Astronomy & Astrophysics manuscript no. organic˙paper˙accepted
(DOI: will be inserted by hand later)
September 29, 2018
arXiv:astro-ph/0406577v1 25 Jun 2004
Organic molecules in protoplanetary disks around T Tauri and
Herbig Ae stars
Wing-Fai Thi1,2,3 , Gerd-Jan van Zadelhoff1,4 , and Ewine F. van Dishoeck1
1
2
3
4
Leiden Observatory, P.O. Box 9513, NL 2300 RA, The Netherlands
Department of Physics and Astronomy, University College London, Gower Street, London WC1E 6BT, U.K.
Sterrenkundig Instituut Anton Pannekoek, Kruislaan 403, 1098 SJ Amsterdam, The Netherlands
Koninklijk Nederlands Meteorologisch Instituut, P. O. Box 201, 3730 AE De Bilt, The Netherlands
Abstract. The results of single-dish observations of low- and high-J transitions of selected molecules from protoplanetary disks around two T Tauri stars (LkCa 15 and TW Hya) and two Herbig Ae stars (HD 163296 and
MWC 480) are reported. Simple molecules such as CO, 13 CO, HCO+ , CN and HCN are detected. Several lines
of H2 CO are found toward the T Tauri star LkCa 15 but not in other objects. No CH3 OH has been detected
down to abundances of 10−9 − 10−8 with respect to H2 . SO and CS lines have been searched for without success. Line ratios indicate that the molecular emission arises from dense (106 –108 cm−3 ) and moderately warm
(T ∼ 20–40 K) intermediate height regions of the disk atmosphere between the midplane and the upper layer,
in accordance with predictions from models of the chemistry in disks. The sizes of the disks were estimated from
model fits to the 12 CO 3–2 line profiles. The abundances of most species are lower than in the envelope around
the solar-mass protostar IRAS 16293-2422. Freeze-out in the cold midplane and photodissociation by stellar and
interstellar ultraviolet photons in the upper layers are likely causes of the depletion. CN is strongly detected in all
disks, and the CN/HCN abundance ratio toward the Herbig Ae stars is even higher than that found in galactic
photon-dominated regions, testifying to the importance of photodissociation by radiation from the central object
in the upper layers. DCO+ is detected toward TW Hya, but not in other objects. The high inferred DCO+ /HCO+
ratio of ∼0.035 is consistent with models of the deuterium fractionation in disks which include strong depletion
of CO. The inferred ionization fraction in the intermediate height regions as deduced from HCO+ is at least
10−11 − 10−10 , comparable to that derived for the midplane from recent H2 D+ observations. Comparison with
the abundances found in cometary comae is made.
Key words. interstellar medium: molecules – circumstellar matter – stars: pre-main-sequence – Astrochemistry
1. Introduction
The protoplanetary disk phase constitutes a key period
in the evolution of matter between the young protostellar
and the mature planetary system stages. Before their incorporation into comets and large bodies, the gas and dust
could have participated in a complex chemistry within the
disk. Studies of the chemistry in disks are therefore important to quantify the chemical composition of protoplanetary material.
The chemical composition of the envelopes around young
protostars is now known with increasing detail thanks to
the combination of rapid advances in detectors and antenna technology and improved models (e.g., van Dishoeck
& Blake 1998; Langer et al. 2000). Part of this gas and
dust settles around the pre-main-sequence star in the form
Send offprint requests to: Ewine F. van Dishoeck, e-mail:
ewine@strw.leidenuniv.nl
of a disk, and after the collapse and accretion onto the
star cease, planets and comets can form by accumulating gaseous and solid material on timescales of a few
million years (e.g., Lissauer 1993; Beckwith & Sargent
1996; Wuchterl et al. 2000). Surveys from the near-infrared
to the millimeter wavelength range have shown that a
large fraction of 1–10 million year old Sun-like pre-mainsequence stars harbors a disk in Keplerian rotation (e.g.,
Beckwith et al. 1990). The masses of these disks (0.001 0.1 M⊙ ) is sufficient to form a few giant gaseous planets.
Single-dish and interferometric observations of molecular species other than CO are starting to reveal the
chemistry in disks around classical T Tauri stars (Dutrey,
Guilloteau & Guélin 1997; Kastner et al. 1997; Simon,
Dutrey & Guilloteau 2000; Duvert et al. 2000; van
Zadelhoff et al. 2001; Aikawa et al. 2003; Qi et al.
2003; Dartois et al. 2003; Kessler et al. 2003; Wilner
et al. 2003). The low-J rotational transitions of simple
molecules (HCN, CN, HNC, H2 CO, HCO+ , CS, ...) are
2
Thi et al.: Molecules in protoplanetary disks
detected, but their abundances relative to H2 are inferred
to be orders of magnitude lower than those observed in
dark clouds. The prevailing explanation of this depletion
involves a combination of freeze-out of the molecules on
grain surfaces in the cold midplane and their photodissociation by ultraviolet and/or X-rays in the upper atmosphere of disks (see Aikawa et al. 1999a, 2002; Bergin et al.
2003). The abundances are enhanced in the intermediate
height regions, which are warm enough for the molecules
to remain in the gas phase. Photodesorption induced by
ultraviolet radiation (Willacy & Langer 2000; Westley et
al. 1995) or X-rays (Najita, Bergin & Ullom 2001) can further populate the upper layers with molecules evaporated
from dust grains.
We present here the results of a survey of several lowand high-J molecular transitions observed toward two
classical T Tauri stars (LkCa 15 and TW Hya) and two
Herbig Ae stars (MWC 480 and HD 163296) using singledish telescopes. In particular, organic molecules such as
H2 CO, CH3 OH and HCN and deuterated species were
searched for. The comparison of the two types of objects
allows the influence of the color temperature of the radiation field on the chemistry to be studied. There are several
advantages in observing high-J transitions over the lowerJ ones. First, detections of CO J = 6 → 5 and H2 show the
presence of a warm upper surface layer in protoplanetary
disks whose temperature is higher than the freeze-out temperature of most volatile molecules (van Zadelhoff et al.
2001; Thi et al. 2001). Combined with the high densities in
disks, this allows the mid-J levels to be readily populated.
Models of flaring disks predict that the upper layer facing directly the radiation from the central star can extend
out to large radii (Chiang & Goldreich 1997; D’Alessio et
al. 1999). Second, by observing at higher frequencies with
single dish telescopes, the lines suffer less beam dilution
entailed by the small angular size of disks, typically 1–3′′
in radius, than at lower frequencies. Also, confusion with
any surrounding low-density cloud material is minimized.
The results for the different molecules are compared
to those found for protostellar objects, in particular the
solar-mass protostar IRAS 16293-2422, which is considered representative of the initial cloud from which the Sun
and the solar nebula were formed. This so-called Class
0 object (André et al. 2000) is younger than the protoplanetary disks studied here, only a few ×104 yr, and its
chemistry is particularly rich as shown by the number of
species found in surveys in the (sub)millimeter range (e.g.,
van Dishoeck et al. 1995; Ceccarelli et al. 2001; Schöier et
al. 2002, Cazaux et al. 2003 and references therein). The
similarities and differences in the chemical composition
between IRAS 16293-2422 and the protoplanetary disks
can be used to constrain the chemical models of disks.
