Äîêóìåíò âçÿò èç êýøà ïîèñêîâîé ìàøèíû. Àäðåñ îðèãèíàëüíîãî äîêóìåíòà : http://www.arcetri.astro.it/~lt/preprints/i20126_deb/pp2.ps.gz Äàòà èçìåíåíèÿ: Tue Sep 11 16:56:51 2007 Äàòà èíäåêñèðîâàíèÿ: Sat Dec 22 07:34:47 2007 Êîäèðîâêà: Ïîèñêîâûå ñëîâà: m 17

ApJ 6dec99, referee and author comments incorporated
The Molecular Outflow and Possible Precessing Jet from the
Massive Young Stellar Object IRAS 20126+4104
D.S. Shepherd 1;2 , K.C. Yu 3 , J. Bally 3 , & L. Testi 4;2
ABSTRACT
We present images of the molecular gas in the IRAS 20126+4104 massive outflow
and examine the interaction between the energetic outflowing material and the
surrounding molecular cloud. Mosaic interferometric images in CO(1--0), 13 CO(1\Gamma0),
C 18 O(1\Gamma0), C 17 O(1\Gamma0), and millimeter continuum emission are compared with
mid­infrared images at 12.5 ¯m & 17.9 ¯m, near­infrared images in the K s band
(2.17 ¯m) and H 2 line emission, and optical Hff and [SII] images. We show that the
molecular outflow is approximately 6 \Theta 10 4 years old with a mass of about 50­60 M fi
and mass outflow rate —
M f ¸ 8 \Theta 10 \Gamma4 M fi yr \Gamma1 . The driving source is located near the
center of the ?
¸ 300 M fi molecular cloud and the mass of the disk plus circumstellar
envelope traced by millimeter continuum emission is ¸ 50 M fi . The outflow appears
to be bounded on most sides by higher density gas traced by C 18 O emission. Shocks
identified by H 2 and [SII] emission knots follow a NW­SE jet close to the young stellar
object and then rotate more N­S along the edges of the CO flow. The most likely
interpretation appears to be that the knots trace the working surfaces of a collimated
jet which precesses through an angle of ¸ 45 ffi . Possible mechanisms that could produce
the jet precession include tidal interactions between the disk and a companion star in
a non­coplanar orbit or an anisotropic accretion event that dramatically altered the
angular momentum vector of the disk.
Subject headings: stars: formation -- nebulae: HII regions -- ISM: jets and outflows --
ISM: molecules
1 National Radio Astronomy Observatory, 1003 Lopezville Rd., Socorro, NM 87801
2 Division of Physics, Mathematics and Astronomy, California Institute of Technology, MS 105­24, Pasadena, CA
91125
3 Center for Astrophysics and Space Astronomy, CB 389, University of Colorado, Boulder, CO 80309
4 Osservatorio Astrofisco di Arcetri, Largo Enrico Fermi 5, I­50125 Firenze

-- 2 --
1. INTRODUCTION
Energetic molecular outflows from massive protostars share some characteristics with their
low­mass counterparts. In particular, both massive and low­mass flows are probably powered by
disk accretion, and the flow energetics scale with the luminosity of the source (c.f. Cabrit & Bertout
1992; Shepherd & Churchwell 1996; Richer et al. 2000). However, the details of the interaction
between accretion and outflow are poorly understood and it is not clear whether outflows from
high and low­mass protostars are powered by the same mechanism (c.f. Rodr'iguez, 1995; Cabrit,
Raga & Gueth 1997; Bachiller & Tafalla 1999; K¨onigl & Pudritz 2000; Shu et al. 2000; Eisl¨offel et
al. 2000; Richer et al. 2000). To address this problem, the properties of outflows from luminous
young stellar objects at a range of evolutionary stages must be studied in detail to compare with
those from lower luminosity systems. In this paper, we focus our attention on a particularly young
outflow from an early B protostar, IRAS 20126+4104.
IRAS 20126+4104 (hereafter I20126), at ff(B1950) 20 h 12 m 41.0 s and ffi (B1950) 41 ffi 04 0 21:0 00 , is
located in a dark cloud in the Cygnus­X region at a kinematic distance of 1.7 pc. The far infrared
luminosity is about 10 4 L fi and the IRAS colors match the selection criteria for ultracompact
H II regions (Wood & Churchwell 1989). The IRAS source is embedded in a 230 M fi NH 3 core
of roughly 0.4 pc diameter (Estalella et al. 1993) and lies close to the origin of a N­S molecular
outflow (Wilking et al. 1990). Near the center of the outflow is a bright, bipolar infrared reflection
nebula with a NW­SE orientation (Hodapp 1994, Cesaroni et al. 1997, 1999, hereafter C97 & C99).
In addition, a well­collimated jet from the protostar is visible in shock­enhanced SiO emission
(C97, C99) and H 2 emission (C97, Ayala, et al. 1998), while H 2 O masers within 1 00 of the source
are aligned along the jet axis (Tofani et al. 1995). Observations by Kawamura et al. (1999) in
the CO(J=7--6) and NH 3 (3,3) transitions imply an underlying jet density of at least 10 5 cm \Gamma3 .
A rotating 10,000 AU NH 3 (1,1) torus with a mass of approximately 20 M fi (Zhang et al. 1998)
surrounds a more compact disk detected in CH 3 CN and thermal dust emission (C99, Hofner et
al. 1999). The roughly 1700 AU diameter inner disk appears to lie along an axis that is nearly
perpendicular to the jet axis. Near the peak of the 3 mm continuum emission, Hofner et al. (1999)
have detected two elongated structures in 3.6 cm continuum emission that are within 1 00 of each
other. The emission is consistent with the presence of two parallel jets of ionized gas. The northern
jet is elongated in the direction of the H 2 O maser emission and is coincident with the disk inferred
from CH 3 CN emission. A disk counterpart is not detected toward the southern jet nor is there
H 2 O maser emission associated with it.
If the same young stellar object (YSO) is producing one of the inner jets with a NW­SE
orientation and the N­S molecular flow, the change in position angle implies that the jet may
be precessing or wandering over an angle of ¸ 40 ffi . Jet flows from low­mass YSOs have been
observed to have an S­shaped symmetry with typical opening angles !
¸ 10 ffi , but a jet precession
angle of 40 ffi is unusually large. It is also possible that there are multiple outflows emerging with
different orientations. Clearly, high resolution observations are are needed to determine the detailed
properties of I20126 and identify which scenario is correct.

