 8: (Nisini et al., 2005)9: White & Hillenbrand (2004); 10: Doppmann et al. (2005); 11: Prato et al. (2009); 12: Chen et al. (2009)13: Oliveira et al. (2009); 14: Ciardi et al. (2005); 15: Meyer et al. (2009); 16: Boogert et al. (2008)17: Grosso et al. (2001); 18: Arce et al. (2006); 19: Pyo et al. (2009); 20: Allen et al. (2002); 21: Ybarra et al. (2006)22: Beckwith et al. (1989); 23: Gomez et al. (2003); 24: Zhang & Wang (2009); 25: Chandler et al. (1996) Table 3.
Observation Log
Target Alt. Names RA a DEC a UT Date t exp (s) airmass FWHM b v bary (km s − ) λ settings S / N d IRS 43 YLW 15 16:27:27 -24:40:51 2008-08-06 960 l d YLW 16a 16:27:28 -24:39:34 2008-04-27 720 1.41 0.56 27 4716 15IRS 44 YLW 16a 16:27:28 -24:39:34 2008-04-30 720 m e YLW 16a 16:27:28 -24:39:34 2008-05-01 1200 1.61 0.86 25 4946 9IRS 44 e YLW 16a 16:27:28 -24:39:34 2008-08-06 1200 1.31 0.33 -17 4868 20IRS 44 e YLW 16a 16:27:28 -24:39:34 2008-08-07 960 1.38 0.32 -17 4716, 4730 30IRS 63 GWAYL 4 16:31:36 -24:01:29 2007-04-26 600 1.02 0.37 28 4716, 4730, 4833, 4929 32WL 12 YLW 2 16:26:44 -24:34:48 2007-09-02 960 1.22 0.31 -20 4716, 4730 48WL 12 YLW 2 16:26:44 -24:34:48 2007-09-04 960 1.25 0.31 -20 4833 45WL 12 YLW 2 16:26:44 -24:34:48 2010-03-03 2280 1.05 0.31 40 4946 60WL 6 YLW 14 16:27:22 -24:29:53 2007-04-27 480 1.20 0.29 27 4716, 4730, 4833 30HH 100 IRS V710 CrA 19:01:51 -36:58:10 2007-09-03 600 1.27 0.43 -17 4716, 4730 170HH 100 IRS V710 CrA 19:01:51 -36:58:10 2007-09-06 480 1.43 0.26 -18 4833 140HH 100 IRS V710 CrA 19:01:51 -36:58:10 2008-08-03 480 1.11 0.44 -6 4710, 4730, 4868,4946 100Elias 32 YLW 17A 16:27:28 -24:27:20 2008-05-03 1200 1.19 0.32 24 4716, 4730, 4833 10Elias 32 e YLW 17A 16:27:28 -24:27:20 2008-08-09 960 1.07 0.45 17 4716 3Elias 32 f YLW 17A 16:27:28 -24:27:20 2008-08-10 960 1.12 0.52 18 4730 7.5Elias 29 e WL 15 16:27:09 -24:37:19 2008-08-09 360 1.37 0.59 -17 4716, 4730, 4868 130CrA IRS 2 CHLT 1 19:01:42 -36:58:32 2007-04-26 480 1.02 0.37 34 4716,4730,4833,4929 110RNO 91 HBC 650 16:34:29 -15:47:02 2007-04-24 600 1.05 0.28 30 4716,4730,4929,4946 55RNO 91 HBC 650 16:34:29 -15:47:02 2007-04-25 1200 1.10 0.26 30 4770,4780 60HL Tau g HBC 49 04:31:38 18:13:58 2007-10-12 600 n h HBC 49 04:31:38 18:13:58 2010-02-19 1080 1.71 0.22 42 4716,4730,4800,4820 150SVS 20 i [EC92] 90 18:29:58 01:14:06 2007-04-24 1200 o j IRAS 04365 + p h Elias 21 16:26:21 -24:23:06 2010-03-04 600 q k HBC 393 04:31:34 18:08:05 2007-10-17 840 1.50 0.82 -4 4716,4730 17 a b In arcsec, includes both the seeing and any real spatial extension in the continuum. c Approximate S / N per pixel in a single setting, typically measured in regions with an average telluric absorption of 10-15%.Position angle of 0 ◦ except where noted: d PA = e PA = f PA = g PA = -150; h PA = i PA = j PA = k PA = Listed position angle is approximate for cases where the slit was aligned with the parallactic angle.
Exposure times same for all settings except where noted: l m n o / q q Fig. 1.
Extraction of two components from the IRS 44 spectra.The main plot shows the extracted counts versus wavelength forIRS 44 E (red) and IRS 44 W (blue) over a small wavelengthrange. The two insets show the spatial profile of the counts inthe cross-dispersion direction for a single pixel (left, within theCO P(21) line, right within the continuum). Two Gaussian pro-files were fit to the cross-dispersed emission profile and weresubtracted to minimize contamination between components dur-ing the extraction. CO emission is detected from the secondary,IRS 44 W (blue) but not from the primary, IRS 44 E (red). Theexcised region is unreliable because of strong telluric CO ab-sorption.each object. Two objects, Elias 29 and L1551 IRS 5, have espe-cially large seeing measurements, which indicates that either theseeing was poor or that the M-band continuum emission is spa-tially extended on 0 . ′′ − . ′′ ff erential atmospheric dispersion between theK and M-bands. Prior to November 2008, our observations weregenerally obtained at low airmasses or at parallactic angle tominimize the e ff ect of atmospheric dispersion. Because most ofour observations did not use adaptive optics, the e ff ect of atmo-spheric dispersion ( ∼
50 mas at airmass of 1.7) is typically muchless than the FWHM of the continuum in the cross-dispersiondirection. A real o ff set between the K-band and M-band con-tinuum emission is also possible for embedded young stars. Incases where the slit was aligned with the parallactic angle, theposition angle did not change by more than 10 ◦ in any observa-tion. All of our observations include strong continuum emission,and e ff ects that may be attributed to di ff erent slit placements arelikely minor.A telluric standard at a similar airmass to our object spec-trum was always observed in an adjacent observation and withan AO correction. Spectra from the object and the telluric stan-dard were wavelength-calibrated using telluric absorption lines.The science spectrum was then divided by the spectrum of thetelluric standard. Di ff erences in spectral resolution between thetelluric standard and the science target were always minor and,when necessary, were accounted for by convolving the spectrumof the telluric standard with a Gaussian profile to match the tworesolutions. Figure 2 demonstrates that high S / N ( ∼ − Fig. 2.
The telluric correction for IRS 43 S + N from two chips inthe 4716 Å setting. The spectrum reaches S / N ∼
100 per pixelin regions with few atmospheric lines (left) and S / N ∼
50 perpixel in regions with many atmospheric lines (right). CO linesare labeled in black and CO lines are labeled in red.tained on dates when the telluric absorption was shifted by largevelocities from the local standard of rest to shift telluric CO ab-sorption away from the central wavelength of CO lines for theobjects. The relative wavelength calibration is accurate to ∼ − .Data were reduced following Pontoppidan et al. (2008). Thepixels on the detector are each 0 . ′′
086 in the dispersion and cross-dispersion direction. When coadding the 2D images, the lightwas shifted onto a common central pixel and resampled onto afiner (0 . ′′ . ′′ § A1.1), which wasmarginally resolved on three nights when the seeing was excep-tional (0 . ′′ ff set between the two components are constant across the detec-tor. The Gaussian profile of one component was subtracted fromthe data. The counts were then extracted from a spectral win-dow around the other component. The extraction windows for Fig. 3.
A sample spectral region, with CO v = − CO v = − O v = − CO v = − Spitzer / IRAC photom-etry for seven objects (van Kempen et al., 2009; Harvey et al.,2007). For binaries, the continuum flux of each companion wascalculated by using the M-band magnitude and the counts ra-tio in the continuum emission measured in our M-band observa-tions. Spatially-extended continuum emission for a given objectwould artificially increase the calculated M-band flux of the cen-tral source. The CO line fluxes are calculated from equivalentwidth measurements and would be overestimated in cases wherethe M-band continuum is spatially extended.Our CRIRES spectra yield a flat M-band continuum flux forevery object, with a relative uncertainty of ∼
30% between 4.5and 5.0 µ m. In the four ISO spectra, the maximum relative fluxdi ff erence between 4.5 and 5.0 µ m is 30%. We also do not cor-rect the absolute or relative CO emission line fluxes for extinc-tion because of the wide range in calculated extinctions for anygiven object, the lack of methodological consistency in extinc-tion calculations for the entire sample, and because extended COemission may not su ff er from the same extinction as the cen-tral source. Applying extinctions that range from A V = − ∼ µ m spectral region for A V =
20 mag).
In several subsections we compare our high-resolution CRIRESspectra to archival VLT / ISAAC M-band spectra of the samesources obtained in 2001–2002. The ISAAC spectra span from 4.53–4.75 µ m with R = − . ′′ To compare extended CO emission to continuum emission fromGSS 30 (see § VLT archive that was obtained with NACO on 19September 2007 as part of
VLT program 079.C-0502 (P.I. Chen).Gaspard Duchˆene (private communication) provided us with anL-band AO image of GSS 30 obtained with
VLT / NACO and pub-lished in Duchˆene et al. (2007).
3. DESCRIPTION OF M-BAND EMISSION FROMEMBEDDED OBJECTS
M-band spectra cover many CO fundamental ( ∆ v =
1) transi-tions, H Pfund- β , H S(9), and the 4.67 µ m CO ice absorptionband. One spectral region of our CRIRES M-band spectra ofembedded objects is shown in Figure 3. The M-band emissionis dominated by a smooth continuum produced by warm dust,likely located in the disk close to the star (e.g., Eisner et al.,2005; Enoch et al., 2009). No photospheric features are detectedfrom any of our targets. CO ice absorption (Pontoppidan et al.,2003) is detected towards most objects in our sample. HLTau lacks any CO ice absorption (see also Whittet et al., 1989;Brittain et al., 2005), which suggests that much of the line-of-sight extinction occurs in a disk or flattened envelope inwhich little solid CO absorption is expected to be observed Fig. 4.
Co-added profiles for CO v = − v = − CO v = − VLT / CRIRES spectrum of HH 100 IRS.The di ff erent observed components are each identified with aphysical region in the system. For scaling, the peak of the v = − ∼ .