At the other extreme, the results for disks can be compared with those found for objects in our solar system, in
particular comets. This will provide more insight into the
evolution of matter from the protoplanetary disk phase to
planetary systems. Unfortunately, the chemical composition of the large bodies in our solar system has changed
since their formation 4.6 Gyr ago. For example, solar radiation triggers photochemical reactions in the atmospheres
of planets, and the release of energy from the radioactive
decay of short-lived elements such as 26 Al causes solids
to melt. Comets, however, could have kept a record of the
chemical composition of the primitive solar nebula because
they spent much of their time in the cold outer region of
the Solar System (the Oort cloud) since their formation
(Irvine et al. 2000, Stern 2003). Comparison of cometary
D/H ratio and the CH3 OH abundances with those in disks
are particularly interesting.
This paper is organized as follows. In § 2, the characteristics of the observed objects are summarized. In § 3, the
observational details are provided. The results are given
in § 4 where a simple local thermodynamical equilibrium
(LTE) and statistical equilibrium analysis is performed. In
this section, we also derive several disk characteristics by
fitting the 12 CO 3–2 lines. In § 5, the molecular abundance
ratios are discussed. In particular, the CN/HCN ratio can
trace the photochemistry whereas the CO/HCO+ ratio is
a tracer of the fractional ionization. Finally, a discussion
on the D/H ratio in the disks compared with that found
in comets or other star-forming regions is presented (see
also van Dishoeck et al. 2003).
2. Objects
The sources were selected to have strong CO J = 3 → 2
fluxes and the highest number of molecular lines detected
in previous observations (Qi 2001; Thi et al. 2001, van
Zadelhoff et al. 2001). LkCa 15 is a solar mass T Tauri
star located in the outer regions of the Taurus cloud. Its
age is estimated to be ∼10 million years, although Simon
et al. (2000) argue for an age of only 3–5 million years.
LkCa 15 is surrounded by a disk whose mass is estimated
to be around 0.03 M⊙ , although a higher mass has been
obtained from the fitting of its spectral energy distribution
(SED) (Chiang et al. 2001). LkCa 15 is one of the strongest
millimeter emitting sources in the sample of T Tauri stars
surveyed by Beckwith et al. (1990) along with GG Tau
and DM Tau.
TW Hya forms part of a young association of stars
that has been discovered only recently and is located at
only ∼56 pc (Webb et al. 1999). TW Hya itself is a classical isolated T Tauri star with a high X-ray flux and a
large lithium abundance. Its large Hα equivalent width
is indicative of active disk accretion at a rate of ∼ 10−8
M⊙ yr−1 (Kastner et al. 2002). Despite its relatively high
age (∼15 Myr), TW Hya is surrounded by a disk of mass
∼ 3×10−2 M⊙ (Wilner et al. 2000) seen nearly face-on
(Weintraub et al. 1989; Krist et al. 2000; Zuckerman et al.
2000).
MWC 480 and HD 163296 were chosen to be representative of Herbig Ae stars. These two objects have the
strongest millimeter continuum emission, with disk masses
similar to those around the two T Tauri stars. All selected objects show gas in Keplerian rotation as revealed
Thi et al.: Molecules in protoplanetary disks
3
Table 1. Stellar characteristicsa
a
Star
SpT
LkCa 15
TW Hya
HD 163296
MWC 480
K7
K8Ve
A3Ve
A3ep+sh
α
(J2000)
04 39 17.8
11 01 51.91
17 56 21.26
04 58 46.27
δ
(J2000)
+22 21 03
−34 42 17.0
−21 57 19.5
+29 50 37.0
D
(pc)
140
56
122
131
M∗
(M⊙ )
0.8
1.0
2.4
2.2
Log(L∗ /L⊙ )
−0.27
−0.60
+1.41
+1.51
Age
(Myr)
11.7
7−15
6.0
4.6
See Thi et al. (2001) for references. The ages are highly uncertain (see also Simon et al. 2000).
Table 2. Disk characteristics
Star
LkCa 15
TW Hya
HD 163296
MWC 480
Disk Massa
(10−2 M⊙ )
3.3 ± 1.5
3.0 ± 2.0
6.5 ± 2.9
2.2 ± 1.0
Radius
(AU)
425
200
310
695
Diameter
(′′ )
6.2
7.0
5.0
10.4
Inclination
(◦ )
57±10
<10
32±5
30±5
Ref.
1
2, 3, 4, 5, 6
7
7
Total gas + dust mass computed from millimeter continuum flux using a dust opacity coefficient κλ =0.01(1.3mm/λ) cm2
g and assuming a standard gas-to-dust mass ratio of 100
References. (1) Qi et al. (2003); (2) Weintraub et al. (1989); (3) Krist et al. (2000); (4) Weinberger et al. (2002); (5) Wilner et
al. (2000); (6) Calvet et al. (2002); (7) Mannings et al. (1996)
a
−1
by CO interferometric observations (Qi 2001; Mannings &
Sargent 1997).
The stellar characteristics of the four objects are given
in Table 1 and the disk characteristics in Table 2. In this
paper, inclination is defined such that 0◦ is a disk viewed
pole-on and 90◦ edge-on. Detailed modeling of the SED
of these objects suggests that the disks are in a state of
dust settling, especially LkCa 15 (Chiang et al. 2001). The
objects were chosen to be isolated from any cloud material
to avoid contamination of the single-dish data.
3. Observations
The observations were performed between 1998 and 2000
with the James Clerk Maxwell Telescope (JCMT)1 located on Mauna Kea for the high-J transitions (850 µm
window) and with the 30-m telescope of the Institut de
Radioastronomie Millimétrique (IRAM) at Pico Veleta for
the lower J lines (1 to 3 mm). At both telescopes, the observations were acquired in the beam-switching mode with
a throw of 120′′ at IRAM and 180′′ at the JCMT in the
azimuth direction. The observations suffer from beam dilution owing to the small projected sizes of the disks of
the order of 5–10′′ in diameter, compared to the beam
size of the JCMT (14′′ at 330 GHz) and IRAM (11.3′′ at
220 GHz). The data were reduced and analyzed with the
SPECX, CLASS and in-house data reduction packages.
1
The James Clerk Maxwell Telescope is operated by the
Joint Astronomy Centre in Hilo, Hawaii on behalf of the
Particle Physics and Astronomy Research Council in the
United Kingdom, the National Research Council of Canada
and The Netherlands Organization for Scientific Research.
The JCMT observations made use of the dual polarization B3 receiver (315–373 GHz) and were obtained mostly
in November–December 1999. The antenna temperatures
were converted to main-beam temperatures using a beam
efficiency of ηmb =0.62, which was calibrated from observations of planets obtained by the staff at the telescope.