-- 3 --
We have therefore made high resolution aperture synthesis maps of the I20126 region in
CO(J=1--0), 13 CO(1\Gamma0), C 18 O(1\Gamma0), and C 17 O(1\Gamma0) to trace the dynamics of the outflow from
the large scale (1 0 \Gamma 2 0 ) to the small scale (inner 20 00 ). We have also obtained images at visible, near­
infrared, and mid­infrared wavelengths to provide an understanding of the shocks within the flow
and the relationship between the outflow and the larger scale molecular cloud. In x2 we describe
the observations. In x3 we present the results and point out key features of the outflow and YSO.
In x4 we discuss the flow components. We present a summary of our findings in x5.
2. OBSERVATIONS
2.1. OVRO observations
Observations in the 1.3 mm & 2.7 mm continuum, CO(1\Gamma0), 13 CO(1\Gamma0), C 18 O(1\Gamma0), and
C 17 O(1\Gamma0) lines were made with the Owens Valley Radio Observatory (OVRO) millimeter wave
array of six 10.4 m telescopes between 1998 February 2 and 1999 February 3. Projected baselines
ranging from 15 to 484 meters provided sensitivity to structures up to about 16 00 . The final ¸ 3 0 \Theta 3 0
mosaic images of both line and continuum emission consist of 13 fields with primary beam ¸ 65 00
(FWHM), spaced 30 00 apart. The total integration time on source was approximately 5.3 and 2.7
hours/pointing center for the CO and 13 CO/C 18 O mosaics, respectively. The integration time on
source at the central position was 10.5 hours for the C 17 O and 3 mm continuum observations, and
2.8 hours for the 1 mm continuum observations. Cryogenically cooled SIS receivers operating at
4 K produced typical single sideband system temperatures of 200 to 600 K. Gain calibration used
the quasar BL Lac and the passband calibrators were 3C 84 and 3C 273. Observations of Uranus,
Neptune, and 3C 273 provided the flux density calibration scale with an estimated uncertainty of
¸ 20%. Calibration was carried out using the Caltech MMA data reduction package (Scoville et
al. 1993). Images were produced using the MIRIAD software package (Sault et al. 1995) and
deconvolved with a maximum­entropy­based algorithm designed for mosaic images (Cornwell &
Braun 1988). The overlap region of the mosaic is corrected for primary beam attenuation.
The spectral band pass for all lines was centered on the systemic local standard of rest velocity
(v LSR ) \Gamma3:5 km s \Gamma1 . CO emission was detected across the entire 77 km s \Gamma1 band pass. To optimize
sensitivity to extended structure, the CO uv data were convolved with a 6 00 taper resulting in a
synthesized beam of 7:53 00 \Theta 7:09 00 (FWHM) at P.A. \Gamma80:6 ffi . At the spectral resolution of 1.3 km s \Gamma1 ,
the RMS noise is 0.12 Jy beam \Gamma1 . The 13 CO and C 18 O uv data were similarly convolved, resulting
in synthesized beams of 7:46 00 \Theta 6:03 00 (FWHM) at P.A. \Gamma9:8 ffi and 7:42 00 \Theta 5:85 00 (FWHM) at P.A.
\Gamma10:6 ffi , respectively. The RMS noise of the 13 CO map is 65 mJy beam \Gamma1 at spectral resolution
2.72 km s \Gamma1 , and the RMS noise in the C 18 O map is 0.14 Jy beam \Gamma1 at spectral resolution
0.68 km s \Gamma1 . The C 17 O maps were convolved with a 2 00 taper which produced a synthesized beam
of 3:43 00 \Theta 3:09 00 (FWHM) at P.A. \Gamma87:8 ffi . The RMS noise of the C 17 O maps is 80 mJy beam \Gamma1
with spectral resolution 0.67 km s \Gamma1 .

-- 4 --
Simultaneous 3 mm continuum observations were made in a 1 GHz bandwidth channel which
for CO observations had a central frequency of 112.8 GHz and for 13 CO observations had a central
frequency of 113.5 GHz. The central frequency of the combined bands is 113.15 GHz. Continuum
observations at 1 mm wavelength were made in two 1 GHz bands with an average frequency of
228.1 GHz. The images had a synthesized beam 2:93 00 \Theta 2:53 00 (FWHM) at P.A. 78:4 ffi and RMS
noise 1.7 mJy beam \Gamma1 at 113 GHz and a beam of 1:18 00 \Theta 0:81 00 (FWHM) at P.A. \Gamma51:2 ffi and RMS
noise 19 mJy beam \Gamma1 at 228 GHz.
2.2. Mid­Infrared Observations
Infrared images of I20126 at 12.5 ¯m and 17.9 ¯m were obtained on UT 1998 Oct 7 at
the W.M. Keck Observatory 5 using the Jet Propulsion Laboratory mid­infrared camera MIRLIN
mounted at the Keck II 10 m telescope. The MIRLIN focal plane array is a 128\Theta128 Si:As Boeing
BIB detector, with a plate scale of ¸0.14 arcsec/pixel, when mounted at the f/40 bent­Cassegrain
focus of the Keck II telescope. The source was observed with the N5 (– eff ¸ 12:5 ¯m, \Delta– ¸ 1:2 ¯m)
and Q s (– eff ¸ 17:9 ¯m, \Delta– ¸ 2:0 ¯m) filters. Data acquisition employed a chop­nod cycle with a
chop frequency of ¸4 Hz, and a few hundred chops per nod. Chop and nod throws of ¸ 10 00 were
oriented in orthogonal directions on the sky. Background emission was subtracted by computing
the double­difference of each chop­nod cycle. Final images were obtained by shifting and adding
the background subtracted frames. Poor weather conditions during the observations precluded
photometric calibration. The 12.5 ¯m image has a total exposure time of 195 seconds and the
17.9 ¯m image has a total exposure time of 12 seconds. The point spread function FWHM was
¸ 0:4 00 and ¸ 0:5 00 at 12.5 ¯m and 17.9 ¯m, respectively. No astrometric calibration could be
derived since there were no other sources within the field. To obtain an approximate position, the
peak of the mid­infrared emission was assumed to be coincident with the 3 mm continuum peak.
This assumption is probably accurate to within about an arcsec.
2.3. Near­Infrared Observations
Observations in K s band and H 2 narrow band were taken with at the Palomar 60­inch (1.5 m)
telescope with the Cassegrain IR Camera (IRCAM) on 1997 October 14 and 16. The Rockwell
256 \Theta 256 HgCdTe (NICMOS­3) array has a read noise of ¸ 50 e \Gamma , a pixel scale at the f/8.75
focal ratio of 0.624 arcsec/pixel, and a field­of­view of 160 00 (Murphy et al. 1995; Persson & Pahre
1995). Images were obtained through a narrow band filter centered at 2.12 ¯m which includes the
5 The W.M. Keck Observatory is operated as a scientific partnership among the California Institute of Technology,
the University of California and the National Aeronautics and Space Administration. The Observatory was made
possible by the generous financial support of the W.M. Keck Foundation.