38, normalized to the adjacent continuum.(Pontoppidan et al., 2005b). Strangely, GSS 30 also shows onlyvery weak ice absorption (see also Pontoppidan et al., 2002).Absorption in gaseous CO and isotopologues is detected to-wards every object. Narrow absorption lines are caused by theenvelope and parent molecular cloud, while broad absorptionarises in the wind.Emission in CO v = − v = − ffi cient to pro-duce detectable emission in CO and C O lines, is detectedfrom 9 sources. These di ff erent components are more easily dis-tinguished in the v = − v = − CO lines, combine to create the CO v = − CO lines are typically narrower than the v = − § § J transitions that were observed and that have minimal ormoderate contamination from telluric absorption. Some objectsshow strong emission in the narrow, optically-thick component,some objects show strong emission in broad, vibrationally ex-cited emission, other objects show both, and still other objectsshow only weak or no emission in CO v = − ∼ L ⊙ , respectively. Table 5 summarizes the properties of Gaussian profiles thatwere fit to the di ff erent isotopic and vibrational lines. The fitswere applied to lines coadded over all unblended rotational lev-els (Fig. 6). The emission from v ′ > − ) and centered at the systemic velocity. Noneof the lines show a double-peaked profile that is characteristic ofKeplerian rotation in a disk with high inclination, although theco-added CO v = − CO emission lines are narrow (FWHM between 10 to 50km s − ), with central velocities shifted between 0 to -15 km s − .The range of velocity centroids for the narrow component sug-gests that our crude classification scheme does not capture thefull range of physical processes that produce this component.In most cases, the CO emission can be clearly distinguised asa broad or narrow component. In a few spectra, the precise clas-sification of CO components is unclear. For example, the narrowcomponent of CO emission from GSS 30 and IRS 44W each in-clude emission from at least two regions based on distinct com-ponents in the line profiles. These cases are discussed in moredetail in the following subsections and in Appendix C.Online Table A.2 presents the line fluxes and equivalentwidths for the di ff erent components in several selected transi-tions. Most CO v = − CO v = − ∼ −
100 K, Jørgensen et al., 2002;Bergin & Tafalla, 2007) and only a ff ects emission in low- J tran-sitions. However, some absorption components are optically-thick for transitions with J >
30. In some cases, multiple COabsorption components are detected. The listed fluxes are mea-sured from Gaussian fits to the emission profile, ignoring anyabsorption component. The FWHM and central velocity are de-termined from fits to the summed profiles (Table 5). The centralvelocities are listed relative to the velocity of CO and C O ab-sorption listed in Table 2. Fitting profiles to the observed emis-sion while ignoring the absorption components yields line fluxesmeasured consistently throughout a given spectrum. In manycases the fit is applied only to a portion of the line profile toavoid absorption features, regions with optically-thick telluricabsorption, contamination from emission in other componentsof the line, and emission in other lines. In each case the 1- σ uncertainty in flux is dominated by the uncertainty in the contin-uum level and does not include the uncertainty in the relative orabsolute flux calibration.In the subsequent subsections, we describe separately theproperties of the broad and narrow emission components,spatially-extended CO emission, H emission, and CO wind ab-sorption detected within our sample. For each CO emission com-ponent, a model of an isothermal, plane-parallel slab of CO gas(see Appendix D) is used to calculate temperatures, column den-sities, and emitting areas, with results in Table 6. Broad emission in v = − (Figure 5). Emission in some v = − − / N. Non-detections of lines from v ′ > Fig. 5.
Line profiles of CO (left) and CO v = − v = − CO line profiles of S CrA S and DR Tau are shown at the top.
Table 4.
Presence of CO emission components a Star L bol Broad Narrow Wind( L ⊙ ) v = − v = − CO CO abs.SVS 20 S 142 n n Y Y nElias 29 41 n n Y Y YL1551 IRS 5 23 n n Y Y nIRS 44 E (18) – n n – nGSS 30 14 n? n? Y Y nHH 100 IRS 15 Y Y Y Y YCrA IRS 2 12 – Y Y Y nIRS 44 W (9) n n Y b Y nHL Tau 6.6 Y Y Y Y nIRS 43 S 6.0 Y Y Y? n nIRS 63 3.3 Y Y Y – YTMC 1A 2.8 Y Y Y Y YWL 12 d e c n YWL 6 2.6 n – n – nRNO 91 2.5 Y Y n n nElias 32 1.1 n n n n yIRS 43 N – n – n – –SVS 20 N (0.27) Y Y Y – n a Ymeans present, n means not present– means not possible to determine b Two distinct narrow components c Weak emission tentatively identified as narrow component d Sep. 2007 e Mar. 2010 generally not significant, assuming the same flux ratio in v = − v = − Counting WL 12 as a detection, despite a non-detection in one ofthe two observations, and GSS 30 as a non-detection, despite some v = − § Fig. 7.
A Venn diagram showing the presence of CO emissionand absorption components from Table 4.In most cases, a component of similar shape and strength asthe v = − v = − v = Fig. 6.
Co-added line profiles from sources with detected CO emission. In each plot the CO v = − CO v = − O v = − CO v = − CO emission line profile to that of CO, C O, or v = −
1, depending on thesource.1 − v = − § v = − v = − v = Table 5.
Average CO Emission Line Properties a Target Lines v bCO (km s − ) FWHM (km s − )Broad ComponentHH 100 IRS c v = − c v = − v = − v = − v = − v = − v = − v = − d
118 (30)WL 12 h v = − v = − v = − v = − v = − v = − d (96) d SVS 20 N v = − v = − d
115 (40)Narrow ComponentHH 100 IRS c CO 0.6 (0.4) 11 (3)HH 100 IRS c v = − CO 2.4 (1.5) 18 (7)Elias 29 e v = − d
19 (3)IRS 63 e v = − CO -15 (6) 42(14)SVS 20 N v = − f CO 0.0 (0.4) 34 (4)CrA IRS 2 f v = − v = − CO (-8.0) d
47 (23)IRS 44 W CO 15.4 (0.4) 14 (1)IRS 44 W C O (15.4) d CO -1.0 (1.0) 32 (4)L1551 IRS 5 CO 0.8 (0.7) 12 (2)L1551 IRS 5 CO (0.8) d
17 (2)L1551 IRS 5 v = − d h v = − v = − h CO, J >
30 -10 (1) 36 (3)GSS 30 SW CO -9.1 (1.0) 6.3 (1.0)GSS 30 NE CO -8.6 (1.0) 6.5 (1.0)IRS 43 S h v = − a Based on Gaussian profiles fit to coadded lines. b Relative to velocity of CO and C O absorption. c From spectrum obtained in August 2008 d Forced value from same component in v = − CO transition e Fit only to red emission f Fit is not good because line is not Gaussian g March 2010 h Tentative classification as narrow component
Table 6.
Physical Properties of CO Emission Region
Star Component T rot (K) log N (CO) a (Area) . b IRS 44 W Narrow c ±
20 19.15 ± .
25 3.6GSS 30 Narrow c ±
15 18 . ± .
10 18.1GSS 30 Extended 250 ±
30 18 . ± . ±
200 – –CrA IRS 2 Narrow 560 ±
70 19 . ± . ±
100 17 . ± . a Column density N in units of log cm − hereand throughout paper, assumes b = . − . b Square root of emitting area, in units of AU c Blueshifted emission − v = − CO lines within our sample.The centroid of the v = − > − (IRS 43 S, HL Tau, and WL 12) each have very low S / Nin the coadded line profiles and as a result have unreliable centralvelocities. In the cross-dispersion direction, the broad emissionfrom HH 100 IRS, RNO 91, and CrA IRS 2 (the three cases withhighest S / N in the broad lines and the highest spatial resolution)is centered at the same location (to within ∼ . . − . / N in broad line emission over a wide range of J levels.The excitation diagram of CO emission from HH 100 IRS showsan upturn in N / (2 J +
1) at low- J , characteristic of optically-thickemission from warm (500-1500 K) gas (see Appendix D). TheCO v = − ff ects. No CO v = − ∼
10 times weaker than the CO v = − v = − ∼ ±
100 K. At this temperature, the lack ofdetectable CO emission in this component requires a columndensity log N (CO) < .
0. In the excitation diagram, the curvefor v = − N (CO) > .
5. The emissionin CO v = − J is underproduced at thistemperature and column density, which may suggest that the ro-tational temperature derived from the v = − v = − ∼ . v = − − ∼ .For CrA IRS 2, the broad component in the v = − v = −
1, 3 −
2, and a few 4 − v = − − ∼ ff erent vibrational levels yields a vibrational temperatureof ∼ Emission in CO v = − CO emission is also seen in CO and C O lines,which indicates that the CO lines are optically-thick. Some ofthese same objects also have a broad emission component. Inonly one case, GSS 30, is CO v = − / N andline-to-continuum contrast in that spectrum.The line profile of narrow emission is centered at the sys-temic velocity for most sources. In addition to emission at thesystemic velocity, IRS 44 and GSS 30 also show narrow com-ponents blueshifted by ∼
10 km s − . Meanwhile, WL 12 (whenthe emission is detected, see variability in Fig. 16) and SVS 20 Sshow only a blueshifted component. In each case, the emissionis relatively narrow, with FWHM between 10–50 km s − . Formost objects, the emission in this component is centered at thesame spatial location on the detector as the continuum emission Fig. 9.
A comparison of the CO P(6) and CO R(6) line pro-files from GSS 30. The two profiles should be similar, includ-ing absorption from the same lower level. The absorption in the CO line is deeper and redder than that seen in the CO line,despite that the CO line cannot be more optically-thick thanthe CO line. Instead, we infer the presence of redshifted ab-sorption from infalling gas at 10–20 km s − . The broader emis-sion is produced far enough from the central source that it doesnot travel through the infalling gas and consequently does notsu ff er from absorption.and is generally not spatially extended. In one case, IRS 44 W,some of the CO is o ff set by 0 . ′′
07 ( ∼ . ′′
22 ( ∼
28 AU). Some veryextended emission is also detected from GSS 30 and IRS 43 andis described in § / N CO spectra.Figure 8 shows the rotational diagram for narrow emissionfrom GSS 30, IRS 44 W, and CrA IRS 2. Relative to the broademission component, the narrow emission component is pro-duced by gas that is cooler, with temperatures of 250–600 K, andmore optically-thick, with column densities of log N (CO) ∼ ff erent processes for dif-ferent objects. For IRS 44 W and GSS 30, the temperature andcolumn density from the rotational diagrams apply only to theblueshifted emission because CO emission is not detected onthe red side of the line profile.Figure 9 compares the CO P(6) line with the CO R(6)line for GSS 30. The two lines should have similar emission andabsorption properties, modulated by the di ff erent optical depths.However, the CO absorption is actually broader on the red sideof the line than the CO absorption. The most likely explanationfor this discrepancy is that the CO emission su ff ers less from COabsorption than does the continuum. The stronger CO emissionthen fills in the absorption more than the weaker CO emission.The absorption, detected out to +
10 km s − , indicates the pres-ence of infalling gas in our line of sight to GSS 30. The M-band spectra of both GSS 30 and IRS 43 show CO emis-sion extended on arcsecond scales. We concentrate here on theextended CO emission from GSS 30 because it is much brighterthan that from IRS 43. The slit position angle for the GSS 30 observation roughly aligns with the ∼ ◦ position angle of theoutflow (e.g. Allen et al., 2002).Figure 10 shows that emission in low- J CO lines is de-tected out to 2 . ′′ CO is also de-tected o ff -source in the SW direction and is barely detected to theNE. Co-adding C O, CO v = −
1, and CO v = − J lineseach yields marginal detections of emission to the SW and non-detections to the NE. Most of the extracted emission is in narrowline profiles (FWHM ∼ − ) centered 2 − − bluewardof the CO absorption. A weak, broad (FWHM ∼
20 km s − ) baseis also detected in the line profile. The CO absorption likely at-tenuates any emission on the red side of the CO line profiles.Figure 11 compares the extent of emission in CO and theM-band continuum to the spatial extent in archival L-band(Duchˆene et al., 2007) and K-band images (3.6 and 2.2 µ m, re-spectively) obtained with VLT / NACO, deconvolved to the ap-proximate spatial resolution of our CRIRES observations. TheK-band continuum and CO line emission is much more spatiallyextended than the L- and M-band continuum emission. The blueK-L color of the nebulosity suggests that the continuum emis-sion is seen in reflected light.Figure 8 shows the rotational diagram for CO line fluxes inthe SW direction, extracted over a region from 0 . ′′
47 to 0 . ′′ ∼
2% the flux of the cen-tral pixel. A linear fit to the CO emission yields a temperatureof 250 ±
30 K. The low CO / CO line flux ratio and the curvedshape of the CO rotational diagram indicates that the COemission is optically-thick, with log N (CO) ∼ .
7. With theseparameters, the C O is just below our detection limit. The linefluxes are reproduced with an emission area of (3.7 AU) , whichis two orders of magnitude smaller than the approximate extrac-tion region of 0 . ′′ × . ′′
3, or ∼ (43 AU) at 120 pc. Lowering theDoppler b parameter from 2.0 to 0.2 km s − decreases the col-umn density to log N (CO) ∼ .