The data were obtained in single sideband mode with
the image side band lines reduced in intensity by about
13 dB (i.e. by a factor of ∼ 20). The sidebands were chosen to minimize the system temperature and to avoid any
unwanted emission in the other sideband. The integration
times range from 5 minutes for the bright 12 CO J = 3 → 2
lines to 8 hours for the faint lines to reach a r.m.s. noise
δTther of 10–20 mK after binning. The backend was the
Digital Autocorrelator Spectrometer (DAS) set at a resolution of ∼ 0.15-0.27 km s−1 (see Tables 3 and 4), and subsequently Hanning-smoothed to 0.3-0.6 km s−1 in spectra
where the signal-to-noise ratio is low. Pointing accuracy
and focus were regularly checked by observing planets,
and was found to be accurate to better than 3′′ rms at the
JCMT.
The estimated total r.m.s. error δT at the JCMT
associated with each line is given by the relation (e.g.,
Papadopoulos & Seaquist 1998):
δT
Tmb
tot
=
"
δT
Tmb
2
ther
+
δT
Tmb
2
drift
+
δT
Tmb
2
syst
#1/2
(1)
where the first term on the right-hand side of the relation
expresses the ratio between the thermal r.m.s. temperature and the main-beam peak temperature averaged over
Nch channels and with a baseline derived from Nbas chan-
4
Thi et al.: Molecules in protoplanetary disks
nels:
δT
Tmb
=
ther
δTther
TA∗
Nbas + Nch
Nbas Nch
1/2
(2)
where δTther is the thermal noise per channel and TA∗ is the
antenna temperature given in Tables 3 and 4. The beamswitching method gives extremely flat baselines such that
Nbas >> Nch and
δT
δTther 1
√
=
(3)
Tmb ther
TA∗
Nch
The data are subsequently binned to Nch = 2 or 4, and a
line detection is claimed whenever
δT
< 0.3
(4)
Tmb ther
observations related to the JCMT ′′ tracking error′′ problem 3 , unless the ion abundance is variable. We use only
the new 2004 data in our analysis. Note that the H13 CO+
and DCO+ data were taken only 1 week apart so that the
analysis of the DCO+ /HCO+ ratio should not suffer from
any potential long-term variability. Further monitoring of
the HCO+ line is warranted.
The IRAM-30m observations were carried out in
December 1998 using the 1–3 mm receivers. The weather
conditions were excellent. The three receivers and a
splitable correlator were used to observe simultaneously
lines at 1.3, 2 and 3 mm. The receivers were tuned singlesideband. Image band rejection was of the order of 10 dB.
Forward (Feff ) efficiencies were measured at the beginning
of each run and have been found to be consistent with
standard values. We measured Feff =0.9, 0.82 and 0.84 at
100, 150 and 230 GHz respectively. The derived beam efficiencies (ηmb = Beff /Feff ) are 0.57, 0.69, and 0.69 at 1, 2
and 3 mm respectively using main-beam efficiencies (Beff )
provided by the IRAM staff. The pointing and focusing
accuracy were regularly checked to ensure pointing errors
< 3′′ (r.m.s.) by observing planets and quasars. TW Hya
is unfortunately located too far south to observe with the
IRAM 30m telescope.
The second therm in Eq. 1 concerns the drift of
the temperature scale due to all errors of stochastic nature, and therefore includes any temperature variation of
the cold loads or any fluctuation of atmospheric opacity. Measurements of spectral standard sources just before or after the source observations allow an estimate of
this drift, which is generally found to be 10–15% and may
be up to 20–25% in difficult parts of the atmospheric window (e.g. H2 D+ amd N2 H+ lines) depending on the conditions. The lines for which calibration sources are available
from measurements by the JCMT staff are mentioned in
Table 52 .
The last term encompasses the systematic error, whose
main contributors are the uncertainty in the value of the
beam-efficiency and pointing errors. As noted above, the
pointing at the JCMT was found to be accurate to better than 3′′ at 345 GHz. Differences in beam efficiencies
and pointing should also be reflected in the spectral standard observations, which generally agree within 10-15% as
noted above. This last term is therefore estimated to contribute at most 10–20%. Taking into account all possible
sources of errors, the overall calibration uncertainties can
be as high as 30–40% for a line detected with 3σ in a difficult part of the spectrum, whereas it is of order 20–25%
for high S/N lines for which spectral standards have been
observed.
As discussed by van Zadelhoff et al. (2001), our HCO+
J = 4−3 intensity measured in 1999 is a factor of 3 weaker
than that obtained by Kastner et al. (1997). More recently,
we have re-observed the HCO+ line in May 2004 and find
intensities on two days which agree with those of Kastner
et al. within 10–20%. For comparison, 12 CO 3-2 spectra
taken in 1999, 2000 and 2004 are consistent within 10%
with the Kastner et al. results taken in 1995 with a different receiver, as are the HCN 4-3 and CN 3-2 results. Thus,
only the 1999 HCO+ result appears anomalously low, perhaps due to unusually large pointing errors during those
The measured antenna temperatures and thermal noise
per channel width are summarized in Tables 3 and 4. The
spectra are displayed in Fig. 2 to 4 on the main-beam temperature scale for the four sources. 12 CO J = 3 → 2 and
13
CO J = 3 → 2 are detected toward all objects. Apart
from TW Hya, the profiles of the 12 CO J = 3 → 2 spectra are double-peaked with peak separations of ∼ 2 km
s−1 for both three objects. The 12 CO J = 3 → 2 spectrum of MWC 480 shows a profile with slightly different
peak strengths. However, the level of asymmetry is not
signifcant compared to the noise. 12 CO J = 3 → 2 observations obtained with 30′′ offsets and position-switching
to an emission-free position are shown in Fig. 5 for the
four objects. The maps around LkCa 15, TW Hya and
MWC 480 confirm that these objects are isolated from
cloud material. The observations at offset positions from
HD 163296 show emission at velocities shifted compared
with the velocity of the star. The extinction to HD 163296
is sufficiently low that the extended low density emission is
unlikely to arise from a foreground cloud. The offset emission is only seen in 12 CO, not in 13 CO or other molecules.
Lines arising from high-J transitions require high critical
densities and are therefore not likely to come from a low
density cloud. A more complete discussion on the possible
contamination by foreground and/or background clouds is
given in Thi et al. (2001).