-- 5 --
H 2 v = 1--0 S(1) line and a broadband K s filter 6 centered at 2.17 ¯m with a bandwidth of 0.33 ¯m
(Persson et al. 1998). Ten dithered exposures were taken in each of the filters, with exposure times
of 120 s and 10 s in the H 2 and K s filters, respectively. On­source exposures were immediately
followed by offsetting the telescope to a sky position. The seeing was about 1:4 00 FWHM. Multiple
sky frames were median combined and then subtracted from the on­source frames. After dividing
by a flat­field, the dithered exposures were median combined to remove detector defects. The
continuum emission in the final H 2 mosaic of dithered images was estimated by comparing the
total flux from 15 stars in the K s and H 2 images, and calculating an averaged flux ratio between
the two filters. The standard deviation of the measured flux ratios was ¸ ! 3%. The K s image was
renormalized by the flux ratio and subtracted from the H 2 image to obtain an H 2 line­only image.
Astrometric calibration was done by comparing common stars found in our field and the STScI
Digital Sky Survey plates.
2.4. Visible Hff and [SII] Observations
We obtained images of the I20126 region on the nights of 1997 October 29­31 at the prime
focus of the Mayall 4 m telescope using the engineering grade MOSAIC 8192x8192 CCD camera.
The field of view with the imaging correctors is approximately 36 0 x 36 0 with a scale of 0.26 00 per
pixel. We obtained 600s exposures through narrow band filters transmitting Hff and [SII]. The Hff
filter is centered on 6563 š A with a band pass of 75 š A, while the [SII] filter is centered on 6723 š A with
a band pass of 80 š A. To minimize the effect of cosmetic blemishes and gaps between the individual
CCDs, we centered I20126 in the middle of the best 2048x4096 CCD. The data was over­scanned,
trimmed, dark subtracted and then flat fielded with sky flats in the standard manner using IRAF
packages. The resulting full MOSAIC image still showed gaps between CCDs and the outer areas
of the image suffered from cosmetic problems that could not be calibrated out. Thus, the images
presented in this paper present only the inner 6 0 \Theta 6 0 field centered on the I20126 outflow.
3. RESULTS
3.1. The I20126 molecular outflow
Figures 1 through 5 present the large­scale morphology and kinematics of the CO outflow and
the infrared and visible emission in the region. The outflow axis is predominantly north­south
(position angle 171 ffi ) with red­shifted gas in the south and blue­shifted gas in the north. The
outflow is centered near the 3 mm millimeter continuum peak (C97). From end­to­end, the flow
6 The Ks filter has half­power transmission at ¸ 2:0 ¯m and 2.3 ¯m rather than ¸ 2:0 ¯m and 2.4 ¯m for the K
filter. The cutoff towards the long wavelength end of the Ks filter results in roughly a factor of 2 reduction in the
thermal background when compared to the K filter.

-- 6 --
measures 2 0 which corresponds to nearly 1 pc at the distance of 1.7 kpc. The red­shifted lobe
is roughly twice as long as the blue­shifted lobe. Velocity structure in both lobes is somewhat
chaotic with a shell of low­velocity gas (v !
¸ 10 km s \Gamma1 ) surrounding clumps of higher velocity gas
(v ? 10 km s \Gamma1 ). The highest velocity red­shifted gas is on the east side of the red lobe while the
highest velocity blue­shifted gas is located less than 10 00 NW of the continuum peak.
Assuming the flow is roughly conical in shape, the measured semi­opening angle is between
about 30 ffi and 40 ffi and there is very little overlap between the red and blue­shifted lobes. Therefore,
the flow does not intersect the plane of the sky and the geometry most closely corresponds to case
2 in Cabrit & Bertout (1986) (see also Chandler et al. 1996). This indicates that a reasonable
range for the inclination angle of the outflow i is between 40 ffi and 50 ffi (measured with respect to
the line of sight). Thus, we adopt the value i = 45 ffi .
Near­infrared reflection nebulae are centered near the continuum source and H 2 emission is
within the reflection nebulae and to the NW and SE of the flow (Fig. 4, see also C97). The positions
of the H 2 knots are listed in Table 1 and are indicated by triangles in Figs. 1 & 2. Within 10 00 of the
central source, the position angle of the outflow defined by the infrared emission is approximately
120 ffi . This is consistent with position angles seen in HCO + , SiO (C97, C99), NH 3 (3,3) (Kawamura
et al. 1999, Zhang et al. 1999), and ionized gas (Hofner et al. 1999). The outer H 2 emission knots
have a more north­south orientation and appear to follow the boundaries of the CO outflow lobes.
The two southern­most H 2 knots have a bright star between them suggesting that they may be
associated with a different flow. However, this star has the colors of a foreground object with no
significant infrared excess (C97) and no millimeter continuum emission is detected at this position
(discussed in x3.2). Thus, the southern H 2 knots are most likely associated with the outflow powered
by I20126.
Mid­infrared emission at 12.5 ¯m and 17.9 ¯m is shown in Fig. 5. The 17.9 ¯m image is
also compared to the structure of H 2 emission near the central star. To estimate the approximate
positions in the mid­IR maps, we assume the brightest peak of the mid­IR emission is coincident
with the location of the millimeter continuum source (see also C99). The NW emission peak is
more extended than the SE peak and the relative brightness of the NW peak increases at the longer
wavelength which may indicate a cooler temperature. The morphology and relative intensities are
consistent with warm emission near the central star and either cooler dust emission or a reflection
nebula stretching toward the blue­shifted outflow lobe in the NW. Shocked H 2 is present within
and just beyond the NW emission region.
The CO outflow is seen in projection against the larger scale molecular cloud (Fig. 2).
Filamentary ridges of Hff and [SII] emission trace the cloud boundary on three sides and the higher
extinction within the cloud is evident by the lack of background stars. The entire 36 0 \Theta 36 0 MOSAIC
field (not shown) is filled with a combination of diffuse and filamentary Hff and [SII] nebulosity.
The I20126 cloud is embedded in a large scale network of [SII] bright filaments extending from
northeast to the southwest across the field of view. The morphology and relatively large [SII]/Hff

-- 7 --
ratio indicates that the filaments may trace shock fronts produced by either an old supernova
remnant or possibly a very large scale superbubble. Many dark clouds are visible in silhouette
against the background star field and nebulosity. Some of these clouds, including the one containing
I20126 shown in Fig. 2, are rimmed by ionization fronts facing towards the east indicating that hot
stars illuminate the eastern rim of the the I20126 cloud. Thus, it is possible that star formation
in this cloud was triggered by the energy release from massive stars associated with the extended
nebulosity in this portion of Cygnus.
As expected, there is little nebulosity associated with the outflow within the confines of the
CO lobes due to the high optical extinction. However, just north of the blue­shifted outflow lobe is
a faint, semi­circular arc of Hff and [SII] emission that resembles a bow shock and extends to the
northern rim of the cloud. In the center of this ``arc'' of emission, an amorphous, [SII]­bright knot
is visible at ff(B1950) 20 h 12 m 39.88 s ffi(B1950) 41 ffi 05 0 15:6 00 . The position of this knot relative to
the CO flow and H 2 emission knots is indicated by a filled diamond in Figs. 1 & 2 and lies within
2 ffi of the outflow axis determined from the CO geometry. There is also a pair of Hff and [SII]
filaments projecting beyond the southeastern rim of the cloud, below the red­shifted CO lobe. We
speculate that these features may correspond to portions of the CO flow which have broken free of
the molecular cloud and have been dissociated or, perhaps, ionized. Our data cannot confirm this
interpretation; however, Hff spectroscopy would resolve this uncertainty.
The combined H 2 and [SII]­bright knots display a remarkable S­shaped symmetry about the
central star. The three inner­most H 2 knots lie along the SiO jet axis identified by C97 while the
outer H 2 and [SII] knots show a rotation of approximately 45 ffi relative to the jet axis.
3.2. Circumstellar material surrounding the central B star
Figure 6 presents 1 mm and 3 mm continuum images of thermal dust emission near the central
star. No other continuum sources were detected in the mosaic image above a level of 3oe (5.1 mJy).
A significant fraction of the 3 mm continuum emission is extended over a 10,000 AU diameter region
as shown in the top panel of Fig. 6 where all baselines were used to recover extended emission.
An integrated flux density of 81.14 mJy and peak flux density of 34.7 mJy beam \Gamma1 is found at
position ff(B1950) 20 h 12 m 40.99 s ffi (B1950) +41 ffi 04 0 21:0 00 . The location of the 3 mm emission peak
is within 0:4 00 of the position reported by C99. The integrated flux density is approximately three
times larger than that found by C99 and is probably due to the larger inner hole in the Plateau
de Bure (PdB) uv coverage which resolves out more extended emission. The center panel of Fig.
6 shows the compact 3 mm emission where much of the extended structure has been resolved out
by using only longer baselines. The total flux density recovered in this map is 62.2 mJy and the
peak flux density is 20.6 mJy beam \Gamma1 . The location of the 3 mm emission peak remains the same.
The central peak of the compact emission is nearly coincident with the H 2 O masers tracing the
northern ionized jet while the extension to the south is coincident with the southern ionized jet
(Tofani et al. 1995, Hofner et al. 1999).