8, thereby increasing the emis-sion area to (11 AU) , which is still ten times smaller than theextraction region. This discrepancy between emitting area andextraction area is somewhat surprising because the CO emissionappears to be smoothly distributed in the slit and the spatially-extended K-band continuum emission is smoothly distributed inthe K-band image. Four objects, WL 6, IRS 44 E, Elias 32, and IRS 43 N, are un-detected in CO emission. For these four objects, the presence ofstrong CO emission relative to the continuum can be definitivelyruled-out in our CRIRES spectra. However, the spectra are in-conclusive regarding the presence of weak CO emission. Eachnon-detection is discussed in detail in the following paragraphs.No CO emission is detected in our CRIRES spectrum of WL6. The telluric absorption falls on the red side of the line. In sev-eral other cases (WL 12, Elias 29, IRS 63), the CO v = − / circumstellarabsorption at low-resolution, leading to a non-detection in theISAAC spectrum (Pontoppidan et al., 2003). The CRIRES spec-trum is consistent with weak CO emission. Fig. 8.
Plots of apparent column density versus excitation for CO emission for the broad and narrow emission components andspatially-extended emission. For the CO v = N (CO) and T , based on models described in § A.4.
Fig. 11. Top:
VLT / NACO K-band (left) and L-band (right, provided by Duchˆene) images of GSS 30, with the CRIRES slit in red.Nebulosity is detected to the NE and SW of the resolved central source.
Bottom:
The cross-dispersion profile of K-band, M-band,and narrow CO emission from GSS 30 along the CRIRES slit. The K-band emission is extracted from the NACO image, deconvolvedto the 0 . ′′ +
50 km s − , with a peak-to-continuum that reaches ∼ .
05. However, this marginal detec-tion is not considered significant because it may be an artifactof the non-standard spectral extraction from this close binary(see § J CO linesin an ISAAC spectrum of Elias 32 (Pontoppidan et al., 2003), soour non-detections of CO emission from Elias 32 on three dif- ferent nights is surprising. Since CRIRES would spectrally re-solve CO emission from the narrow absorption line, the emissioncomponent should be detected at velocities that lack CO absorp-tion. The equivalent width required to explain the CO emissionin the ISAAC spectrum is inconsistent with the non-detection ofCO emission in the CRIRES spectrum of Elias 32. Either theequivalent width of the CO emission is variable or the ISAACspectrum detected mostly extended CO emission. Some CO isdetected in blueshifted absorption to Elias 32 (see § We observed Elias 32 three times, with each spectrum having lowS / N. The highest S / N spectrum for Elias 32 was obtained on 3 May2008, when the telluric CO absorption was located at +
23 km s − . Aswith WL 6, the non-detection of CO emission on that date is not signif-icant. From the spectrum obtained on 10 August 2008, we find a limit12erczeg et al.: CO Emission from Embedded Objects Fig. 10.
Spectra around the CO v = − CO v = − ff -source spectra are ex-tracted from 0 . ′′ − . ′′
93 from the central source. The 2D spec-tral images are obtained from coadding images over CO and CO lines. Both CO and continuum emission are detected o ff -source in both directions along the slit. The emission is strongestto the SW. The o ff -source CO line emission is very narrow(FWHM ∼ . − ) and likely su ff ers from CO absorptionin our line of sight.The object IRS 43 N is a very faint companion to IRS 43S. The non-detection is not significant, although the equivalentwidth must be at least 3 times smaller than the equivalent widthof CO emission seen from IRS 43 S. S(9) Emission
Emission in the H S(9) line at 4.6947 µ m is detected from 9objects in our sample (see Figure 12 and Table 7). Most H lineshave small or negligible blueshifts and a FWHM of ∼ − − . Figure 13 shows the spatial distribution of the H emis-sion in the cross-dispersion direction, along with the spatial dis-tribution of continuum and CO emission. The H emission fromfive stars (IRS 63, Elias 29, SVS 20 S, L1551 IRS 5, and WL 6)is not extended along the slit direction and is consistent with a that the emission on the red side of the central absorption is less than7% above the continuum flux. Weak CO emission from several otherembedded objects would not have been detected with this S / N. Fig. 12. H line profiles from our sample, including non-detections from CrA IRS 2 and HL Tau. The H emission fromSVS 20 S and WL 6 should be considered as tentative detectionsbecause the line is very broad.point source. The H emission is extended in the slit for the other4 objects with detected emission (GSS 30, WL 12, IRS 43, IRS44 W). The observed H emission from IRS 43 S is producedby molecular gas that extends > . ′′ >
190 AU) to the S of thecentral object. The H emission from WL 12 extends 1 . ′′ ∼ emission from IRS 44 is associated with IRS44 W, similar to the case for CO emission. Weak H emission isextended by up to ∼ . ′′ ∼
140 AU) to the E of IRS 44 W. Theequivalent width of H emission is ∼
50% smaller on our nightof good seeing than on our two nights with poor seeing and thanin the ISAAC spectrum (Table E.3). This result confirms that theS(9) emission from IRS 44 is spatially-extended because poorseeing increases the ratio of o ff -slit emission to on-slit emission.Some on-source non-detections are not significant. The CO3–2 R(10) line at 4.6958 µ m could mask any weak on-source H emission from CrA IRS 2 and HH 100 IRS. The non-detectionof H emission from Elias 32 is limited by poor S / N. On theother hand, Bitner et al. (2008) detected S(9) emission from HLTau at a level that should have been detectable in our spectrum . Table 7. H S(9) Line Detections
Target v H a FWHM EW Extended?km s − km s − km s − Detected H linesIRS 44 E + W b -4.6(0.5) 22(3) 9.4(2.0) YIRS 44 E + W c -4.9 (0.9) 23(5) 9.9(2.4) YIRS 44 E + W d -6.6 (0.6) 25(3) 6.7(1.3) YElias 29 -9.5 (2.5) 26 (8) 0.9 (0.2) N e IRS 43 S + N -5.1 (1.6) 18(4) 1.7(0.3) YIRS 63 f -10.2 (3.5) 14(4) 0.7(0.3) NWL 12 f -17(1) 26(3) 16(2) YL1551 IRS 5 -17 (3) 22(5) 5.4(1.4) N e GSS 30 g -7.9 (0.7) 18(3.5) 3.2(0.4) YBroad H lines: spurious detections? h WL 6 1.5(2.0) 89(10) 6.9(1.7) NSVS 20 S 19 (2) 70(14) 5.3 (1.0) NNon-detectionsElias 32 low S / NCrA IRS 2 blended with COHH 100 IRS No detection, telluric correction?HL Tau No detectionRNO 91 No detectionTMC 1A No detection, bad telluric correctionSVS 20 N low S / N a Relative to velocity of CO and C O absorption. b c d e Low FWHM in continuum for CRIRES observation f low S / N g C O line detected at -30 km s − h Uncertain detection because of poor telluric correction
The H S(9) emission reported in Bitner et al. (2008) could bespatially-extended and not located within our slit, similar to thespatial distribution of H µ m emission from HLTau (Takami et al., 2007; Beck et al., 2008). However, in our2010 observation of HL Tau, the slit was aligned with the posi-tion angle of the outflow and the H µ m emission,yet no extended H emission was detectable. The Bitner et al.(2008) spectrum of HL Tau includes a stronger unidentified fea-ture redward of the S(9) line, which may alternately point to aspurious detection in their spectrum.The broad, redshifted H emission from SVS 20 S is not con-firmed in the ISAAC spectrum and may be a spurious detectionintroduced by a poor telluric correction. A strong telluric CO absorption line is located at +
85 km s − of the H line ( + − relative to the source velocity at the time of the obser-vation). The H emission from WL 6 is also particularly broad,with a FWHM of 89 km s − . The ISAAC spectrum includes asimilar detection for WL 6, so this broad emission may be real.The H S(8) line at 5.0529 µ m is included in a wavelengthsetting that was used to observe only GSS 30. Any emission inthis line is severely blended with strong emission in the CO v = − S(8) / S(9) line flux ratio upper limit of ∼ . ∼
700 K. Using a FWHM =
12 km s − , as measured by Bitner et al. (2008), fora Gaussian line centered at the systemic velocity, we measure an equiv-alent width upper limit of ∼ . − , roughly 2.5 times weaker thanthat detected in the TEXES spectrum. Some emission could be locatedat -10 to -30 km s − in the on-source spectrum but not be detectablebecause of a poor telluric correction. Our March 2010 observation ofHL Tau has a poor telluric correction at the H line, presumably be-cause of variable telluric absorption and is not usable for the analysisof on-source emission. No o ff -source H emission was detected in thatobservation. Fig. 13.
Spatial profiles of M-band continuum emission, coadded CO v = − S(9) emission from foursources. The H emission from these four sources is much moreextended than the CO and continuum emission, and is reminis-cent of H v = − emission from Elias 29, L1551 IRS 5, IRS 63,and WL 6 is not extended and not shown here. The CO emissionfrom GSS 30 shown here is the blueshifted component between-20 and -10 km s − . The CO emission on the line wings is notspatially extended beyond the continuum emission. Within our sample, six objects (HH 100 IRS, IRS 63, WL 12,Elias 29, Elias 32, and TMC 1A) show wind absorption in CO v = − ∼
100 km s − .The wind absorption to Elias 32, IRS 63, and WL 12 is slower,with speeds up to ∼
50 km s − . These di ff erences, and the lackof CO wind absorption in the majority of our sample, may beattributable to the di ff erent viewing inclinations, compositions,and wind velocities between sources.The wind absorption is typically weakest at J ′′ = J ′′ =
10 to 20. For HH 100 IRS, the absorption isstrongest in two distinct components at -65 to -85 km s − . Theabsorption seen to TMC 1A is strongest at low velocity and be-comes undetectable at . −
65 km s − . To calculate temperaturesand column densities, we assume that the wind absorption is op-tically thin. An alternate explanation, which applies to the slowwind seen from TMC 1A described below, is that the wind onlyattenuates a fraction of the M-band continuum emission. In thiscase, the absorption in low J and high J lines would still needto be optically thin because the absorption depth is largest forlines with 10 < J <
20. For TMC 1A, an optically-thick, low-velocity ( < − ) component is seen in low- J lines and in CO lines. Since the spectrum does not go to 0 at low velocitiesin the CO lines, the low-velocity absorption must obscure onlyabout half of the M-band continuum emission produced by theinner disk.The temperature of the wind for HH 100 IRS and TMC 1Ais 1260 ±
100 K and 1260 ±
40 K, respectively (Fig. 15). Fig. 15shows the column density of CO in 1 km s − bins in the windsseen to HH 100 IRS and TMC 1A. The log of the total CO col-umn density to TMC 1A, summed from 20–65 km s − , is 19.6,and for HH 100 IRS, summed from 60–90 km s − , is 18.4.Figure 16 shows CO v = − Fig. 15.
Rotational diagram for wind absorption from HH 100 IRS (left) and TMC 1A (middle) covering the -55 to -80 km s − region that includes the double-peaked wind absorption. The temperature of the absorbing gas for both stars is 1260 ±
100 K. Theright panel shows the column density of CO in 1 km s − bins. Fig. 14.
Wind absorption towards four sources in CO 1–1 lines,with high- to low- J from top to bottom.changes between the two observations, with stronger absorptionat ∼ −
70 km s − in August 2008. The shape and strength of theemission line profile remained the same, to within observationaluncertainties. For WL 12, some of the absorbing gas appears tohave gone into emission, which could happen with a change inviewing angle of the wind or if the opacity of the wind decreased,allowing some emission to escape.