2
see also http://www.jach.hawaii.edu/JACpublic/JCMT/
Heterodyne observing/Standards/
3
http://www.jach.hawaii.edu/JCMT/Facility description/
Pointing/problem transit.html
4. Results
4.1. General characteristics
Thi et al.: Molecules in protoplanetary disks
5
Table 3. Antenna temperatures and noise per channel width for the T Tauri stars
Telescope
dv a
(km s−1 )
TA∗
(K)
a
δTther
(K)
LkCa 15
12
CO J = 2 → 1
12
CO J = 3 → 2
13
CO J = 3 → 2
C18 O J = 2 → 1
C18 O J = 3 → 2
HCO+ J = 4 → 3
H13 CO+ J = 4 → 3
DCO+ J = 5 → 4
CN J = 3 27 → 2 25
HCN J = 4 → 3
CS J = 7 → 6
H2 CO J = 212 → 111
H2 CO J = 303 → 202
H2 CO J = 322 → 221
H2 CO J = 312 → 211
H2 CO J = 515 → 414
CH3 OH J = 2K → 1K
CH3 OH J = 42 → 31 E+
CH3 OH J = 5K → 4K
N2 H+ J = 4 → 3
H2 D+ J = 110 → 111
IRAM30m
JCMT
JCMT
JCMT
JCMT
JCMT
JCMT
JCMT
JCMT
JCMT
JCMT
IRAM30m
IRAM30m
IRAM30m
IRAM30m
JCMT
IRAM30m
IRAM30m
IRAM30m
JCMT
JCMT
12
JCMT
JCMT
JCMT
JCMT
JCMT
JCMT
JCMT
JCMT
JCMT
JCMT
JCMT
JCMT
JCMT
JCMT
JCMT
JCMT
JCMT
JCMT
0.10
0.13
0.28
0.21
0.14
0.26
0.13
0.13
0.27
0.26
0.27
0.16
0.10
0.10
0.10
0.27
0.12
0.11
0.09
0.12
0.12
0.402
0.258
0.074
<0.084
<0.054
0.055
<0.078
<0.048
0.070
0.051
<0.038
0.021
0.025
<0.036
0.066
0.030
<0.014
<0.036
<0.036
<0.454
<0.372
(2σ)
(2σ)
(2σ)
(2σ)
(2σ)
0.159
0.051
0.026
0.042
0.027
0.018
0.039
0.024
0.018
0.022
0.019
0.008
0.021
0.018
0.016
0.016
0.007
0.018
0.018
0.227
0.186
(2σ)
(2σ)
(2σ)
(2σ)
(2σ)
(2σ)
(2σ)
(2σ)
(2σ)
0.066
0.241
0.046
0.230
0.087
0.016
0.023
0.075
0.080
0.028
0.049
0.034
0.044
0.037
0.023
0.355
0.298
0.101
(2σ)
(2σ)
(2σ)
(2σ)
(2σ)
(2σ)
TW Hya
CO J = 3 → 2b
12
CO J = 3 → 2c
13
CO J = 3 → 2
HCO+ J = 4 → 3b
HCO+ J = 4 → 3c
H13 CO+ J = 4 → 3
DCO+ J = 5 → 4
CN J = 3 72 → 2 25
HCN J = 4 → 3
H13 CN J = 4 → 3
HNC J = 4 → 3
DCN J = 5 → 4
H2 CO J = 312 → 211
H2 CO J = 515 → 414
CH3 OH J = 7K → 6K
N2 H+ J = 4 → 3
H2 D+ J = 110 → 111
SO J = 88 → 77
0.13
0.13
0.14
0.13
0.13
0.13
0.13
0.14
0.13
0.13
0.13
0.13
0.10
0.13
0.14
0.12
0.12
0.13
1.853
1.679
0.185
0.656
1.197
0.050
0.104
0.597
0.369
<0.056
<0.089
<0.068
<0.088
<0.074
<0.046
<0.710
<0.596
<0.202
a
dv is the spectral resolution of the data in km s−1 ; TA∗ is the antenna temperature (not corrected for beam efficiency) and
δTther is the thermal noise per dv channel. b Data taken in November 1999. c Data taken in May 2004.
High-J lines of various molecules are detected in
the disks. Lines with high signal-to-noise ratio toward
LkCa 15, HD 163296 and MWC 480 show a doublepeak structure corresponding to emission from a disk in
Keplerian rotation viewed under an inclination angle i
(Beckwith & Sargent 1993). The line profiles observed toward TW Hya are well fitted by a single gaussian, consistent with a disk seen almost face-on. The profiles show
no evidence of extended velocity wings characteristic of
molecular outflows in any of the objects. The velocity in-
tegrated main-beam temperatures for the four sources are
summarized in Table 6. This table includes the energy of
the upper level of the transitions, their critical densities
and frequencies, the telescope at which they were observed
and the beam
P size. The critical densities are defined as
ncr = Aul / l qul , where Aul is the Einstein A coefficient
of the transition u → l and qul the downward rate coefficient. They have been computed in the optically thin
limit at 100 K using the molecular data listed in Jansen et
al. (1994) and Jansen (1995). For optically thick lines, the
6
Thi et al.: Molecules in protoplanetary disks
Table 4. Same than Table 3 but for the Herbig Ae stars.
Telescope
dv
(km s−1 )
TA∗
(K)
Tther
(K)
HD 163296
12
CO J = 3 → 2
13
CO J = 3 → 2
HCO+ J = 4 → 3
CN J = 3 27 → 2 25
HCN J = 4 → 3
H2 CO J = 515 → 414
H2 CO J = 212 → 111
H2 CO J = 303 → 202
H2 CO J = 312 → 211
CH3 OH J = 2K → 1K
JCMT
JCMT
JCMT
JCMT
JCMT
JCMT
IRAM30m
IRAM30m
IRAM30m
IRAM30m
12
JCMT
JCMT
JCMT
JCMT
JCMT
JCMT
JCMT
JCMT
JCMT
JCMT
JCMT
JCMT
IRAM30m
IRAM30m
IRAM30m
0.23
0.28
0.26
0.27
0.26
0.27
0.17
0.11
0.10
0.12
0.920
0.225
0.128
0.115
<0.074
<0.204
<0.018
<0.030
<0.018
<0.014
(2σ)
(2σ)
(2σ)
(2σ)
(2σ)
(2σ)
0.130
0.052
0.030
0.018
0.037
0.102
0.009
0.015
0.009
0.007
MWC 480
CO J = 3 → 2
13
CO J = 3 → 2
C18 O J = 3 → 2
HCO+ J = 4 → 3
DCO+ J = 5 → 4
CN J = 3 27 → 2 25
CS J = 7 → 6
HCN J = 4 → 3
H2 CO J = 312 → 211
H2 CO J = 515 → 414
H2 D+ J = 110 → 111
N2 H+ J = 4 → 3
H2 CO J = 212 → 111
H2 CO J = 322 → 221
CH3 OH J = 42 → 31 E+
critical densities are lowered by roughly the optical depth
of the line.
The upper limits are computed assuming a main beam
temperature which corresponds to twice the r.m.s. noise
level in a 0.3 km s−1 bin and a line profile similar to that
derived from fitting the 13 CO J = 3 → 2 lines.