-- 8 --
Figure 7 presents the spectral energy distribution (SED) of I20126 from 3.6 cm to 2.2 ¯m,
including fluxes from the literature. The dashed line is our best­fit SED using a constant dust
temperature of 44 K and a dust emissivity index, fi = 1:5. The integrated bolometric luminosity
L bol is ¸ 10 4 L fi . The estimated flux at 1 and 3 mm from free­free emission is ¸ 0:2 mJy, assuming a
spectral index of \Gamma0:1 between 3.6 cm (Hofner et al. 1999) and millimeter wavelengths. This is well
below the 1oe RMS of 19 and 1.7 mJy beam \Gamma1 in the 1 and 3 mm continuum images, respectively.
Thus, we assume the millimeter flux density is due to thermal dust emission. This is in good
agreement with C99 and Hofner et al. (1999). Following the method of Hildebrand (1983), the mass
of gas and dust is estimated from the millimeter continuum emission using M gas+dust = Fš D 2
Bš(T d ) Ÿš
where D is the distance to the source, F š is the continuum flux density due to thermal dust emission
at frequency š, B š is the Planck function at temperature T d . Assuming a gas­to­dust ratio of 100,
the dust opacity per gram of gas is taken to be Ÿ š = 0:004( š
245GHz ) fi cm 2 g \Gamma1 based on measurements
by Kramer et al. (1998). We assume that the emission is optically thin, and the temperature of
the dust can be characterized by a single value. Using values of T d = 44 K and fi = 1:5 determined
from the SED, we find the mass of gas and dust associated with the 3 mm continuum emission is
approximately 52 M fi .
The integrated flux at 1 mm is 393.9 mJy with peak flux of 254.9 mJy beam \Gamma1 at position
ff(B1950) 20 h 12 m 40.001 s ffi (B1950) +41 ffi 04 0 21:00 00 (bottom panel of Fig. 6). The source is
unresolved with a deconvolved size of less than 1:06 00 \Theta 0:83 00 at P.A. --81 ffi . Again, assuming
T d = 44 K and fi = 1:5, we find that the mass of gas and dust associated with the 1 mm continuum
emission is approximately 22 M fi . The size of the 1 mm emission region is significantly less than
that found at 3 mm which indicates the more extended flux is probably resolved out. Thus, as
expected, the 1 mm mass estimate is lower than that derived from the 3 mm continuum emission.
3.3. Outflow Mass Estimates
To estimate the CO optical depth in the flow and examine the dense gas morphology, mosaic
images of 13 CO, C 18 O, and C 17 O emission were also obtained along with the CO mosaic presented
in x3.1. Comparisons between the outflow morphology traced by CO and 13 CO red and blue­shifted
gas and the C 18 O integrated intensity are presented in Fig. 8. Channel maps of 13 CO and C 18 O
emission are shown in Fig. 9. The remnant molecular core traced by C 18 O emission is fractured and
filamentary and appears to encase the outflow on most sides. In contrast, the 13 CO emission near
v LSR nearly fills the mosaiced field. The morphology of the integrated C 17 O emission tracing the
highest density gas is shown in Fig. 10. Only emission near the continuum source is detected with
peak position ff(B1950) 20 h 12 m 41.04 s ffi(B1950) +41 ffi 04 0 20:0 00 , peak flux density 506 mJy beam \Gamma1
km s \Gamma1 , and total flux 1.2 Jy.
The CO optical depth as a function of velocity is determined using CO and 13 CO spectra
that have been convolved with a 40 00 beam (Fig. 11). We assume 13 CO is optically thin at all
velocities which is probably valid in the line wings; however, 13 CO is likely to be optically thick in

-- 9 --
the line core. Further, C 18 O cannot be used to estimate the optical depth of the 13 CO emission
because the two isotopes have very different morphologies over the same velocity range. Thus, mass
estimates based on CO and 13 CO emission within \Sigma2:7 km s \Gamma1 of v LSR represent a lower limit.
The optical depth at red and blue­shifted velocities is estimated from the CO/ 13 CO ratio of the
spectra centered on the red and blue­shifted outflow lobes, respectively. In channels where no 13 CO
emission is detected, we assume the CO is optically thin. Similarly, the spectra constructed from
emission at the central position were used to estimate a lower limit to the optical depth within
\Sigma2:7 km s \Gamma1 of v LSR . The CO line profile also shows a dip near v LSR indicating that the line
suffers from self­absorption and/or the interferometer map is missing some extended flux at lower
velocities.
The mass associated with CO line emission is calculated following Scoville et al. (1986)
and measuring the flux density in each velocity channel (corrected for the mosaic primary beam
attenuation). We assume the gas is in LTE, at a temperature of 44 K with [CO]/[H 2 ] = 10 \Gamma4 ,
and [CO]/[ 13 CO] = 68.4, [CO]/[C 18 O] = 513.4, and [CO]/[C 17 O] = 2824 at the galacto­centric
distance of 8.1 kpc (Wilson & Rood 1994). Table 2 summarizes the physical properties of the
molecular gas in the flow. The total flow mass M f is given by P
M i , momentum P is P
M i v i ,
and energy E is 1
2
P
M i v 2
i where M i is the flow mass in velocity channel i, and v i is the central
velocity of the channel relative to v LSR . A characteristic flow timescale t d is R f = ! V ?, where
! V ? is P=( P
M i ) (Cabrit & Bertout 1990) and R f is the flow radius. Finally, the mass outflow
rate —
M f is P
M i =t d and force F is P=t d . The total molecular mass in the flow (jvj ? 2:7 km s \Gamma1 )
is approximately 53 M fi . This value compares reasonably well with the outflow mass estimate
of 67 M fi by Wilking et al. (1990) made with a single dish telescope. Given the differences in
the assumptions above, it does not appear that the OVRO interferometer is missing a significant
amount of extended emission in the outflow.
Using emission within \Sigma2:7 km s \Gamma1 of v LSR within our mosaic field, we calculate a lower limit
to the quiescent molecular cloud mass associated with 13 CO emission to be 249 M fi . The molecular
mass associated with total C 18 O emission is 104.3 M fi and the momentum and kinetic energy in
the C 18 O ridges is approximately 190 M fi km s \Gamma1 yr \Gamma1 and 4 \Theta 10 45 ergs, respectively. Finally, the
mass associated with C 17 O emission is 7.3 M fi .
The 13 CO mass within \Sigma2:7 km s \Gamma1 of v LSR combined with the CO outflow mass (jv LSR j ?
2:7/kms) represents a lower limit to the total cloud mass of ¸ 300 M fi . Approximately 18% of the
molecular cloud material is participating in the high­velocity outflow from I20126.
4. DISCUSSION
The I20126 CO outflow, with a mass of approximately 50 M fi , appears to be bounded on most
sides by higher density gas traced by C 18 O emission. The large­scale molecular flow has a N­S axis
while the near­infrared reflection nebula is oriented NW­SE along the jet axis. Shocks identified