4. DISCUSSION In §
3, we describe four di ff erent components of CO rovibra-tional lines from embedded young stars: (1) broad, vibrationally-excited emission ( § § § § ff er-ent components of warm CO gas in the YSO. The broad CO emission seen from the embedded objects in oursample is similar to the warm CO emission seen from CTTSs.Of the 12 CTTSs with detected CO v = − v = − − and the line fluxes range from (0 . − × − erg cm − s − (0 . − × − L ⊙ at 140 pc). Similarly, in our sample most ob-jects with broad CO v = − v = − − and line fluxes that range from ∼ (2 − × − ergcm − s − (2 − × − L ⊙ for a 120 pc distance and extinction A M = COemission was only detected to one source. In our spectra, thebroad component is not detected in CO lines. The measuredtemperatures are roughly consistent with the CO temperaturesin the CTTS sample in Salyk et al. (2009). The centroids of thespectral line profile are consistent with the source velocities, andthe emission is centered on the source and not spatially extendedbeyond ∼
10 AU. The most likely origin for this broad compo-nent of CO emission from embedded sources is the disk.That vibrational levels v = , , ∼ v ≥ v ′ ≥ CO. The lackof broad CO emission is therefore consistent with the presenceof UV-excited gas and places a limit of log N (CO) <
18 in thisinner disk region.Emission from the v ′ = v = − Fig. 16. Left:
Comparing the CO fundamental line profile from HH 100 as seen in observations separated by 11 months. The profilesshown here are the co-added lines of P(4), P(6), P(7), and P(8). No significant di ff erence is seen in the emission line profile betweenthe two epocsh. In August 2008 the wind absorption at -75 km s − was deeper than that seen in September 2007. Right:
Co-addedCO emission from two observations of WL 12. The observation of September 2007 covered low-J lines while the observation ofMarch 2010 covered high-J lines, so the comparison of line profiles is not direct. Some of the wind absorption in 2007 appears togo into emission in 2010.from CTTSs (Najita et al., 2003; Bast et al., 2011) and HerbigAeBe stars (Brittain et al., 2007, 2009; van der Plas et al., 2009).That emission from v ′ > ffi ciently strong to excite warm CO in the inner disk.Figure 17 compares the rotational diagram of number ofCO molecules for HH 100 IRS, DR Tau, and TW Hya. Thefluxes for TW Hya were measured by Salyk et al. (2007). Thefluxes from DR Tau were measured from the CRIRES spec-trum presented in Bast et al. (2011). TW Hya is an ∼ ∼ ∼ . × − M ⊙ yr − (Herczeg & Hillenbrand, 2008). DR Tau is a youngerstar with an accretion rate of 3 × − M ⊙ yr − (Gullbring et al.,2000), comparable to the 10 − M ⊙ yr − estimated for HH 100IRS by Nisini et al. (2005). The strength of CO v = − (Najita et al., 2003;Salyk et al., 2007; Pontoppidan et al., 2008; Salyk et al., 2009;Najita et al., 2009; H¨ugelmeyer et al., 2009; Bast et al., 2011;Brown et al., in prep.). Most CO lines from embedded objectsare consistent with Gaussian profiles. Some Gaussian profilesmay be consistent with a Keplerian disk viewed face-on, as is the case for CO lines from several CTTSs (Pontoppidan et al.,2008; Goto et al., 2011).To investigate the absence of double-peaked line profiles,Bast et al. (2011) selected a subsample of 8 CTTSs with brightCO lines and high accretion rates. In that sample, the combi-nation of the narrow central peak and the lack of any spatially-extended emission, with upper limits of a few AU, indicates thepresence of some gas near the star that is not in Keplerian rota-tion. The subset of embedded objects that have low bolometricluminosities is likely similar to the sample of Bast et al.. In both,the broad line widths indicate an origin near the where the in-ner disk is likely truncated. The fluxes suggest an emitting areaof ∼ . Bast et al. (2011) were unable to explain the linefluxes, profiles, and lack of any spatially-extended emission withan origin in the surface layers of a warm disk in Keplerian rota-tion and with a surface temperature described by a power law.One possible explanation attributes much of the line flux to adisk wind (Pontoppidan et al., 2011).Regardless of explanation, the large velocities in the line pro-file are likely explained by Keplerian broadening. For the diskexplanation, the Keplerian broadening is in situ , while for thedisk wind explanation the Keplerian broadening would applyto the inner radius of the launch region. Within the context ofthe disk / disk wind interpretation, the inner radii R in (sin i ) forthe CO emission are calculated from the velocity of the best-fitGaussian profile at 20% the peak flux (Table 8) and a centralmass 1 M ⊙ . Although a direct measurement of this location inthe line profile is typically not possible in our sample becauseof insu ffi cient S / N, all broad line profiles are consistent with aGaussian profile. About 93% of the warm CO emission is pro-duced beyond this inner radius. Alternately, we solve for inclina-tion by setting R in = .
05 AU for every source. The listed innerradii and inclinations su ff er from significant uncertainties. Najita et al. (2003) suggest that the CO emission from RW Aur isdouble-peaked, but this inference was based on modeling a small spec-tral region that did not show a clear double-peaked line profile. Wefind no evidence for a double-peaked CO emission profile in a CRIRESspectrum of RW Aur A (Brown et al., in preparation).16erczeg et al.: CO Emission from Embedded Objects
Fig. 17.
A comparison of CO emission from the broad compo-nent of HH 100 IRS (black circles), DR Tau (blue diamonds),and TW Hya (red squares). The emission in CO v = − v = − Molecular emission is often detected in winds from young stars(e.g. Shang et al., 2007; Beck et al., 2008; Davis et al., 2010).As described in Panoglou et al. (submitted), winds can have amolecular component for one of the following reasons: (i) en-trainment of molecular gas in the envelope or cloud by the wind,(ii) molecular formation within the wind itself, or (iii) launch-ing of molecular gas in a disk wind. Typically, observations ofmolecules in winds lack su ffi cient spatial resolution to discrim-inate between these possibilities. Our data includes diagnosticsof the wind near the launch region and also of the slow-movingmolecular gas extended from the star. The fast blueshifted CO absorption detected to HH 100 IRS,Elias 29, and TMC 1A suggests that the MHD wind itself isat least partially molecular when it is launched. The slower COabsorption detected to Elias 32, WL 12, and IRS 63 could beexplained by either an MHD disk wind or a slower thermal diskwind. When detected, wind absorption usually occurs close tothe star because the density within the wind decreases with ra-dius squared as the wind expands. Molecular gas from the enve-lope that gets entrained in the outflow is highly unlikely to pro-duce absorption close to the star and at such high velocities. COwind absorption in the M-band has also been detected from high-mass YSOs (Mitchell et al., 1990; Thi et al., 2010). The pres-ence of CO in the wind suggests that the wind from these sourcesfavors launching from the disk rather than the stellar chromo-sphere or corona (Matt & Pudritz, 2005; Ferreira et al., 2006).Although any significant mass loss from the star itself wouldhave to be cooler than 10 K (Matt & Pudritz, 2007), a chromo-spheric wind would likely be mostly neutral. As an alternativeto a wind that is molecular where it is launched, Glassgold et al.(1991) suggest that CO, but not H , may form within the wind For simplicity, the term winds used here is intended to incorporateall disk winds, MHD winds, and jets from a YSO. itself. However, in their models a high CO abundance requiresa mass loss rate > − M ⊙ yr − , which is larger than the massloss rates inferred for typical Stage 1 sources (Bontemps et al.,1996).Models of MHD disk winds by Panoglou et al. (submitted)demonstrate that molecules can survive within the wind whenthe mass loss / accretion rate are su ffi ciently large ( M acc ∼ − M ⊙ yr − ), in material with an ionization fraction that is highenough to couple the molecular gas to the magnetic field. Themeasured temperature of 1260 ±
100 K is roughly consistentwith the predicted wind temperatures for objects with accre-tion rates of 10 − − − ˙ M yr − . Wind absorption in CO hasnot been detected previously to CTTSs. Several CTTSs showon-source FUV H emission with velocities up to 50 − − (Herczeg et al., 2006). The frequency of CO absorptionin winds from embedded objects and deficiency of CO absorp-tion in winds from CTTSs is consistent with the survival of COrequiring large outflow rates to shield the outflow from irradia-tion by the central star.That CO wind absorption is detected toward only some ofthe objects in our sample is consistent with the MHD windbeing non-spherical and with a relatively wide opening anglenear the star. In a survey of the He I λ ◦ (Close et al., 1997; Lucas et al., 2004) , alsocontrasts sharply with the deep, fast He I absorption seen to thestar (Edwards et al., 2006). For both CrA IRS 2 and HL Tau, thelikely interpretation is that the molecular fraction is low in thewind. For HL Tau, the discrepancy between the deep He I windabsorption and the lack of any detectable CO wind absorptioncould also be explained if 4.8 µ m continuum is seen through avery di ff erent line-of-sight than the 1.1 µ m continuum emission. Powerful jets and winds from young stars can carve out cavitieswithin the circumstellar envelope (e.g., Whitney & Hartmann,1993; Wood et al., 2001; Ybarra et al., 2006). The interaction re-gion between the cavity and the jet / wind entrains some cold gas,which is likely seen in bipolar molecular outflows with veloci-ties of a few km s − . At the surface of the cavity wall, gas canbe heated by shocks and UV photons from the central star (e.g.Spaans et al., 1995; van Kempen et al., 2009).Within our sample, H emission is spatially-extended, always asymetrically about the star. The spectral lineprofiles are typically centered within 10 km s − of the systemicvelocity. The extended component is likely produced in the en-velope or nearby cloud material that is shocked by powerful out- (Furlan et al., 2008) suggest an i ∼ ◦ for the HL Tau disk frommodelling the SED, but we consider the near-IR polarimetry a morereliable method to measure inclination because inclination is degeneratewith many other parameters in broadband SED fitting and because thecombination of high extinction and lack of ice absorption suggests diskattenuation rather than envelope extinction. 17erczeg et al.: CO Emission from Embedded Objects flows. Greene et al. (2010) detected H ′′ scales. Compared with the H − ) and central velocities that di ff er by ∼ −
10 km s − .Of the objects with extended H emission, the slit PA waswell-aligned (better than 20 ◦ ) with the outflow axis only forGSS 30. The spatial distribution is likely analogous to thedistribution of warm H emission seen towards other embed-ded YSOs and environmentally-young CTTSs that drive power-ful outflows (e.g. McCaughrean & Mac Low, 1997; Davis et al.,2001; Walter et al., 2003; Beck et al., 2008; Neufeld et al., 2008;Lahuis et al., 2010; Greene et al., 2010). The spatially unre-solved H emission could be produced by a disk, as is thecase for weak H rovibrational emission from a few CTTSs(Bary et al., 2008). However, because the H emission of-ten includes a significant contribution from surrounding enve-lope / cloud material, a careful analysis would be required to usethe line as a disk probe. In light of our results, the origin of H emission from GSS 30, and perhaps Elias 29 and HL Tau, byBitner et al. (2008) should be considered the wind / envelope in-teraction region rather than the disk.Spatially-extended CO emission is also detected from twoobjects, GSS 30 and IRS 43, that show extended H emission.The extended CO line emission traces cooler gas and is spec-trally more narrow than the H line emission. That extended COemission is not detected more frequently is likely explained bydensities that are much lower than the critical density of 5 × cm − required to populate the v = / H abundance ratio could be negligible. On the other hand,the presence of extended CO emission from two objects may in-dicate high densities in the associated outflows. Narrow CO emission is detected from 9 of 18 embedded ob-jects within our sample. The optical depth of this componentis larger than that in the broad component, as evidenced by thehigh CO / CO line ratios. Narrow line profiles have a widerange of properties, indicating that this classification is overlybroad. For GSS 30, IRS 44 W, WL 12, and SVS 20 S, some orall of the narrow CO emission is blueshifted by ∼
10 km s − .For other objects the emission centroid is consistent with thesystemic velocity. The lack of spatially-extended emission indi-cates that this emission is produced relatively close to the centralstar, although some narrow emission from IRS 44 W is slightlyextended, and CO emission from GSS 30 is likely emitted be-yond infalling CO absorption. The di ff erent velocities of the nar-row component indicates the emission is produced by di ff erentphysical processes, depending on the star.This narrow component does not have a counterpart for thetypical CTTSs observed by Najita et al. (2003) and Brown et al.(in prep.). The mechanism that produces an optically-thick col-umn of ∼
400 K gas for these evolutionarily young objects doesnot heat a similar amount of gas in more evolved CTTSs. Onthe other hand, CO emission is commonly detected within theBast et al. (2011) sample of high accretion rate CTTS. In par-ticular, the CO emission from S CrA S is remarkably similar tothat seen from HH 100 IRS, IRS 63, and TMC 1A. All four ob-jects show broad emission in CO v = − CO transitions, and optically-thick CO v = − Table 8.