The ion HCO+ is detected in all sources. Toward
TW Hya, H13 CO+ is also seen, and the ratio of integrated
fluxes HCO+ /H13 CO+ of 24 is lower than the interstellar
isotopic ratio [12 C]/[13 C] ∼ 60, indicating that HCO+ line
is optically thick (van Zadelhoff et al. 2001). Emission from
the CN radical in the J = 3 → 2 line is surprisingly strong
in all objects compared to HCN J = 4 → 3. HNC J = 4 → 3
was only searched toward TW Hya and has not been detected. The line intensities for some species in this object
differ with previous observations by Kastner et al. (1997)
(see van Zadelhoff et al. 2001, §3). In general, the two
Herbig Ae stars display a less rich chemistry than the two
classical T Tauri stars. In particular, HCN is not detected
in either source in our observations. Qi (2001), however,
reports detection of HCN J = 1 → 0 toward MWC 480
with the Owens Valley Millimeter Array (OVRO).
Several lines of H2 CO are seen toward LkCa 15 with
the IRAM 30-m and JCMT, but not toward the other
three disks. Deep searches for various CH3 OH lines with
the IRAM-30m and JCMT down to very low noise levels
0.13
0.28
0.14
0.26
0.13
0.27
0.27
0.53
0.43
0.27
0.12
0.12
0.21
1.37
1.37
0.498
0.102
<0.062
0.052
<0.106
0.038
<0.040
<0.024
<0.030
<0.032
<0.090
<0.126
<0.016
<0.006
<0.006
(2σ)
(2σ)
(2σ)
(2σ)
(2σ)
(2σ)
(2σ)
(2σ)
(2σ)
(2σ)
(2σ)
0.073
0.016
0.031
0.017
0.053
0.017
0.020
0.012
0.015
0.016
0.045
0.063
0.008
0.003
0.003
did not yield any detections. So far, CH3 OH has only been
seen toward LkCa 15 through its J = 5K → 4K and J =
423 → 313 lines using OVRO (Qi 2001). This illustrates
the power of interferometers to detect minor species in
disks, because, with a beam diameter of 1′′ –2′′ , the fluxes
are much less beam-diluted, in addition to a potentially
larger collecting area when a great number of dishes is
available.
No CS J = 7 → 6 line nor lines of SO2 , some of which
occur fortuitously in other settings (e.g., near H2 CO 351
GHz), were detected toward LkCa 15. A deep limit on SO
was obtained toward TW Hya (Fig. 2 and Table 6).
Finally, DCO+ is detected for the first time in a disk,
as reported by van Dishoeck et al. (2003). The DCO+
J = 5 → 4 line is observed toward TW Hya with a strength
similar to that of H13 CO+ J = 4 → 3, but the line is not
detected toward LkCa 15 and MWC 480, where H13 CO+
is also not seen. There is a hint of a feature near the DCN
J = 5 → 4 line toward TW Hya, but it is offset by a few km
s−1 . Deeper integrations or interferometer data are needed
to confirm this. Note also the high critical density needed
to excite the DCN J = 5 → 4 line (∼ 4.8 107 cm−3 ), which
may make it more difficult to detect than DCO+ . The
ground-state line of ortho-H2 D+ at 372 GHz was searched
toward three sources (LkCa 15, TW Hya and MWC 480)
in a setting together with N2 H+ , but neither was detected.
Thi et al.: Molecules in protoplanetary disks
7
Table 5. Calibration sources used for the JCMT observations
Line
LkCa15
12
CO J = 3 → 2
N2071 IR
13
CO J = 3 → 2
W3(OH)
C18 O J = 2 → 1
C18 O J = 3 → 2
CRL618
W3(OH)
HCO+ J = 4 → 3
H13 CO+ J = 4 → 3
W3(OH)
IRC+10216
DCO+ J = 5 → 4
TW Hya
N2071 IR
IRC+10216
IRC+10216
...
...
HD 163296
IRC+10216
IRAS16293-2422
NGC6334I
G34.3
IRAS16293-2422
...
...
G34.3
CN J = 3 27 → 2 25
N1333IR4
IRC+10216
CRL618
N2071 IR
IRC+10216
N2071 IR
IRAS16293-2422
CRL618
HCN J = 4 → 3
CRL618
IRC+10216
CRL618
IRAS16293-2422
IRAS16293-2422
HNC J = 4 → 3
DCN J = 5 → 4
...
...
H13 CN J = 4 → 3
not available
CS J = 7 → 6
SO J = 88 → 77
H2 CO J = 312 → 211
H2 CO J = 515 → 414
not observed
...
...
CRL618
CH3 OH J = 7K → 6K
N2 H+ J = 4 → 3
H2 D+ J = 110 → 111
...
not available
not available
N2071 IR
N2071 IR
IRC+10216
...
IRAS16293-2422
...
...
N2071 IR
N2071 IR
IRC+10216
not observed
not available
not available
...
...
...
...
MWC 480
W3(OH)
W3(OH)
...
...
N1333IR4
W3(OH)
...
...
CRL618
G34.3
W3(OH)
CRL618
OMC1
...
...
...
...
...
...
...
IRAS16293-2422
not observed
...
...
CRL618
...
...
...
...
not available
not avaialble
Note.– When a observation is carried out over more than one run, different standard sources are used depending on their
availability. The entry ... indicates that the source was not observed in that particular line (cf. Table 6), so no standard was
needed.
Because of the poor atmosphere and higher receiver noise
at this frequency, the limits for both H2 D+ and N2 H+
are not very deep, except toward MWC 480. Recently,
Ceccarelli et al. (2004) have published the detection of
the H2 D+ 372 GHz line from the DM Tau disk using the
Caltech Submillimeter Observatory, together with a tentative feature from the TW Hya
R disk. Their integrated line
intensity toward TW Hya is TMB dV = 0.39 ± 0.12 K km
s−1 , compared with our 2σ limit of 0.20 K km s−1 . Taking
into account the smaller beam dilution in the JCMT beam
and the measurement uncertainties, the difference between
these two data sets is about a factor of two.
4.2. Disks properties and molecular abundances
4.2.1. Disk mean density
The mean density can be constrained from line ratios
of molecules with high dipole moments such as HCO+ ,
CN, HCN or H2 CO. A simple excitation analysis was
performed using an escape probability code described in
Jansen et al. (1994, 1995). The code computes the statistical equilibrium population of the rotational levels given
the kinetic temperature, volumn density and column density. Integrated temperatures of low-J transitions from Qi
(2001) were used to complement our high-J data. Both
sets of data were corrected for beam dilution by multiplying the observed velocity integrated main-beam temperatures by the size of the beam as listed in Table 6.
The analysis for LkCa 15 and TW Hya has been performed previously by van Zadelhoff et al. (2001) using
HCO+ and HCN, and takes both the radial and vertical
density structure of the disk into account. Here H2 CO is
also used as a diagnostic for LkCa 15 adopting the same
method. Consistent with their results, we find that the
densities in the regions probed by our observations range
from 106 to 108 cm−3 . This density refers to the region where the molecular lines are emitted. The
fractions of mass in a given density interval for various
disk models are shown in Fig. 3 of van Zadelhoff et al.