-- 10 --
by H 2 and [SII] emission knots trace the jet close to the source and then rotate to a more N­S
orientation farther from the source following the edges of the CO flow.
4.1. The driving source of the molecular outflow
Two protostars spaced by about 1 00 (1500 AU) apart lie near the center of the I20126 outflow
(Hofner et al. 1999). Each source produces a separate ionized outflow with a NW­SE orientation.
We detect 3 mm, & 1 mm dust continuum centered on the northern source (hereafter I20126 N), and
it is associated with H 2 O masers along the ionized jet axis(Tofani et al. 1995). I20126 N is near the
center of a massive, ¸ 1700 AU diameter disk detected in CH 3 CN(5--4) (C99) and 7 mm continuum
emission (Hofner et al. 1999). The disk has a position angle of approximately 46 ffi which is nearly
perpendicular to the ionized jet (P.A. 117 ffi ). I20126 N is also clearly the center of the outflow
activity seen in H 2 , SiO, NH 3 , and CO(J=7--6) (Tofani et al. 1995, C97, Zhang et al. 1998, 1999,
C99, Hofner et al. 1999, Kawamura et al. 1999). The southern ionized outflow (hereafter I20126 S)
is weaker and there are no other signs of accretion disk or strong outflow activity centered on this
position. Our 3 mm continuum image shows that there is warm dust emission coincident with
I20126 S (center panel in Fig. 6); consistent with the presence of an embedded protostar. Saraceno
et al. (1996) have shown that the millimeter flux is proportional to the bolometric luminosity of
the embedded protostar. The lack of detectable 1 mm or 7 mm continuum emission and weaker
3 mm emission from I20126 S thus implies that the source itself is probably of lower luminosity
than I20126 N. The mass outflow rate from I20126 S is also expected to be lower since outflow
energetics ( —
M f , F and mechanical luminosity in the flow) scale with source luminosity (Cabrit
& Bertout 1992, Shepherd & Churchwell 1996, and references therein). The scenario of a binary
system with the primary having a more energetic outflow and more massive circumstellar disk
with, presumably, a higher disk accretion rate is consistent with what is found for many lower
mass binaries: circumprimary disks appear to have longer lifetimes and higher accretion rates than
circumsecondary disks (Mathieu et al. 2000, and references therein) and this result is supported
by numerical simulations of binary disk evolution (Lubow & Artymowicz 2000, and references
therein). This comparison indicates that I20126 N is probably dominating the outflow energetics
and luminosity.
A noteworthy property of the I20126 binary system is that both sources produce an ionized
jet with approximately the same position angle of ¸ 117 ffi . The orientation matches that of the
molecular jet, however, the jet axis is rotated by approximately 54 ffi relative to the CO outflow axis.
No other continuum source that could drive the massive molecular outflow has been detected in
the entire mosaiced field at 1 or 3 mm. Nor has any other potential driving source been detected
in 7 mm or 3.6 cm continuum emission (Hofner et al. 1999). Thus, it appears that the 53 M fi
flow with a N­S orientation is most likely powered by one or both young stellar objects in the
I20126 binary system even though the jet orientation is more NW­SE. In the remaining discussion
we assume, that the massive N­S outflow is dominated by the I20126 N protostar and that the

-- 11 --
contribution of the I20126 S flow to the total observed flow energetics is relatively minor. This
assumption should not introduce a significant error since a low­mass star of the same kinematic
age as the I20126 outflow (6:4 \Theta 10 4 years) with an average —
M f of 10 \Gamma7 to 10 \Gamma6 M fi yr \Gamma1 would
produce an outflow mass of only 0.006 to 0.06 M fi . This would contribute less than 0.1% to the
total mass of the flow and would not likely be distinguished within the energetic lobes of a 53 M fi
outflow.
The mass outflow rate —
M f ¸ 8:1 \Theta 10 \Gamma4 M fi yr \Gamma1 provides a means to obtain a rough estimate
of the expected bolometric luminosity of I20126 N (Shepherd & Churchwell 1996). The measured
—
M f implies an L bol of the driving star of about 1300 L fi which corresponds to a B3.5 spectral
type (Thompson 1984). In comparison, the bolometric luminosity determined from the SED is
¸ 10 4 L fi . The large difference between L bol (total) and the expected L bol of the driving source
supports the conclusion by C99 that most of the luminosity in the system may be due to accretion.
Thus, the outflow parameters are consistent with a scenario in which a mid to early­B type star
with L bol ¸ 1300 L fi is driving the flow while most of the luminosity is produced by accretion.
4.2. The nature of the jet
We suggest that the Hff and [SII] emission arc and central knot trace shocks associated with
a high­velocity component of the blue­shifted outflow that is breaking out of the near­side of the
cloud. The combined H 2 and [SII]­bright knots display an S­shaped symmetry about the central
star. The three inner­most H 2 knots lie along the SiO jet axis while the outer H 2 and [SII] knots
show a rotation of approximately 45 ffi relative to the jet axis. One possible explanation for this
is that the entire lobe could be filled with strong, shock­excited emission and the observed knots
are simply regions in which the extinction is relatively low. However, it is not likely that such
well­defined symmetry in the knot placement would exist if they were simply caused by random
extinction. Further, to explain the disparate orientations of the N­S CO flow and NW­SE jet
produced by a single source, one would have to invoke the special and somewhat unlikely geometry
that the CO outflow must be redirected by dense cloud material in a symmetric manner for both
the red and blue­shifted lobes. Thus, although extinction is clearly an important factor which
determines whether visible or infrared emission is present, it does not appear to be the primary
cause of the S­shaped symmetry in the knots.
A more promising explanation is that the H 2 and [SII] knots trace the working surfaces of the
I20126 N jet, as is the case for many outflows from lower mass young stellar objects (e.g. Davis
et al. 1994, Bally et al. 1995; Zinnecker et al. 1997). H 2 is also often observed as bright bow
shocks at the ends of jet flow segments and as fainter emission along the flow axis. Spectroscopic
studies of the kinematics of H 2 and CO indicate that this happens mostly via prompt entrainment
in the leading bows and to some extent via turbulent entrainment in shear layers along the flow axis
(Eisl¨offel 1997, and references therein). Although the situation is less clear for massive outflows
because the morphology is often more complex, it appears that H 2 may trace a steady jet in the