Inner Radius of disk CO Emission a Target FWHM a v binner R in (sin i ) c incl. d km s − km s − (AU) ◦ Broad EmissionIRS 63 92 70 0.17 33HH 100 IRS 80 61 0.23 28IRS 43 S 146 111 0.070 58CrA IRS 2 54 41 0.49 19WL 12 98 74 0.16 34HL Tau 130 99 0.088 49TMC 1A 96 73 0.16 34SVS 20 N 100 76 0.15 35RNO 91 165 125 0.056 71Narrow Emission e GSS 30 42 32 0.87 –HH 100 IRS 11 8.4 13 –Elias 29 18 13 5.3 –IRS 63 28 21 2.0 –CrA IRS 2 34 26 1.3 –IRS 44 W 32 24 1.5 –IRS 43 S 42 32 0.87 –L1551 IRS 5 12 11 7.3 – a CO v = − b half-width at 20% the peak of a Gaussian with listed FWHM. c Assumes M ∗ = M ⊙ d Inclination if the inner radius of CO emission is 0.050 AU E Narrow component may not be in Keplerian rotation accretion rates are likely not much more evolved than many ofthe embedded sources in our sample, including HH 100 IRS,TMC 1A, and HL Tau. Four Stage I YSOs with high luminosi-ties show some blueshifted CO emission, which indicates thatthe mechanism that produces the warm CO continues to changewith increasing bolometric luminosity.The component could be explained by a thick layer of gasthat is either in a thermal disk wind and / or the disk surface.Evaporative winds typically produce emission lines with cen-troids at 0 −
10 km s − (Alexander, 2008; Owen et al., 2010).Di ff erent viewing anlges could explain the range of measuredvelocity centroids. The lack of significant emission in lines from v ′ > Most of the objects with large bolometric luminosities emit inthe narrow component but not in the broad component (seeTable 4 and Fig. 18). The sources with small bolometric lumi-nosities typically emit in the broad component. A few objectswith intermediate bolometric luminosities emit in both compo-nents. Within these categories, the correlation between bolomet-ric luminosity and the strength of the broad component ( v = − CO line flux).These comparisons are not definitive because the sample size issmall and the bolometric luminosities have large uncertainties.Other YSO properties must also determine the strength of theemission, as IRS 44 E lacks strong CO emission despite a highbolometric luminosity. However, bolometric luminosity clearlya ff ects what regions emit in CO fundamental transitions.For CTTSs and low-luminosity embedded objects, the broad,vibrationally-excited emission indicates that molecules survivein the inner disk near the disk truncation radius. On the other Fig. 18.
The luminosity in the CO v = − v = − CO and v = − − in the narrow component and 100 km s − in the broad component.hand, the non-detection of this broad emission from someembedded objects indicates an absence of warm CO gas inKeplerian rotation near the star. Of the high-luminosity objectsin our sample that lack broad CO emission lines, disks havebeen detected around GSS 30 and L1551 IRS 5 in the sub-mm(Looney et al., 1997; Jørgensen et al., 2009) and around GSS 30,Elias 29, and IRS 44 E in near-IR polarimetry (Beckford et al.,2008). The measured disk inclinations of Elias 29 ( ∼ ◦ ) andGSS 30 ( ∼ ◦ ) from Beckford et al. (2008) rule out that thenarrow emission simply results from disks seen face-on.For YSOs with high bolometric luminosity, molecules maybe photodissociated out to larger radii than is typical for CTTSs.As a consequence, perhaps no warm, broad CO emission isdetected because the inner disk is mostly atomic or ionized.Similarly, Herbig AeBe stars frequently emit in CO fundamentallines at large radii (e.g. Brittain et al., 2009; van der Plas et al.,2009) but in tracers of warmer gas, such as [O I] λ ff ers from that ofmature disks. For stars with blueshifted CO emission, the regionwhere the CO emission is produced may be active or turbulentenough so that the disk surface evaporates. The deficiency ofstrong, UV-excited lines from v ′ > / or younger) are probably quite di ff erent.Any CO gas in the inner disks of these objects should have beendetected. If the inner disk has not yet settled into Keplerian ro-tation, then CO gas in the inner disk, if present, could produce the narrower, optically-thick emission. Alternately, dust that isoptically-thick at 4.8 µ m could shield any warm CO emission,particularly if the inner disk does not have a temperature inver-sion because of significant viscous heating.
5. CONCLUSIONS
We present an overview of fundamental CO emission in high-resolution
VLT / CRIRES spectra of 18 embedded young stars,obtained to explore the warm inner disks as they are still form-ing. CO emission was detected from 14 of the 18 stars. The fournon-detections are only significant in ruling out large equivalentwidths in the lines. We find the following results:1. On-source CO emission is produced in narrow and broadcomponents, both of which are single-peaked. Often only one ofthese components is detected from a given source. The narrowemission component is typically detected from embedded ob-jects with high bolometric luminosity, while the broad emissioncomponent is typically detected from objects with low bolomet-ric luminosity.2. The broad emission (FWHM range of 50–160 km s − ) isdetected in lines v = − v = − / N is suf-ficient, lines with v ′ >
2. This emission traces warm ( ∼ ∼ and is centered at the systemic veloc-ity. This emission is reminiscent of the warm, single-peaked COemission detected from more mature disks around CTTSs withhigh accretion rates, and likely traces gas in the inner regions ofthe disk or in a slow, thermal disk wind. The vibrational excita-tion is higher than the rotational excitation, which suggests thatUV photoexcitation of the disk surface may be responsible foremission from levels with v ′ > − ) is de-tected in CO, CO and, when S / N is su ffi cient, C O and CO v = − − , indicating an origin in a slow out-flow. Other lines are at the systemic velocity, either because thegas is not escaping the disk or because of viewing angle of thewind. The range of line profiles indicate that this classificationincludes multiple physical processes. The faint emission in CO v = − > L ⊙ . Either warm CO gas is not presentin the inner disk of these objects due to photodissociation, anywarm CO is shielded by dust that is optically-thick at 4.8 µ m, orthe disk has not yet settled into Keplerian rotation. If the molec-ular fraction is low near the star, the gaseous disk should be de-tectable in atomic or ionized gas lines.5. Very narrow, spatially extended CO emission from GSS30 and IRS 43 likely traces cold molecular gas, perhaps in denseregions where the winds interact with the envelope and nearbymolecular material. The extended CO emission is especiallybright from GSS 30, indicating some unique physics or mor-phology for that source which is not typical of the rest of oursample.6. Winds are detected in CO absorption towards 6 ob-jects, with velocities of 10 to 100 km s − and temperatures of1260 ±
100 K. Because wind absorption usually occurs closeto the star, these detections indicate that the wind is partiallymolecular when it is launched. A molecular wind is consistentwith a wind launched from the disk but not with an accretion-driven wind launched from the stellar chromosphere or corona.
The fraction of objects with CO detected in the wind (6 of 18)suggests that the wind may have a wide opening angle near thestar.7. H S(9) emission is commonly detected in our sample. Thespatial extents of ∼ − ′′ seen in about half of these detectionsindicate that the H emission is produced in winds interactingwith surrounding material rather than the disk.We have analyzed M-band CO emission as a probe of gasin the inner few AU of disks around young stars that are still inthe embedded phase of pre-main sequence evolution. The dataquality in these observations, obtained with a sensitive, high-resolution instrument on an 8m telescope on nights with goodseeing, will be di ffi cult to improve upon with existing instru-mentation, although larger samples and spectral imaging withcurrent capabilities could clarify the interpretation of narrow COemission and spatially-extended CO emission. Spectral imagingwith future instruments, such as NIRSpec on JWST or METISon the Extremely Large Telescope , holds significant potentialfor understanding the morphology and production mechanismof CO emission around GSS 30 and other sources, while spec-troastrometry of embedded objects obtained with a laser-guidedAO system could identify the location of the di ff erent CO emis-sion components, following Pontoppidan et al. (2008). Despiteits revolutionary capabilities in observations of young disks, ALMA will be unable to trace warm gas within a few AU of thecentral star, where terrestrial planets are thought to form.
6. Acknowledgements
GJH is indebted to Sylvie Cabrit for a careful read and de-tailed comments and discussion. GJH also thanks Tom Greeneand Mary Barsony for valuable discussion of K-band spectra ofembedded objects, Richard Alexander and Sean Matt for a dis-cussion of the temperature of stellar winds, Gaspard Duchenefor providing the VLT / NACO L-band image of GSS 30, AdwinBoogert and Isa Oliveira for reducing and providing a
Spitzer -IRS spectrum of RNO 91, which we used to estimate the fluxat 4.8 µ m. We appreciate a useful and prompt report fromthe anonymous referee and helpful comments from the editor,Malcolm Walmsley. We also thank Jeanette Bast, Bill Dent,Geo ff Blake, Wing-Fai Thi, Alain Smette, and Ulli K¨aufl for helpin carrying out the observations. Astrochemistry in Leiden issupported by a Spinoza grant from the Netherlands Organizationof Scientific Research (NWO). The authors wish to recognizeand acknowledge the very significant cultural role and reverencethat the Theresenwiese has always had within the Max Planckcommunity.