(2001). In all models (Chiang & Goldreich 1997; D’Alessio
8
Thi et al.: Molecules in protoplanetary disks
Fig. 1. Line profiles observed toward LkCa15. The dashed line indicates the velocity of the source. Note the different
antenna temperature scales for the different features.
et al. 1999; Bell et al. 1997), most of the gas is located in
the region of the disk where the density is greater than
106 cm−3 . Those densities are sufficient for most transitions studied here to be thermalized. We refer to the paper
Thi et al.: Molecules in protoplanetary disks
9
Fig. 2. Line profiles observed toward TW Hya. Note the different antenna temperature scales for the different features.
of van Zadelhoff et al. (2001) for a detailed discussion on
the disk models, the densities derived from line ratios and
the disk location where the lines are expected to become
optically thick.
4.2.2. Disk mean temperature
The mean kinetic temperatures are less well constrained:
the ratio 13 CO J = 3 → 2 / 13 CO J = 1 → 0 of 1.35 ±
0.4 suggests that Tkin ∼ 20–40 K for LkCa 15 in the region where the 13 CO emission originates, assuming that
both lines are optically thin (van Zadelhoff et al. 2001).
The bulk material where CO emits is therefore on average moderately warm and the density is high enough
that the level populations can be assumed to be thermalized for most cases. The ratios of 2.4 ± 0.7 for MWC 480
and 1.7 ± 0.5 for HD 163296 indicate similar temperature ranges. From the CO J = 6 → 5 / J = 3 → 2 ratios
presented in Thi et al. (2001), it is found that the upper
layers of disks have temperatures in the range 25–60 K.
The gas temperatures derived from the H2 data in Thi et
al. (2001) are slightly higher for disks around Herbig Ae
10
Thi et al.: Molecules in protoplanetary disks
Fig. 3. Line profiles observed toward MWC 480. Note the different antenna temperature scales for the different features.
stars than around T Tauri stars, as expected if the disks
are heated by the radiation from the central star. The
H2 CO J = 303 → 202 /J = 322 → 221 ratio is potentially a
good temperature indicator, but the J = 303 → 202 line
has only been detected toward LkCa 15. The limit on the
line ratio constrains the temperature to be below 200 K.
A mean temperature of 25 K is adopted in the remaining
parts of the paper.
4.2.3. Disk size
The disk sizes are important ingredients for comparing
the observed column densities with models. Since sizes
are notoriusly difficult to derive from low S/N interfer-
ometer maps, an attempt has been made to infer them
directly from our model profiles. Two methods have been
employed. First, since the 12 CO 3–2 emission line is optically thick, it probes the surface temperature profile of
the disk (van Zadelhoff et al. 2001). Using the method described by Dutrey et al. (1997) an estimate of the disk size
can be made from the 12 CO J = 3 → 2 lines:
Z
Tmb
#
2
2
π Rout
− Rin
cos i Ω−1
dv = Tex (ρδv)
a
D2
"
(5)
where Rin and Rout are the inner and outer radii, δv is
the local turbulent velocity (between 0.1 and 0.2 km −1 )
and ρ a geometrical factor of the order of 1.5. We adopt
Thi et al.: Molecules in protoplanetary disks
11
Fig. 4. Line profiles observed toward HD 163296. Note the different antenna temperature scales for the different
features.
here ρδv = 0.3 km s−1 , Rin =0 AU, Tex =25, 30, and 50 K
as the mean disk excitation temperature. The values for
the inclinations i, distances D (in AU) and beam sizes Ωa
are provided in Table 1 and 2. The derived disk sizes are
given in Table 7. Our estimates for Tex =30 K are similar to published values except for HD 163296, which we
find to have twice the size found by Mannings & Sargent
(1997), who measured it directly from their 12 CO 1–0 map.
Spectra of 12 CO 3–2 emission line were also generated
using a standard parametric disk model as described by,
e.g., Beckwith & Sargent (1993). The code uses a raytracing method and assumes that the population of the
rotational levels is in Local Thermodynamic Equilibrium.
All disks have a power-law density profile of the form
n(r) = n0 (r/1 AU)−2.5 . The exact value of n0 cannot be
constrained by fitting optically thick lines and we assume a
typical value of 5 1013 cm−3 at 1 AU in the mid-plane. The
disk is in hydrostatic equilibrium in the vertical direction.
The inner radius is set to 0.01 AU and is not a significant
parameter. The free parameters of this model are the excitation temperature Tin at 1 AU, the inclination i and
the outer radius Rout . For simplicity, the gas temperature
in the disk is assumed to have a radial profile power index of 0.5 and isothermal in the vertical direction. The
best fits are found using a downhill simplex method (e.g.,
Press et al. 1997). Figure 6 shows the observed spectra
and their best fits obtained with the parameters reported
in Table 7. The outer radii found by this ray-tracing model
are smaller than those from the optical depth model with
Tex = 30 K, which can be ascribed to additional contri-
butions from warmer gas at large radii not taken account
in the isothermal disk model. Note that Tin and Rout are
probably degenerate: Table 7 gives two sets of values for
LkCa 15 that can both fit the spectra. Only high signal-tonoise spatially resolved interferometer images can lift this
degeneracy. The larger outer radius (and smaller inner radius temperature) is adopted, which is closer to that found
by direct fitting of interferometric maps (Qi et al. 2003).
The inclinations are consistent with published values (see
Table 2).
4.3. Column densities
Given the sizes, molecular column densities can be derived
from the observed line strengths. Two additional assumptions need to be made: the excitation temperature and
the line optical depth. The line ratio analysis shows that
the lines arise from sufficiently high density regions (105 –
107 cm−3 ) that they can be assumed to be thermalized to
first order, although small deviations are expected in the
surface layers (see below). Therefore, a single excitation
temperature of 25 K is adopted to allow easy comparison
between the disks.
The optical depth can be estimated from the ratio
of lines from two isotopologues, assuming that the two
species have the same excitation temperature. Such data
are available for a few species and lines, and the results
are summarized in Table 8. It is seen that both the 12 CO
and H12 CO+ lines are very optically thick. An alternative method is to compare the size of the optically thick
12
Fig. 5.
Thi et al.: Molecules in protoplanetary disks
12
CO J = 3 → 2 maps toward TW Hya, LkCa 15, HD 163296 and MWC 480.
blackbody which accounts for the line flux to the actual
disk radius derived from the optically thick 12 CO line.
We adopt again the approach of Dutrey et al. (1997) and
rewrite Eq.(5) as follows:
sR
Tmb dv
1
×
Tex ρδv
π cos i
(6)
where θa is the main-beam diameter at half power in arcsec. For a Gaussian shape, the solid angle is given by
Ωa = 1.133θa2.
Assuming Tex =25 K, all derived radii are significantly
smaller than the CO disk sizes, except for CN J = 3 27 → 2 52 .
This would suggest that the lines from species other
than CO are less thick in the outer parts of the disks
(R >300 AU) to which our data are most sensitive, but
we consider the direct determination through the isotopologue ratios more reliable.