-- 12 --
W75 N flow (Davis et al. 1998). Assuming that the H 2 traces the path of the I20126 N jet, then
the jet appears to be precessing through an angle of about 45 ffi .
Independent evidence that supports a large jet precession angle in I20126 N comes from
estimates of the inclination angle of the flow. A model of the SiO jet shows that the inclination
angle is approximately 10 ffi with respect to the plane of the sky (C99). The kinematic age of the jet
is only about 2000 yrs old; thus i = 10 ffi provides an estimate of the current inclination of the jet.
In contrast, the CO outflow with i ¸ 45 ffi represents the time­averaged inclination angle over the
6 \Theta 10 4 yr life time of the molecular flow. Assuming the lobes are roughly symmetric about the flow
axis, the inclination angle of the CO flow could not be significantly less than 40 ffi without producing
overlapping red and blue­shifted emission in each outflow lobe. The difference of more than 30 ffi
between the inclination of the young jet and the time­averaged molecular flow is consistent with
the presence of a strongly precessing jet.
4.3. What can cause a disk/jet system to precess by 45 ffi ?
Some outflows from low­mass young stellar objects show ``wiggling'' knots which can be
interpreted as resulting from a precessing jet. The angle of jet precession is generally less than
about 10 ffi (e.g. Bally et al. 1996, Eisl¨offel & Mundt 1997, Gueth & Guilloteau 1999) although
a jet precession angle of ¸ 40 ffi is observed from the low luminosity source in L1228 (Bally et al
1995) and a jet with a 45 ffi precession angle is observed in the from the Herbig Ae/Be star PV Ceph
(Reipurth, et al. 1997, Gomez, et al. 1997). Fendt & Zinnecker (1998) summarize proposed bending
mechanisms of protostellar jets seen in low­mass flows. They conclude that two mechanisms, in
particular, can create jet precession angles ¸ 1 ffi to 2 ffi : precession due to a binary companion in a
coplanar orbit and Lorentz forces between the electric current in the ionized jet and an interstellar
magnetic field. However, precession angles greater than a few degrees cannot be generated with
these mechanisms.
For more luminous systems, significantly larger precession angles (up to 90 ffi ) could be generated
through radiation­induced warping of protostellar accretion disks if the disk is illuminated by a
sufficiently strong central radiation source (Armitage & Pringle 1997). Dramatic changes in the jet
orientation could also be caused by anisotropic accretion events, and precession could be induced
by tidal interactions between the disk from which the outflow originates and a close companion star
in a noncoplanar orbit (e.g. Terquem et al. 1999 and references therein). Here, we consider each
possibility to determine if the scaled parameters of a massive star and cloud core would create a
larger jet precession angle.
Radiative­induced warping
If the accretion disk is significantly warped near the star where the jet is produced, and
the jet remains perpendicular to the surface of the warped disk, then the jet would precess as
the system rotates. To examine whether the I20126 N disk is likely to be unstable to radiative

-- 13 --
warping, we follow the method of Armitage & Pringle (1997). We assume the disk is optically
thick, geometrically thin, and in Keplerian rotation. In general, if the luminosity is mostly derived
from accretion within the disk and boundary layer, then the disk is expected to be stable against
warping. However, if the mass flux is a strongly decreasing function of radius, the outer parts of
the disk could be warped by the stronger radiation emitted from the inner disk. In the extreme
case, we assume all luminosity is generated at or near the surface of the star. The critical radius
beyond which the disk is unstable to warping is given by
R crit = 3j 2
` M ?
M fi
' / —
M acc
10 \Gamma9 M fi yr \Gamma1
! 2 ` L ?
10 L fi
' \Gamma2
AU; (1)
where j is the ratio of the (R,z) component of the disk viscosity to the (R,OE) component of the disk
viscosity and —
M acc is the mass accretion rate. We assume j = 1 to obtain a rough estimate although
j could be an order of magnitude different if the turbulence in the disk is anisotropic (Pringle 1992).
If j = 10, then our estimate of R crit would increase by a factor of 100. For I20126 N, L ? ú 10 4 L fi ,
M ? ú 10 M fi , and we take —
M acc ¸ 10 \Gamma3 M fi yr \Gamma1 (C99). We find R crit ¸ 3 \Theta 10 7 AU. For disk radii
less than R crit , the disk is expected to be stable to warping. This is significantly larger than the
expected disk radius of ¸ 800 AU, thus, radiation does not appear to be able to warp the inner
disk of I20126 N and cause the jet to precess.
Anisotropic accretion events
Young, early B­type stars are usually surrounded by a dense cluster of lower mass stars (Testi
et al. 1999). Thus, it is likely that I20126 N formed in the presence of lower mass condensations
which could have evolved to be lower mass stars or they could have merged or interacted with the
more massive I20126 N protostar (see, e.g., Bonnell et al. 1998, Stahler et al. 2000, and references
therein). If a collision between a condensation and a massive YSO disk occurs during an anisotropic
accretion event, then it is possible the disk could be re­oriented. To obtain a rough estimate of the
angular momentum likely to reside in a disk around I20126 N, we assume the disk has a radius of
800 AU and is in Keplerian rotation about a 10 M fi central star with a disk mass of ¸ 4 M fi (roughly
10% of the mass detected in thermal dust emission). The total angular momentum in the disk would
be ¸ 2:3 \Theta 10 55 g cm 2 s \Gamma1 . In a dramatic scenario, we consider a 1 M fi condensation impacting the
disk at a radius of 800 AU and velocity 5 km s \Gamma1 . The angular momentum of the impactor would
be L impact ¸ 1:2 \Theta 10 55 sin OE g cm 2 s \Gamma1 where OE is the angle between the impactor trajectory and
the disk plane. If OE = 90 ffi and the impactor angular momentum is totally absorbed by the disk,
the disk angular momentum vector would be re­oriented by roughly ` = tan \Gamma1
i L impact
L disk
j
¸ 30 ffi .
Thus, an anisotropic accretion event of this magnitude could explain the large jet precession angle
observed in the I20126 N outflow.
Tidal interactions due to a companion in a non­coplanar orbit
The secondary star in a binary YSO system could produce tidal forces that act to truncate and
distort the disk (Lubow & Artymowicz 2000) and it can cause the circumprimary disk to precess
about the orbital axis (e.g. Terquem et al. 1999). The exact dynamics of the I20126 N disk/jet
precession will depend upon the stellar masses, the orbital radius of the binary, the circumstellar