Appendix A: Resolved Binaries
In our long-slit spectra, spatial information is obtained in thecross-dispersion direction. Most observations are consistent withpoint-source emission. The objects IRS 44, IRS 43, and SVS 20are known binaries that were resolved in our data and are de-scribed below. Interferometry of L1551 IRS 5 at 7-mm indicatestwo distinct sources of dust continuum emission separated by ∼ . ′′ . ′′ A.1. IRS44
IRS 44 was definitively identified as a close binary in L-band AOimaging (Duchˆene et al., 2007), which revealed two componentswith a flux ratio close to 1.0 and a separation of 0 . ′′
30 at a PAof 87 ◦ . High-resolution K-band images also suggested that IRS44 is a binary with a similar separation and a flux ratio of ∼ / N and uncertain detections ofthe secondary.Of our five CRIRES observations of IRS 44, the binary wasresolved only on the nights of 6–7 Aug. 2008, when the see-ing was superb ( ∼ . ′′ . ′′ ∼ ◦ o ff set in po-sition angle between the slit and the binary. Assuming the binaryPA from Duchˆene et al. (2007), then IRS 44 E is ∼ . ± . µ m ice band, absorption reduces the con-tinuum flux from IRS 44 E by ∼
18% and from IRS 44 W by ∼ A.2. IRS43
IRS 43 was resolved as a 0 . ′′
57 binary with an L-band flux ratioof 3.1 mag and a PA of 336 ◦ by Duchˆene et al. (2007). In ourM-band long-slit spectra, the continuum emission from IRS 43is concentrated on the primary. Some extended emission is de-tected to the N, in the direction of the faint secondary. If thisemission is attributed to the secondary component, then it is ∼ . ± . µ m CO ice band absorbs 80% of the photons emit-ted from both the continuum and from the secondary (or ex-tended nebulosity) to the N. We infer that the ice absorptionlikely occurs in a cloud that envelops both objects. A.3. SVS20
SVS 20 is a 1 . ′′
51 binary with a PA of 9.9 ◦ (Eiroa et al., 1987;Haisch et al., 2002, 2006). The separation on our slit is ∼ . ′′ ff erence of 1.25, similarto the magnitude di ff erence seen at other optical and IR wave-lengths. The luminosity ratio of ∼
500 listed in Table 2 is likelymuch too high.The properties of the stellar components of SVS 20 are un-certain. SVS 20 S is considered the primary component be-cause it is brighter in the near-IR and mid-IR. Weak K-bandabsorption features indicate that the primary is likely an early-G star (Doppmann et al., 2005), although as with IRS 63, thisspectral type is tentative because the photospheric velocity isdiscrepant by 10 km s − from the v lsr of the HCO + emission(Gregersen et al., 2000). Oliveira et al. (2009) found that one ofthe components is an M4 star with A V = from both objects. The low extinction suggests that whichevercomponent is the M-dwarf is not embedded in the envelope andperhaps dominates the optical flux as a result. The CO ab-sorption lines are also more optically thick to SVS 20 S thanto SVS 20 N. On the other hand, the CO ice absorption has alarger opacity, τ ∼
1, to SVS20 S than to SVS 20 N (see alsoPontoppidan et al., 2003). The brightness di ff erence between thetwo components, ∼ . ∼ . µ m (Ciardi et al., 2005; Haisch et al., 2006), also does not sug-gest a large di ff erence in extinction or bolometric luminosity be-tween the two stars. For the purposes of this paper, we assumethat both objects are at a similar evolutionary state and embeddedin an envelope. A third, faint component in the SVS 20 systemis separated from SVS 20 S by 0 . ′′ Appendix B: Variable CO Ice Absorption from WL 6
For WL 6, the depth of CO ice absorption decreased from τ = . τ = .
51 in the CRIRESspectrum. From Gaussian fits to the absorption band, the centralwavelength and FWHM remained similar, indicating no signifi-cant change in the ice composition. No other significant changeswere detectable between the two observations.WL 6 is a point-source in K- and L-band AOimaging (Ratzka et al., 2005; Duchˆene et al., 2007).Alves de Oliveira & Casali (2008) find variability of ∼ . − . Appendix C: Notes on CO Emission from IndividualTargets In §
3, we discussed the CO emission from individual stars inmostly generic terms. However, the interpretation of the narrowemission component from several objects is somewhat compli-cated. In the following subsections, we describe details of thenarrow component from three objects, GSS 30, IRS 44 W, andCrA IRS 2.
C.1. GSS30
Our CRIRES spectrum of GSS 30 covers 4.645–4.768 µ m and5.036–5.158 µ m, with a large gap that excludes CO lines with J ′ = −
34 from our spectrum. Figure C.1 compares the COline profile for low- J lines ( J ′ <
10) to the scaled line profiles of CO with high- J (35 < J ′ < CO, C O, and CO v = − CO low- J lines are characterized by red- and blue-shiftedemission with absorption at the cloud velocity. No absorption isdetected in lines with J >
34. C O is detected in absorption butnot in emission.The v = − CO, C O, CO with high − J , and CO with low- J , though interveningabsorption complicates the analysis. The broader component isdominant in the high- J CO lines. The Gaussian profiles com-pared with the other lines in Figure C.1 are kept in the sameflux ratio as measured from the v = − Fig. C.1.
A comparison of CO line profiles from GSS 30, as ex-tracted on-source. Two Gaussian profiles are fit to the emissionin coadded v = − CO v = − J , with low- J , CO, and C O. The two components in the v = − J line the broader component dominates. Some narrowextended emission can also be seen as excess emission in low- J CO and CO lines at -5 km s − . Both of these Gaussian com-ponents are considered narrow, optically-thick absorption withinour crude classification scheme.the absorption, some minor di ff erences can be seen between theGaussian fits and the line emission. Very narrow emission in CO and CO transitions extends o ff -source in both the NEand SW within in the slit and is discussed in § v = − CO / CO lines is ∼
8, which in-dicates that the CO lines are optically-thick. The flux ratio of CO / C O lines of ∼ −
10 is similar to the abundance ratio of8.1 in the local ISM (Wilson, 1999) and suggests that both the Not to be confused with the broad component discussed in § CO and C O emission are optically-thin. The measured tem-peratures for CO and C O are 315 and 340 K. The CO / COline flux ratios yield a total CO column density log N (CO) = b ∼ . − and T =
340 K, assuming that the blueshifted CO emission is dominated by this component and not thebroader emission component. At this column density, the ob-served CO emission lines are nearly optically-thick, but notsu ffi ciently enough to reduce the CO / C O flux ratio. Thefluxes in C O v = − CO v = − ff erfrom any opacity e ff ects until log N (CO) > .
5. At 340 K, the v = − ∼
450 K) at this high column density wouldsu ffi ciently populate the v = v = − J lines of CO and CO is also not producedat such cool temperatures. The total emitting area is roughly (22AU) for log N (CO) = . + log(cos i ). At a distance of 120pc, this spatial extent should be marginally resolvable. The ob-served CO emission is unresolved to a FWHM ∼ . ′′
08 (10 AU at120 pc). The di ff erence could be reconciled with a larger b -valueor a large viewing angle for our line-of-sight to the CO slab. Onthe red side of the line, the column density is di ffi cult to assessbecause the non-detection of CO emission is complicated bythe red absorption.When analyzing spatially-extended CO emission from GSS30 observed with ISAAC ( R = , / disk interface at 10–50 AU from the central star. Thisemission could then be reflected o ff the outflow cavity walls.This shock is expected to have a temperature plateau at ∼ O lines (Neufeld & Hollenbach, 1994). Inthis scenario, the on-source and o ff -source line profiles shouldbe similar, but Pontoppidan et al. (2002) was unable to test thisprediction because the CO lines were unresolved with ISAAC.With the much higher resolution of CRIRES, we find that theo ff -source CO spectral line profiles are significantly narrowerand have a higher peak-to-continuum ratio than the on-sourceCO line profiles. Although the extended nebulosity is bright inthe K-band continuum, K-band pumping into v = v = − ff -source.An alternate but contrived explanation to reconcile the discrep-ancy in emission area versus extraction area is to invoke a veryclumpy medium. Although the emission appears to be smoothlydistributed within the slit, the spatial resolution may not be suf-ficient to detect many di ff erent clumps. C.2. IRS44W
The CO emission line profiles from IRS 44 W are similar tothose from GSS 30. Figure 6 shows that CO emission is seenon both sides of the CO absorption. However, the CO and C Oemission is seen only shortwards of the CO absorption. Whenthe CO emission line profile is scaled to the CO profile (seeinset in Fig. 6), a deficit in flux is seen in the CO emission at < −
30 km s − , at − − (between the line peak and deepabsorption), and on the entire red side of the line profile.The spatial profile of emission on the red wing of CO linesalso di ff ers from the spatial profile of the blue wing of COlines and of the CO lines. The emission on the red wing of CO lines is consistent with the location of continuum emis- sion from the secondary star. However, the CO emission is lo-cated 0 . ′′
07 (8 AU at 120 pc) W of the continuum emission fromthe secondary star with a FWHM of ∼ . ′′
22 (26 AU at 120 pc).The blueshifted component is slightly stronger on the nights withpoor seeing. The unresolved line equivalent widths are strongerin the ISAAC spectrum, which used a wider aperture and wasobtained in worse seeing. These results both support the pres-ence of spatially-extended CO emission. The spectral and spatialinformation requires two physically distinct components for COemission from IRS 44 W.For the blueshifted component, fits to the CO and C Oline both yield temperatures of ∼
330 K. The CO lines cannotbe too optically thick, which places an upper limit on the columndensity of log N (CO) . .
4. The CO lines are very optically-thick, with fluxes that should be considered upper limits, whichplaces a lower limit on the column density of log N (CO) & . ∼ . N (CO)and T because CO and CO v = − C.3. CrA IRS2
CrA IRS 2 is classified here as an embedded object based on theSED (Nutter et al., 2005), although insu ffi cient evidence existsin the literature to confirm the presence of an envelope.The bright CO and CO emission lines from CrA IRS 2have non-Gaussian profiles with centroids that are redshifted by ∼ − from the systemic velocity systemic and, when fitwith a Gaussian profile, have FWHM of ∼
26 km s − . The v = − § CO line pro-file to every CO and CO line in the spectrum. The median CO line profile was calculated by coadding all CO lines thatare not a ff ected by absorption. The resulting fits are typicallygood, except that high- J CO lines have broader peaks. Thelack of any C O emission, with an upper limit of 15% of theflux in CO lines, indicates that the CO lines are not optically-thick. The CO fluxes in the rotational diagram yield a best-fittemperature of T =
560 K. At this temperature, the CO / COline ratios indicate a column density log N (CO) ∼ O emission would be marginally detected.The total emitting area is (1.0 AU) . Appendix D: Synthetic CO Line Fluxes
For optically-thin gas, the temperature and number of COmolecules can be directly measured from an excitation dia-gram. Transitions become optically thick as the column densityof the medium increases, thereby changing the line flux ratiosand, as a consequence, direct temperature measurements. ForCO emission from disks around CTTSs and Herbig AeBe stars,CO v = − CO lines are detected, such excitation dia-grams can be explained by emission produced in optically-thickgas or in gas with a large temperature gradient. The strengths of CO and CO v = − v = − We model the CO line fluxes by calculating the emissionexpected from an isothermal, 1D slab of pure CO gas. The vi-brational excitation and rotational excitation of the CO gas aredescribed by the same temperature. The molecular data was ob-tained from Chandra et al. (1996). Throughout the layer, absorp-tion occurs in a Voigt profile with a Doppler broadening pa-rameter b , which includes thermal and turbulent broadening.Emission lines have a Gaussian profile with a FWHM of 1 . × b .The density is assumed to be high enough so that collisions dom-inate the excitation, allowing the gas to maintain local thermalequilibrium. The layer is dust-free. The fraction of photons thatescape from the total layer is calculated as a function of columndensity in each transition.The primary benefits of an analysis of CO line opacities aremeasuring accurate temperatures, approximate CO column den-sities, and rough surface areas for the emission regions. Withinthis paper, the Doppler parameter is set to b = . − . Alower b parameter yields a faster increase in opacity, so resultingcolumn densities would be smaller. The column densities in thissimple 1D slab model correspond to the amount of material ata specific temperature in our line of sight. For a slab model, thevertical column density would be log N = log N mod + log cos i ,where i is the incidence angle along our line of sight into the slaband N mod is the column density of the model Characterizing thisgeometrical complication is beyond the scope of the simplisticapproach adopted here.Figure D.1 shows the excitation diagram, normalized to N J (2 J + A ul of the R(0) line, for a range of total CO column den-sities, T = b = . − . Each individual tran-sition starts to become optically thick at log N J ∼
14 and iscompletely opaque at log N J >
18. At 1000 K, the rotationalpopulation peaks at J =
13, so that transitions with J ′′ ∼ J line fluxes are weaker than expected, relative to low- J andhigh- J lines. The opacity in the R-branch transitions increasesfaster than the opacity in P-branch transitions because the lowerlevels of R-branch transitions are more populated than thosein P-branch transitions. A more extreme example of the diver-gent P- and R-branch lines occurs in the Orion BN / KL region(Gonzalez-Alfonso et al., 2002).For a slab with T = J ′ <
30 or the combined set of lines with J ′ < < J ′ <
30, are .