When the medium is slightly optically thick (τ <3) and
Tex >> TCMB with TCMB = 2.73 K, the column density
of the upper level Nu is given by:
8πkν 2 W Ωa + Ωs
τ
Nu =
(7)
hc3 Aul
Ωs
1 − e−τ
Rline (AU) = 106.4
θa
1′′
D
100 pc
where ν is the frequency
R of the transition, τ is the mean
optical depth, W = Tmb dv is the integrated line intensity expressed in K km s−1 , and Ωa and Ωs are the
telescope main-beam and the source solid angles, respectively. The ratio (Ωa + Ωs )/Ωs is the beam dilution factor.
Since the inferred disk sizes are much smaller than the
beam sizes, (Ωa + Ωs )/Ωs ≃ Ωa /Ωs . The Einstein Aul coefficient of the transition in units of s−1 coefficient is given
by:
64π 4 ν 3 Sµ2
(8)
Aul =
3hc3
gu
where µ is the dipole moment of the molecule in Debye,
S is the line strength, and gu = gJ gK gI is the statistical
weight of the upper level. Finally, the factor
τ
(9)
β −1 =
1 − e−τ
is the escape probability in the so-called Sobolev or Large
Velocity Gradient approximation.
The column density in level u is related to the total
column density N by:
Nu =
N
gJ gK gI e−Eu /Tex
Qrot (Tex )
(10)
Thi et al.: Molecules in protoplanetary disks
13
Table 6. Integrated line intensities
Eupper
(K)
na
crit
(cm−3 )
16.6
33.2
31.7
15.8
31.6
HCO+ J = 4 → 3
H13 CO+ J = 4 → 3
DCO+ J = 5 → 4
CN J = 3 27 → 2 52
HCN J = 4 → 3
H13 CN J = 4 → 3
HNC J = 4 → 3
DCN J = 5 → 4
CS J = 7 → 6
SO J = 88 → 77
H2 CO
H2 CO
H2 CO
H2 CO
H2 CO
H2 CO
Line
12
CO
CO
13
CO
C18 O
C18 O
12
J =2→1
J =3→2
J =3→2
J =2→1
J =3→2
J = 212 → 111
J = 303 → 202
J = 322 → 221
J = 312 → 211
J = 312 → 211
J = 515 → 414
CH3 OH
CH3 OH
CH3 OH
CH3 OH
J = 2K → 1K
J = 42 → 31 E+
J = 5K → 4K
J = 7K → 6K
N2 H+ J = 4 → 3
H2 D+ J = 110 → 111
R
Tmb dv (K km s−1 )
Beam
(′′ )
Cal.d
LkCa15
IRAM30m
JCMT
JCMT
JCMT
JCMT
10.7
13.7
14.3
21.5
14.3
...
yes
yes
yes
yes
1.82
1.17
0.39
<0.20
<0.14
...
1.98
0.24
...
...
...
3.78
0.94
...
...
...
2.88
0.57
...
...
356.734
346.998
360.169
340.248
354.506
345.339
362.630
362.046
342.883
344.310
JCMT
JCMT
JCMT
JCMT
JCMT
JCMT
JCMT
JCMT
JCMT
JCMT
13.2
13.6
13.1
13.9
13.3
13.6
13.0
13.0
14.0
13.7
yes
yes
yes
no
yes
no
no
yes
no
no
0.26
<0.13
<0.10
0.67
0.25
...
...
...
<0.08
...
1.26
0.07
0.11
1.14
0.49
<0.04
<0.05
<0.03
...
<0.10
1.10
...
...
0.95
<0.20
...
...
...
...
...
0.35
...
...
0.29
<0.07
...
...
...
<0.08
...
1.0(5)
4.7(5)
2.3(5)
4.5(5)
4.5(5)
1.7(6)
140.839
218.222
218.475
225.697
225.697
351.768
IRAM30m
IRAM30m
IRAM30m
IRAM30m
JCMT
JCMT
17.5
11.3
11.3
10.9
22.2
13.4
...
...
...
...
no
yes
0.17
0.14
<0.10
0.10
...
0.29
...
...
...
...
<0.05
<0.04
<0.10
<0.30
...
<0.30
...
<0.20
<0.40
...
<0.06
...
...
<0.09
6.9
45.4
34.8
65.0
2.6(3)c
3.7(4)
4.5(4)
1.3(5)
96.741
218.440
241.791
338.409
IRAM-30m
IRAM30m
IRAM30m
JCMT
25.4
11.3
10.2
13.9
...
...
...
yes
<0.05
<0.10
<0.10
...
...
...
...
<0.02
<0.03
...
<0.10
...
...
<0.20
...
...
44.7
104.3
4.4(6)
1.2(6)
372.672
372.421
JCMT
JCMT
12.7
12.7
no
no
<0.10
<0.10
<0.30
<0.20
...
...
<0.05
<0.05
ν
(GHz)
Telescope
2.7(3)
8.4(3)
8.4(3)
2.7(3)
8.4(3)
230.538
345.796
330.587
219.560
329.330
42.8
41.6
51.8
32.7
42.5
41.4
43.5
52.1
65.8
87.7
1.8(6)
1.8(6)
3.0(6)
6.0(6)
8.5(6)
8.5(6)
8.5(6)
4.8(7)b
2.9(6)
1.8(6)
21.9
21.0
68.1
33.5
33.5
62.5
TW Hya
HD 163296
MWC 480
Note.– The dots indicate not observed. When a line is not detected, a 2σ upper limit on Tmb in a 0.3 km s−1 bin is
computed and the same profile as the 13 CO J = 3 → 2 line is assumed. The beam size (HPBW) is computed for the
IRAM-30m using the fitting formula HPBW(′′ )=2460/frequency(GHz). a Unless specified, the critical densities are
taken from Jansen (1995) for Tkin =100 K assuming optically thin lines. a(b) means a × 10b . b Computed using the
collisional rate coefficients for HCN. c Derived assuming the collisional rate coefficients of Peng & Whiteoak (1993).
Observation of calibration sources before and/or after the object. More details are given in Table 5.
Table 7. Disk sizes from the
LkCa 15
LkCa 15d
TW Hya
HD 163296
MWC 480
a
b
c
d
12
d
CO J = 3 → 2 integrated intensities
Literaturea
425
425
200
310
695
Tex =25 K
620
...
238
778
722
Radius (AU)
Tex =30 K
559
...
215
701
650
Tex =50 K
422
...
162
530
492
i
57
57
3.5
65
30
Disk model
c
Rout
T (1 AU)b
290
170
450
100
165
140
680
180
400
170
See Table 2 for references
Gas temperature at 1 AU. The radius of disk is taken to be 0.01 AU.