-- 14 --
and circumbinary disk masses and radii, as well as the period of precession. Making some very
simplistic assumptions, our data allow us to estimate if this scenario can cause a disk precession of
order of tens of degrees. Assuming the disk precesses as a rigid body and the disk surface density
is uniform, the precession frequency of a Keplerian disk can be expressed as
! p = \Gamma
15
32
M s
M p
` R
D
' 3
cos ffi
s
G M p
R 3 (1 \Gamma e 2 ) \Gamma3=2 ; (2)
where M s is the mass of the secondary, M p is the mass of the primary, R is the disk radius, D is
the semi­major axis of the binary, e is the orbital eccentricity, and ffi is the angle between the disk
plane and the orbit of the secondary (Terquem et al. 1999).
The morphology of the H 2 and [SII] knots suggest that less than or on the order of half
a precession period has been completed over the 6 \Theta 10 4 yr age of the outflow. Thus, as a
rough estimate, the period of the precession is approximately 1:2 \Theta 10 5 yrs which corresponds
to a precession frequency of \Gamma1:6 \Theta 10 \Gamma12 s \Gamma1 . We take R = 800 AU (C97), D = 1400 AU
(Hofner et al. 1999), and M p ¸ 10 M fi (corresponding to a mid to early B star). We find that
M s cos ffi ú 9 (1 \Gamma e 2 ) 3=2 M fi . For the case of circular orbits (e = 0), the secondary must have
roughly the same mass as the primary to induce the observed jet precession. If the orbit is highly
eccentric with e = :75, then a significantly lower mass star M s cos ffi ú 2:6 M fi could produce
the observed precession in the I20126 N jet. Although these values can serve only as an order of
magnitude estimate, they are reasonable for what is expected in a cluster of massive stars. Thus,
the jet precession could be caused by the presence of the binary companion if it is in a non­coplanar
orbit.
More sensitive and higher­resolution 1 mm and 7 mm continuum observations to estimate the
circumstellar mass around the secondary would be useful to help constrain the system parameters
and verify which interpretation, anisotropic accretion or a binary in a non­coplanar orbit, is the
correct one.
5. SUMMARY
The massive molecular outflow from the I20126 region appears to be dominated by a single
early B­type protostar (I20126 N) with L ? ¸ 10 3 L fi while most of the total luminosity of the
system (¸ 10 4 L fi ) is likely produced by accretion. Approximately 1500 AU from I20126 N is
another source (I20126 S) with centimeter emission that is consistent with the presence of an
ionized jet, yet there is little indication of additional outflow activity. Based on the absence of
1 mm and 7 mm continuum emission, weaker 3 mm emission, and no detectable molecular outflow,
we conclude that I20126 S is less luminous than I20126 N. Thus, I20126 N is likely producing both
the N­S molecular flow and the inner jet oriented along a NW­SE axis. Shocks identified by H 2 and
[SII] emission knots appear to trace the jet as it precesses through a 45 ffi angle. Although I20126 S
does not appear to contribute significantly to the observed outflow kinematics, it may produce

-- 15 --
strong enough tidal forces that cause the circumprimary disk to precess about the orbital axis and,
in turn, cause the observed jet to precess by an angle of 45 ffi over the life time of flow. Alternatively,
jet precession could be induced by a dramatic, anisotropic accretion event which alters the angular
momentum vector of the disk. Additional observations are needed to determine which explanation
is correct.
Research at the Owens Valley Radio Observatory is supported by the National Science
Foundation through NSF grant number AST 96­13717. Star formation research at Owens Valley
is also supported by NASA's Origins of Solar Systems program, Grant NAGW­4030, and by the
Norris Planetary Origins Project. We would like to thank D. Koerner, B. Goodrich, and the Jet
Propulsion Laboratory MIRLIN team for their help with the Keck mid­infrared observations, A.
Sargent for critically reading an early version of the manuscript, and the referee, R. Bachiller, for
useful comments that improved this paper.
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-- 18 --
Table 1: Measured Positions of H 2 and [SII] Emission Knots
ff(B1950) ffi (B1950)
(h m s) ( ffi 0 00 )
H 2 --1 20 12 39.83 41 04 42.83
H 2 --2 20 12 40.38 41 04 25.00
H 2 --3 20 12 40.87 41 04 22.45
H 2 --4 20 12 41.67 41 04 17.40
H 2 --5 20 12 42.69 41 03 50.92
H 2 --6 20 12 42.91 41 03 33.56
Hff/[SII] 20 12 39.88 41 05 15.60
Table 2: IRAS 20126 Outflow Parameters
Distance 1.7 kpc
CO radius of red­shifted outflow 0.64 pc
CO radius of blue­shifted outflow 0.35 pc
Inclination Angle ¸ 45 ffi
Outflow Mass (corrected for optical depth):
Red­shifted 33.2 M fi
Blue­shifted 19.8 M fi
53.0 M fi
Dynamical time scale 6:4 \Theta 10 4 yr
—
M f 8:1 \Theta 10 \Gamma4 M fi yr \Gamma1
Momentum 403 M fi km s \Gamma1
Kinetic Energy 5:1 \Theta 10 46 ergs
Momemtum Supply Rate (Force) 6:0 \Theta 10 \Gamma3 M fi km s \Gamma1 yr \Gamma1
Mechanical Luminosity 5.9 L fi
Circumstellar Mass M (gas+dust)
1 52 M fi
1 Based on OVRO 3 mm continuum flux

-- 19 --
Figure Captions
Figure 1. CO red­shifted (thick lines) and blue­shifted (thin lines) emission contours from
34.9 km s \Gamma1 to 4.9 km s \Gamma1 and --10.9 km s \Gamma1 to --41.9 km s \Gamma1 , respectively. The maps have an RMS
of 1.23 Jy beam \Gamma1 km s \Gamma1 . Contours begin at 6.2 Jy beam \Gamma1 km s \Gamma1 (5oe). Blue­shifted contours
continue with spacings of 5oe while red­shifted contours have spacings of 10oe. The synthesized
beam, shown in the bottom right corner, is 7:53 00 \Theta 7:09 00 at P.A. \Gamma80:6 ffi . The position of the 3 mm
continuum peak is indicated by a black cross. Triangles represent the locations of H2 emission
knots (listed in Table 1). The diamond symbol above the blue­shifted outflow lobe represents the
location of the bright emission knot seen in Hff and [SII] emission. H 2 O maser positions identified
by Tofani et al. (1995) are within 2 00 of the continuum peak position and are not distinguishable
on this larger scale map. A scale size of 0.25 pc is represented by the bar in the upper right.
Figure 2. Top: First moment maps of the CO emission in the red and blue­shifted outflow
lobes. The color wedge on the right shows the velocity range. Symbols are the same as in Fig. 1.
Bottom: Hff and [SII] emission in color scale with the lowest CO red and blue­shifted emission
contour from Fig. 1. The boundary of the cloud is seen as a bright rim around the cloud core. The
CO outflow is seen in projection against the cloud core. The position of the 3 mm continuum peak
is indicated by a cross. A scale size of 0.5 pc is represented by the bar in the lower right.
Figure 3. CO channel maps at 1.3 km s \Gamma1 spectral resolution between \Gamma41:9 and 34.9 km s \Gamma1 .
The velocity is indicated in the upper left of each panel. The RMS is 0.12 Jy beam \Gamma1 and the
peak emission is 18.5 Jy beam \Gamma1 . In the top 20 and bottom 20 panels, contours are plotted from
\Gamma5; 5; 10; 15; 20; 25; 30; 40oe and continue with a spacing of 10oe. In the central 20 panels, contours
begin at \Gamma10; 10oe and continue with a spacing of 20oe. The last panel (velocity ­41.9 km s \Gamma1 ) shows
the synthesized beam in the bottom right corner (7:53 00 \Theta 7:09 00 at P.A. \Gamma80:6 ffi ). A scale size of
0.25 pc is represented by a bar in the lower left corner. In all maps, the cross shows the position of
the 3 mm continuum peak.
Figure 4. Left: K s band image of the bright bipolar reflection nebula centered on the millimeter
continuum source. Right: H 2 line image (with continuum subtracted) showing bright H 2 knots
within the bipolar nebula. An isolated knot is also present in the north­west and two knots are
located in the south­east. The position of the 3 mm continuum source is marked with a cross in
each image.
Figure 5. Mid­infrared images at 12.5 ¯m and 17.9 ¯m displayed with a grey scale square root
stretch (top). The 12.5 ¯m map RMS is 2.1 counts/pixel with contours of 5, 10, 15, 20, 30, 40,
50, and 60 \Theta RMS and the 17.9 ¯m map RMS is 20 counts/pixel with contours of 4, 6, 8, 10, 12,
14, 16, and 18 \Theta RMS. The three filled triangles represent H 2 O maser positions from Tofani et
al. (1995). The filled star shows the assumed position of the millimeter continuum peak. In the
bottom panel, 17.9 ¯m image contours are compared with the H 2 emission grey scale from Fig 3.
Figure 6. Continuum emission near the central source in I20126. The top panel shows 3 mm