200 K lower than the input model temperatures.Temperature fits to only the high- J (15 < J <
30) lines aretempting but can be misleading, with best fits as much as ∼ J <
10 can yield temperatures of <
200 K and should becompletely avoided.
Appendix E: Online only TablesReferences
Acke, B., van den Ancker, M.E., & Dullemond, C.P. 2005, A&A, 436, 209Alexander, R.D., Clarke, C.J., & Pringle, J.E. 2006, MNRAS, 369, 216Alexander, R.D. 2008, MNRAS, 391, L64Allen, L.E., Myers, P.C, Di Francesco, J., Mathieu, R., Chen, H., & Young, E.2002, ApJ, 566, 993Alves de Oliveira, C., & Casali, M. 2008, A&A, 485, 155Andr´e, P., Ward-Thompson, D., & Barsony, M. 1993, ApJ, 406, 122Andrews, S.M., & Williams, J.P. 2007, ApJ, 671, 1800Arce, H.G., Shepherd, D., Gueth, F., Lee, C.-F., Bachiller, R., Rosen, A., &Beuther, H. 2007, Protostars and Planets V, B. Reipurth, D. Jewitt, and K.Keil (eds.), University of Arizona Press, Tucson, 951 pp., 2007., p.245-260Arce, H.G., & Sargent, A.I. 2006, ApJ, 646, 1070
Fig. D.1.
Synthetic CO v = − T = N (CO) = T = J ′′ ∼
10 are the first to become optically-thick, thereby reducing theflux in mid- J lines relative to those with J < J lines. Bary, J.S., Weintraub, D.A., Shukla, S.J., Leisenring, J.M., & Kastner, J.H. 2008,ApJ, 678, 1088Bast et al., 2011, submittedBate, M.R. 2010, MNRAS, 404, L79Beck, T.L., McGregor, P.J., Takami, M., & Pyo, T.-S. 2008, ApJ, 676, 472Beckford, A.F., Lucas, P.W., Chrysostomou, A.C., & Gledhill, T.M. 2008,MNRAS, 384, 907Beckwith, S.V.W., Sargent, A.I., Koresko, C.D., & Weintraub, D.A. 1989, ApJ,343, 393Bergin, E.A., & Tafalla, M. 2007, ARA&A, 45, 339Bitner, M.A., et al. 2008, ApJ, 688, 1326Blake, G.A., & Boogert, A.C.A. 2004, ApJ, 606, L73Bontemps, S., Andr´e, P., Terebey, S. & Cabrit, S. 1996, A&A, 311 858Boogert, A.C.A., et al. 2008, ApJ, 678, 985Brinch, C., Crapsi, A., Jørgensen, J.K., Hogerheijde, M.R., & Hill, T. 2007,A&A, 475, 915Brinch, C., Jørgensen, J.K., & Hogerheijde, M.R. 2009, A&A, 502, 199Brittain, S.D., Rettig, T.W., Simon, T., Kulesa, C., DiSanti, M.A., & Dello Russo,N. 2003, ApJ, 588, 535Brittain, S.D., Rettig, T.W., Simon, T., & Kulesa, C. 2005, ApJ, 626, 283Brittain, S.D., Simon, T., Najita, J.R., & Rettig, T.W. 2007, ApJ, 659, 685Brittain, S.D., Najita, J.R., & Carr, J.S. 2009, ApJ, 702, 85Brown, D.W., Chandler, C.J., Carlstrom, J.E., Hills, R.E., Lay, O.P., Matthews,B.C., Richer, J.S., & Wilson, C.D. 2000, MNRAS, 319, 154Brown, J.M., et al., in prepCabrit, S., Guilloteau, S., Andre´e, P., Bertout, C., Montmerle, T., & Schuster, K.1996, A&A, 305, 527Calvet, N., D’Alessio, P., Hartmann, L., Wilner, D., Walsh, A., & Sitko, M. 2002,ApJ, 568, 1008Carr, J.S., Tokunaga, A.T., Najita, J., Shu, F.H., & Glassgold, A.E. 1993, ApJ,411, L37Carr, J.S., Tokunaga, A.T., & Najita, J. 2004, ApJ, 603, 213Chandler, C.J., Carlstrom, J.E., Scoville, N.Z., Dent, W.R.F., & Geballe, T.R.1993, ApJ, 412, L71Chandler, C.J., Terebey, S., Barsony, M., Moore, T.J.T., Gautier, T.N. 1996, ApJ,471, 308Chandra, S., Maheshwari, V.U., & Sharma, A.K. 1996, A&AS, 117, 557Chapman, N.L., Mundy, L.G., Lai, S.-P., & Evans, N.J. 2009, ApJ, 690, 496Chen, J.-H., Evans, N.J., Lee, J.-E., & Bourke, T.L. 2009, ApJ, 705, 1160Chrysostomou, A., Clark, S.G., Hough, J.H., Gledhill, T.M., McCall, A., &Tamura, M. 1996, MNRAS, 278, 449Ciardi, D.R., Telesco, C.M., Packham, C., G´omez Martin, C., Radomski, J.T., deBuizer, J.M., Phillips, C.J., & Harker, D.E. 2005, ApJ, 629, 897Close, L. M., Roddier, F., Northcott, M. J., Roddier, C., & Graves, J. E. 1997,ApJ, 478, 766Crapsi, A., van Dishoeck, E.F., Hogerheijde, M.R., Pontoppidan, K.M., &Dullemond, C.P. 2008, A&A, 486, 245
Table E.1.
ONLINE ONLY: Wavelength Settings
Setting Wavelengths (nm) J ′ range a Chip 1 Chip 2 Chip 3 Chip 44710 4639–4663 4670–4693 4700–4722 4728–4749 0–84716 4645–4669 4676–4699 4706–4728 4734–4755 0–94730 4660–4684 4690–4713 4720–4742 4747–4768 0–104770 4702–4725 4731–4754 4760–4781 4787–4807 4–144780 4712–4735 4742–4764 4770–4791 4796–4816 5–154831 4768–4789 4796–4817 4822–4842 4847–4866 11–204833 4770–4791 4798–4819 4824–4844 4849–4868 11–204868 4806–4827 4833–4853 4858–4877 4883–4901 15–234929 4841–4871 4879–4908 4916–4944 4951–4978 18–294946 4858–4888 4896–4924 4933–4960 4968–4994 20–315115 5036–5063 5070–5096 5103–5128 5135–5158 35–42 a Range of J ′ values for CO v = − Davis, C.J., Hodapp, K.W., & Desroches, L. 2001, A&A, 377, 285Davis, C.J., Gell, R., Khanzadyan, T., Smith, M.D., & Jenness, T. 2010, A&A,511, 24de Zeeuw, P.T., Hoogerwerf, R., & de Bruijne, J.H.J. 1999, AJ, 117, 354Doppmann, G.W., Ja ff e, D.T., & White, R.J. 2003, AJ, 126, 3043Doppmann, G.W., Greene, T.P., Covey, K.R., & Lada, C.J. 2005, AJ, 130, 1145Duchˆene, G., Bontemps, S., Bouvier, J., Andr´e, P., Djupvik, A.A., & Ghez, A.M.2007, A&A, 476, 229Dullemond, C.P., Hollenbach, D., Kamp, I., & D’Alessio, P. 2007, Protostars andPlanets V, B. Reipurth, D. Jewitt, and K. Keil (eds.), University of ArizonaPress, Tucson, 951 pp., 2007., p.555-572Dzib, S., Loinard, L., Mioduszewski, A.J., Boden, A.F., Rodriguez, L.F., &Torres, R.M. 2010, ApJ, 718, 610Edwards, S., Fischer, W., Hillenbrand, L., & Kwan, J. 2006, ApJ, 646, 319Eiroa, C., Lenzen, R., Leinert, C., & Hodapp, K.-W. 1987, A&A, 179, 171Eisner, J.A., Hillenbrand, L.A., Carpenter, J.M., & Wolf, S. 2005, ApJ, 635, 396Enoch, M.L., Corder, S., Dunham, M.M., & Duchene, G. 2009, ApJ, 707, 103Enoch, M.L., Evans, N.J., Sargent, A.I., Glenn, J., Rosolowsky, E., & Myers, P.2008, ApJ, 684, 1240Evans, N.J., et al. 2009, ApJSS, 181, 321Ferreira, J., Dougados, C., & Cabrit, S. 2006, A&A, 453, 785Flaccomio, E., Stelzer, B., Sciortino, S., Micela, G., Pillitteri, I., & Testi, L. 2009,A&A, 505, 695Furlan, E., et al. 2008, ApJ, 176, 184Glassgold, A.E., Mamon, G.A., & Huggins, P.J. 1991, ApJ, 373, 254Gomez, M., Stark, D.P., Whitney, B.A., & Churchwell, E. 2003, ApJ, 126, 863Gonzalez-Alfonso, E., Wright, C.M., Cernicharo, J., Rosenthal, D., Boonman,A.M.S., & van Dishoeck, E.F. 2002, A&A, 386, 1074Goto, M., Usuda, T., Dullemond, C.P., Henning, T., Linz, H., Stecklum, B., &Suto, H. 2006, ApJ, 652, 758Goto, M., et al., 2011Greene, T.P., Barsony, M., & Weintraub, D.A. 2010, ApJ, accepted. astro-ph: // IAU Symposium 243 , ed. J. Bouvier &I. AppenzellerMayama, S., Tamura, M., Hayashi, M., Itoh, Y., Ishii, M., Fukagawa, M.,Hayashi, S., Oasa, Y., & Kudo, T. 2007, PASJ, 59, 1153McCaughrean, M.J., & Mac Low, M.-M. 1997, AJ, 113, 391McClure, M.K., Furlan, E., Manoj, P., Luhman, K.L., Watson, D.M., Forrest,W.J., Espaillat, C., Calvet, N., D’Alessio, P., Sargent, B., Tobin, J.J., &Chiang, H.-F. 2010, ApJS, 188, 75Meyer, M., & Wilking, B.A. 2009, PASP, 121, 350Mitchell, G.F., Maillard, J.-P., Allen, M., Beer, R., & Belcourt, K. 1990, ApJ,363, 554Najita, J.R., Carr, J.S., Glassgold, A.E., Shu, F.H., & Tokunaga, A.T. 1996, ApJ,462, 919Najita, J.R., Carr, J.S., & Mathieu, R.D. 2003, ApJ, 589, 931Najita, J.R., Carr, J.S., Glassgold, A.E., & Valenti, J.A. 2007, Protostars andPlanets V, B. Reipurth, D. Jewitt, and K. Keil (eds.), University of ArizonaPress, Tucson, 951 pp., 2007., p.507-522Najita, J.R., et al. 2009, ApJ, 697, 957Natta, A., Testi, L., & Randich, S. 2006, A&A, 452, 245Neufeld, D.A., & Hollenbach, D.J. 1994, ApJ, 428, 170Neufeld, D.A., & Yuan, Y. 2008, ApJ, 678, 984Nisini, B., Antoniucci, S., Giannini, T., & Lorenzetti, D. 2005, A&A, 429, 543Nutter, D.J., Ward-Thompson, D., & Andr´e, P. 2005, MNRAS, 357, 975Oliveira, I., Pontoppidan, K.M., van Dishoeck, E.F., Overvzier, R.A., Hern´andez,J., Sicilia-Aguilar, A., Eiroa, C., & Montesinos, B. 2009, ApJ, 691, 672Owen, J., & Ercolano, B., Clarke, C.J., & Alexander, R.D. 2010, MNRAS, 401,1415Padgett, D.L., Brandner, W., Stapelfeldt, K.R., Strom, S.E., Terebey, S., &Koerner, D. 1999, AJ, 117, 1490Panoglou, D., Cabrit, S., Pineau de Forets, G., Garcia, P.J.V., Ferreira, J., &Casse, F. 2010, A&A, submitted. e r cze g e t a l . : C O E m i ss i on fr o m E m b e dd e d O b j ec t s T a b l e E . . ON L I N E ON L Y : E qu i v a l e n t W i d t h s a nd F l ux e s i n S e l ec t e d L i n e s a Star Comp b P(2) 4682.64 P(5) 4708.77 P(8) 4735.87 P(11) 4763.99 P(17) 4823.3109 P(26) 4920.41 CO R(15) 4650.4256EW Flux EW Flux EW Flux EW Flux EW Flux EW Flux EW FluxL1551 IRS5 N 9.9 (1.3) 0.61 15.3 (1.3) 0.95 18.6 (0.6) 1.16 17.5 (1.8) 1.1 – – – – 2.7 (0.3) 0.16Elias 29 N 3.2 (0.1) 1.2 1.95 (0.13) 0.74 3.1 (0.1) 1.2 2.6 (0.2) 0.99 2.9 (0.3) 1.2 – – 0.32 (0.7) 0.12RNO 91 B 15 (3) 0.23 10 (4) 0.16 18 (3) 0.28 – – 18 (2) 0.29 – –HH 100 IRS B 6.8 (0.4) 1.2 4.8 (0.3) 0.86 6.6 (0.3) 1.2 5.6 (0.4) 1.01 – – 7.2 (0.5) 1.3 – –HH 100 IRS N 23 (2) 4.1 42.6 (3) 7.6 53 (3) 9.5 57 (3) 10.2 40 (3) 7.4 – – 0.29 (0.3) 0.051GSS 30 Blue 19.1 (0.1) 6.9 25.2 (0.1) 9.1 27.5 (0.1) 10.0 27. 