-- 20 --
continuum emission using baselines ranging from 15 to 242 m. The map has an RMS of
1.7 mJy beam \Gamma1 . Contours begin at \Gamma2; 2; 4oe and continue with spacings of 2oe. The greyscale
is plotted on a linear scale from 1.7 mJy beam \Gamma1 to 3.47 mJy beam \Gamma1 . The synthesized beam
is 2:93 00 \Theta 2:53 00 at P.A. 78:4 ffi . The central panel shows the 3 mm continuum emission using the
longer baselines (70 to 484 m) to show the most compact emission distribution. The map has an
RMS of 2.1 mJy beam \Gamma1 . Contours begin at \Gamma2; 2; 3oe and continue with spacings of 1oe. The
greyscale is plotted on a linear scale from 2.1 mJy beam \Gamma1 to 2.06 mJy beam \Gamma1 . The synthesized
beam is 1:65 00 \Theta 1:16 00 at P.A. 73:3 ffi . The bottom panel shows the 1 mm continuum emission with
resolution 1:18 00 \Theta 0:81 00 at P.A. \Gamma51:2 ffi . The map has an RMS of 19.0 mJy beam \Gamma1 . Contours
begin at \Gamma3; 3; 5oe and continue with spacings of 2oe. The greyscale is plotted on a linear scale from
19 mJy beam \Gamma1 to 254.9 mJy beam \Gamma1 . Synthesized beams are shown in the bottom right corners
of each panel. A scale size of 5100 AU is represented by the bar in the lower left. The three filled
triangles near the continuum peak represent H 2 O maser positions from Tofani et al. (1995). In
the center panel, the two thick, parallel lines represent the size and orientation of the ionized jets
discovered by Hofner et al. (1999).
Figure 7. The spectral energy distribution for I20126 from 3.6 cm to 2.2 ¯m. The dashed line
represents our best­fit to the far­infrared and millimeter points with T d = 44 K and fi = 1:5.
Symbols are (roughly from left to right): filled pentagon--VLA, Hofner et al. (1999); arrow--VLA,
upper limits at 2 cm and 1.3 cm from Wilking et al. (1989) and Tofani et al. (1995), respectively;
\Theta--PdB, Cesaroni et al. (1999); star--OVRO, Wilking et al. (1989); open pentagon--OVRO, this
work; filled square--SCUBA, Cesaroni et al. (1999); open square--NRAO­12m, Walker et al. (1990);
open circle--IRAS detections at 100, 60 and 25 ¯m; filled circle--MAX, Cesaroni et al. (1999);
arrow--IRAS upper limit at 12 ¯m; and filled triangles--ARNICA, Cesaroni et al. (1997).
Figure 8. Comparison of the red and blue­shifted emission in CO and 13 CO with the C 18 O
integrated intensity. The CO contours in the left panel are from Fig. 1. 13 CO red­shifted (thick
lines) and blue­shifted (thin lines) emission contours are from 0.58 km s \Gamma1 to 6.02 km s \Gamma1 and
--7.58 km s \Gamma1 to --13.0 km s \Gamma1 , respectively. 13 CO maps have an RMS of 0.35 Jy beam \Gamma1 km s \Gamma1 ,
contours begin at 1.4 Jy beam \Gamma1 km s \Gamma1 (4oe) and continue with spacings of 0.7 Jy beam \Gamma1
km s \Gamma1 . The synthesized beam is 7:46 00 \Theta 6:03 00 at P.A. \Gamma9:8 ffi . C 18 O integrated intensity map
from \Gamma5:21 km s \Gamma1 to \Gamma1:8 km s \Gamma1 is shown in grey contours and greyscale. The map has an RMS
of 0.27 Jy beam \Gamma1 km s \Gamma1 , contours begin at 3oe and continue with spacings of 1oe. The greyscale is
plotted on a linear scale from 0.27 Jy beam \Gamma1 km s \Gamma1 to 2.54 Jy beam \Gamma1 km s \Gamma1 . The synthesized
beam is 7:42 00 \Theta 5:85 00 at P.A. \Gamma10:6 ffi . Synthesized beams for each map are shown in the bottom
right. Symbols are the same as in Fig. 1.
Figure 9. Top: 13 CO channel maps at 2.72 km s \Gamma1 spectral resolution between \Gamma15:7 and
8.74 km s \Gamma1 . The velocity is indicated in the upper left of each panel. The RMS is 65 mJy beam \Gamma1
and the peak emission is 4.16 Jy beam \Gamma1 . Contours begin at \Gamma5; 5oe and continue with a spacing
of 5oe. The last panel (velocity ­15.7 km s \Gamma1 ) shows the synthesized beam in the bottom right
corner (7:46 00 \Theta 6:03 00 at P.A. \Gamma9:8 ffi ). Bottom: C 18 O channel maps at 0.68 km s \Gamma1 spectral

-- 21 --
resolution between \Gamma6:57 and ­0.43 km s \Gamma1 . The RMS is 140 mJy beam \Gamma1 and the peak emission is
1.55 Jy beam \Gamma1 . Contours are plotted at \Gamma4; 4; 6; 8; 10; 12oe. The synthesized beam (7:42 00 \Theta 5:85 00
at P.A. \Gamma10:6 ffi ) is shown in the last panel. The cross shows the position of the 3 mm continuum
peak.
Figure 10. C 17 O integrated intensity map from \Gamma5:5 km s \Gamma1 to \Gamma3:5 km s \Gamma1 shown in contours
and greyscale. The map has an RMS of 80 mJy beam \Gamma1 km s \Gamma1 . Contours begin at \Gamma2; 2oe and
continue with spacings of 1oe. The greyscale is plotted on a linear scale from 160 mJy beam \Gamma1
km s \Gamma1 to 506 mJy beam \Gamma1 km s \Gamma1 . The synthesized beam, not shown, is 3:43 00 \Theta 3:09 00 at P.A.
\Gamma87:7 ffi . The position of the 3 mm continuum peak is indicated by a cross and the H 2 O maser
positions of C97 are denoted by solid triangles. A scale size of 5100 AU is represented by the bar
in the upper left.
Figure 11. CO (thin lines) and 13 CO (thick lines) spectra convolved with a 40 00 beam. The top
three panels show the spectra in the blue­shifted lobe (20 h 12 m 40 s 41 ffi 04 0 35 00 B1950), on the center
position (20 h 12 m 41 s 41 ffi 04 0 21 00 ), and in the red­shifted lobe (20 h 12 m 40 s 41 ffi 03 0 51 00 ), respectively.
The vertical dashed line shows v LSR = \Gamma3:5 km s \Gamma1 . The bottom panel shows the optical depth
derived from the spectra above.