6(0.3) 10.1 – – – – 3.3 (0.2) 1.16GSS 30 Red 24.0 (0.3) 8.6 27.0 (0.3) 9.7 37.5 (0.3) 13.6 45.8 (0.4) 16.7 – – 12.1 (0.4) < < . IRS 44 W Blue 36.5 (2.0) 0.42 – – – – – – 33 (2) 0.52 – – 5.4 (1.3) 0.082IRS 44 W Red 91 (8) 2.6 67 (7) 1.9 67 (8) 1.9 47 (9) 1.3 80 (10) 2.3 – – – –CrA IRS 2 T 14.4 (0.7) 1.2 18.2 (0.3) 1.5 21.1 (0.4) 1.8 24.2 (0.4) 2.1 23.9 (0.2) / NSVS 20 N Low S / NTMC 1A N Low S / NWL 6 No CO emission detectedElias 32 No CO emission detectedIRS 44 E No CO emission detectedIRS 43 N No CO emission detected Equivalent width ( ∼ σ error) in km s − and Flux in 10 − erg cm − − for 6 CO and 1 CO line.Fluxes are calculated from line equivalent widths and continuum flux, and are overestimated when continuum is spatially extended beyond the slit width.Flux uncertainty also includes ∼
30% relative uncertainty in continuum flux across the M-band. b Component, Br = broad, N = Narrow, NB = Narrow Blue, T = Total P(37) 5.053 µ m Upper limit from other CO lines because R(15) from GSS 30 has a high upper limit. P(18) 4.834 µ m P a s c u cc i , I ., H o ll e nb ac h , D ., N a jit a , J ., M u ze r o ll e , J ., G o r ti , U ., H e r cze g , G . J ., H ill e nb r a nd , L . A ., K i m ., J . S ., C a r p e n t e r , J . M ., M e y e r , M . R ., M a m a j e k , E . E ., & B ou w m a n , J . , A p J , , P a s c u cc i , I l ., & S t e r z i k , M . , A p J , , P on t opp i d a n , K . M ., S c h ¨ o i e r , F . L ., v a n D i s ho ec k , E . F ., & D a r t o i s , E . , A & A , , P on t opp i d a n , K . M ., F r a s e r , H . J ., D a r t o i s , E ., T h i , W . - F ., v a n D i s ho ec k , E . F ., B oog e r t , A . C . A , d ’ H e nd ec ou r t , L ., T i e l e n s , A . G . G . M ., & B i ss c hop , S . E . , A & A , , P on t opp i d a n , K . M ., e t a l . , c on f . p r o c ... P on t opp i d a n , K . M ., D u ll e m ond , C . P ., v a n D i s ho ec k , E . F ., B l a k e , G . A ., B oog e r t , A . C . A ., E v a n s , N . J ., K e ss l e r- S il acc i , J . E ., & L a hu i s , F . , A p J , , erczeg et al.: CO Emission from Embedded Objects Table E.3.
Comparing ISAAC and CRIRES Equivalent Widths a Star Seeing P(6) EW H S(9) EWCRIRES ISAAC CRIRES ISAAC CRIRES ISAACHH 100 IRS 0.44 0.41 9 b b – –WL 12 0.31 0.43 – – 16 10.9WL 6 0.29 0.88 – – 6.9 10.5IRS 43 S + N 0.32 0.58 13 19 1.7 4.0IRS 44 E + W 0.32 0.72 9 20 9.9 16.4IRS 63 0.37 0.56 -1.7 4 0.7 –Elias 32 0.32 0.58 – 21 – 3.8RNO 91 0.28 0.52 17 11 – –SVS 20 S 0.42 0.46 1.1 – 4.2 –SVS 20 N 0.42 0.46 – – – 8.5GSS 30 0.30 0.53 49 58 3.2 12.2 a Seeing in ′′ , Equivalent width in km s − b P(9) rather than P(6)
Pontoppidan, K.M., Blake, G.A., van Dishoeck, E.F., Smette, A., Ireland, M.J.,Brown, J. 2008, ApJ, 684, 1323Pontoppidan, K.M., Blake, G.A., van Dishoeck, E.F., Smette, A., Ireland, M.J.,Brown, J. 2008, ApJ, 684, 1323Pontoppidan, K.M., et al. Messenger paperPrato, L., Lockhart, K.E., Johns-Krull, C.M., & Rayner, J.T. 2009, AJ, 137, 3931Pyo, T.-S., Hayashi, M., Kobayashi, N., Terada, H., Tokunaga, A.T. 2009, ApJ,694, 654Ratzka, T., K¨ohler, R., & Leinert, C. 2005, A&A, 437, 611Rettig, T.W., Haywood, J., Simon, T., Brittain, S.D., & Gibb, E. 2004, ApJ, 616,L163Rieke, G.H., & Leobfsky, M.J. 1985, ApJ, 288, 618Robitaille, T.P., Whitney, B.A., Indebetouw, R., Wood, K., & Denzmore, P. 2006,ApJS, 167, 256Rodriguez, L.F., D’Alessio, P., Wilner, D.J., Ho, P.T.P., Torrelles, J.M., Curiel,S., et al. 1998, Nature, 395, 355Saito, M., Kawabe, R., Kitamura, Y., & Sunada, K. 2001, ApJ, 547, 840Saucedo, J., Calvet, N., Hartmann, L., & Raymond, J. 2003, ApJ, 591, 275Salyk, C., Blake, G.A., Boogert, A.C.A., & Brown, J.M. 2007, ApJ, 655, L105Salyk, C., Blake, G.A., Boogert, A.C.A., & Brown, J.M. 2009, ApJ, 699, 330Shang, H., Li, Z.-Y., & Hirano, N. 2007, Protostars and Planets V, B. Reipurth,D. Jewitt, and K. Keil (eds.), University of Arizona Press, Tucson, 951 pp.,2007., p.261-276Shu, F., Najita, J., Ostriker, E., Wilkin, F., Ruden, S., & Lizano, S. 1994, ApJ,429, 781Skrutskie, M.F., et al. 2006, AJ, 131, 1163.Sloan, G.C., Kraemer, K.E., Price, S.D, & Shipman, R.F. 2003, ApJS, 147, 379Smith, R., et al. 2011, submittedSpaans, M., Hogerheijde, M.R., Mundy, L.G., & van Dishoeck, E.F. 1995, ApJ,455, L167Stark, D.P., Whitney, B.A., Stassun, K., & Wood, K. 2006, ApJ, 649, 900Takami, M., Beck, T.L., Pyo, T.-S., McGregor, P., & Davis, C. 2007, ApJ, 670,L33Terebey, S., van Buren, D., Hancock, T., Padgett, D.L., & Brundage, M. 2001,ASP Conference Series 243,
From Darkness to Light , eds. T. Montmerle & P.Andr´eThi, W.-F., van Dishoeck, E.F., Pontoppidan, K.M., & Dartois, E. 2010,MNRAS, 406, 1409Tobin, J.J., Hartmann, L., & Loinard, L. 2010, ApJ, accepted. arXiv:1008.3429v1Torres, R.M., Loinard, L., Mioduszewski, A.J., & Rodriguez, L.F. 2009, ApJ,698, 242Valenti, J.A., Fallon, A.A., & Johns-Krull, C.M. 2003, ApJS, 147, 305van Boekel, R., G¨udel, M., Henning, T., Lahuis, F., & Pantin, E. 2009, A&A,497, 137van der Plas, G., van den Ancker, M.E., Acke, B., Carmona, A., Dominik, C.,Fedele, D., & Waters, L.B.F.M. 2009, A&A, 500, 1137van Dishoeck, E.F., & Black, J.H. 1988, ApJ, 334, 771van Kempen, T.A., van Dishoeck, E.F., G¨usten, R., Kristensen, L.E., Schilke,P., Hogerheijde, M.R., Boland, W., Nefs, B., Menten, K.M., Baryshev, A., &Wyrowski, F. 2009, A&A, 501, 633van Kempen, T.A., van Dishoeck, E.F., Salter, D.M., Hogerheijde, M.R.,Jørgensen, J.K., & Boogert, A.C.A. 2009, A&A, 498, 167van Kempen, T.A., van Dishoeck, E.F., Hogerheijde, M.R., & G¨usten, R. 2009,A&A, 508, 259van Kempen, T.A., et al. 2010, A&A, 518, L121 Walter, F.M., et al. 2003, AJ, 126, 3076Watson, D.M., Bohac, C.J., Hull, C., Forrest, W.J., et al. 2007,
Nature , 448, 1026Weintraub, D.A., Tegler, S.C, Kastner, J.H., & Rettig, T. 1994, ApJ, 423, 674White, R.J., & Hillenbrand, L.A. 2004, ApJ, 616, 998Whitney, B.A., & Hartmann, L. 1993, ApJ, 402, 605Whittet, D.C.B., Adamson, A.J., Duley, W.W., Geballe, T.R., ney, B.A., &Hartmann, L. 1993, ApJ, 402, 605Wilner, D.J., Ho, P.T.P., & Rodriguez, L.F. 1996, ApJ, 470, L117Wilson, T.L. 1999, Rep.Prog.Phys., 62, 143Wood, K., Smith, D., Whitney, B., Stassun, K., Kenyon, S.J., Wol ff , M.F., &Bjorkman, K.S. 2001, ApJ, 561, 299Wu, P.-F., Takakuwa, S., & Lim, J. 2009, ApJ, 698, 184Yang, J., Jiang, Z., Wang, M., Ju, B., & Wang, H. 2002, ApJS, 141, 157Ybarra, J.E., Barsony, M., Haisch, K.E., Jarrett, T.H., Sahai, R., & Weinberger,A.J. 2006, ApJ, 647, L159Zhang, M., & Wang, H. 2009, ApJ, 138, 1830, M.F., &Bjorkman, K.S. 2001, ApJ, 561, 299Wu, P.-F., Takakuwa, S., & Lim, J. 2009, ApJ, 698, 184Yang, J., Jiang, Z., Wang, M., Ju, B., & Wang, H. 2002, ApJS, 141, 157Ybarra, J.E., Barsony, M., Haisch, K.E., Jarrett, T.H., Sahai, R., & Weinberger,A.J. 2006, ApJ, 647, L159Zhang, M., & Wang, H. 2009, ApJ, 138, 1830