Values adopted in subsequent analysis
Alternative fit to LkCa15 data
where Qrot (Tex ) is the rotational partition function, gJ is
the rotational statistical weight factor equal to 2J + 1 for
diatomic or linear molecules, gK is the K-level degeneracy, gI is the reduced nuclear spin degeneracy, and Eu is
the energy of the upper level expressed in K. For linear
molecules, gK =gI =1 for all levels. Formaldehyde, H2 CO,
is an asymmetric top molecule with I=1/2. Combining
the above equations, the total column density can be expressed as:
τ
1.67 × 1014
Eu /Tex
Q
(T
)e
N=
rot ex
νµ2 S
1 − e−τ
Z
Ωa
(11)
Tmb dv
Ωs
14
Thi et al.: Molecules in protoplanetary disks
Fig. 6. Observed (full lines) and simulated (dahes lines) 12 CO J = 3 → 2 spectra. The observations are normalized to
the peak values. The simulations are for a single temperature profile disk model (see Table 7 for parameters).
Table 8. Optical depths estimates from line ratios between isotopologues
12
CO J = 3 → 2
13
CO J = 3 → 2
HCO+ J = 4 → 3
H13 CO+ J = 4 → 3
HCN J = 4 → 3
H13 CN J = 4 → 3
LkCa15
TW Hya
26.8
0.44
19.4
< 0.32
0.75
0.012
8.3
0.14
3.8
0.06
< 5.1
< 0.085
HD 163296
18.9
0.31
...
...
...
...
MWC 480
14.9
0.25
...
...
...
...
Note.– We assume that the isotopologues have the same excitation temperature and that [12 C]/[13 C]=60.
This formula is similar to that of Blake et al. (1987) and
Turner (1991) but with the introduction of the escape
probability and beam dilution factor.
For linear molecules the line strength is equal to the rotational quantum number J. The rotational energy level
structure of the two linear molecules CN and HCN are
more complicated than those for CO. The spin of the unpaired electron for CN (S=1/2) and the nuclear spin of
14
N (I=1) lead to fine- and hyperfine splitting of the ro-
tational levels. The observed CN 340 GHz line is a blend
of three lines that account for ≃ 55% of all the hyperfine
lines arising from the level J=3 (Simon 1997). The line
strength is taken from Avery et al. (1992). An advantage
of the hyperfine splitting is that it decreases the optical
depth in each individual component. The HCN J = 4 → 3
line is also a blend of hyperfine lines but we assume that
all the flux is included in the observed line. Other constants used to derive the column densities are taken from
Thi et al.: Molecules in protoplanetary disks
existing catalogs (Pickett et al. 1998) and are summarized
in Thi (2002a). The rotational partition functions were
calculated using the formulae for each molecule in Gordy
& Cook (1984).
The radical CN and the molecule HCN have different
critical densities and the HCN J = 4 → 3 line may be subthermally excited in the upper layer, so that the inferred
N (CN)/N (HCN) ratio varies strongly with density. This
effect, which can be up to a factor of 2 in the CN/HCN
abundance ratio, has been corrected using the statistical
equilibrium calculations described above for the inferred
range of temperatures and densities. It should be noted
that this correction assumes that the CN and HCN lines
come from the same location inside the disks, which is
probably not the case. In disk models, CN peaks more toward the lower density surface layers than HCN because
CN is mostly formed by radical reactions and photodissociation of HCN (Aikawa et al. 2002). This effect would
lead to higher CN/HCN abundance ratios than presented
here.
Table 9 summarizes the beam-averaged column densities and upper-limits for the observed molecules, adopting
the disk sizes derived from the fits to the 12 CO J = 3 → 2
spectra using the isothermal disk model (see parameters
in rightmost columns of Table 7). A single excitation temperature Tex = Tkin = 25 K and an optical depth of τ = 1
are assumed for all lines. For optically thick lines with
τ ≫ 1, such as those of HCO+ and perhaps HCN, the derived column densities by these formulae are clearly lower
limits. Wherever available, column densities derived from
isotopic data have been used in these cases.
4.4. Molecular abundances
It is well known that disk masses derived from CO measurements are much lower than those obtained from millimeter continuum observations assuming a gas/dust ratio
of 100. Because CO is subject to photodissociation and
freeze-out, one cannot adopt the canonical CO abundance
of CO/H2 =10−4 found for molecular clouds; instead, the
disk masses Mdisk are assumed to be given by the millimeter continuum observations (see Thi et al. 2001, Table 2).
It should noted that, in the optically thin limit, the
abundances are independent of the disk size. The derived
abundances are summarized in Table 10. As noted above,
the abundances derived from the highly optically thick
HCO+ and HCN lines are likely to be underestimated by
up to an order of magnitude. For molecules that are detected in all four disks (CN, HCO+ ), the abundances vary
significantly from object to object. The non-detection of
HCN toward the Herbig Ae stars confirms the low abundances in these cases, although the high critical density of
the HCN J = 4 → 3 line may also play a role. The upper
limits are much lower in the case of TW Hya owing to the
small distance of this object and its narrow lines.
The derived abundances and depletion factors are
roughly consistent with the large range of values given
15
by van Zadelhoff et al. (2001), especially for their colder
models. A full comparison between the two studies is difficult since van Zadelhoff et al. performed more detailed
radiative transfer modeling with a varying temperature in
the vertical direction. Also, they adopted a smaller disk
radius of 200 AU compared with our value of 450 AU for
LkCa 15. This leads to higher abundances to reproduce
the same line flux, at least for optically thick lines.
For comparison, we have also re-derived the abundances in the disk around DM Tau, where many of the
same species have been detected by Dutrey et al. (1997).
Their tabulated velocity integrated flux densities have
been converted to velocity integrated main beam temperatures via the relation:
Z
Z
λ2 (cm) −1
Ωa
(12)
Tmb (K) dv = 10−23 × F (Jy) dv
2k
where F is the flux density in Jansky, k is the Boltzmann
constant in erg K−1 , λ the wavelength in cm, Ωa is the
main-beam solid angle and dv is in km s−1 . A total
disk mass of 0.018 M⊙ has been used, computed using
Equation (6) of Thi et al. (2001) and a continuum flux
of 110 mJy at 1.3 mm (Guilloteau & Dutrey 1998). The
abundances are reported in Table 10. Our new abundances
estimates are within a factor of 4 of those deduced by
Dutrey et al. (1997) who used a different method to derive their abundances.
The last column of Table 10 contains the abundances
found in the cold outer region of the protostellar envelope
of IRAS 16293–2422. The latter abundances seem to be
higher than those in disks by at least an order of magnitude, even taking into account the fact that the disk abundances may be underestimated because of optical depth
effects. A noticeable exception is CN, which has a higher
abundance in all four disks and in DM Tau than in IRAS
16293-2422.
5. Discussion
5.1. Molecular depletion
Two processes have been put forward to explain the low
molecular abundances in disks. First, in the disk midplane,
the dust temperature is so low (< 20 K) and the density
so high (nH > 109 cm−3 ) that most molecules including
CO are frozen onto the grain surfaces. This possibility is
supported by the detection of large amounts of solid CO at
infrared wavelengths in the disk around the younger class I
object CRBR 2422.8–3423 (Thi et al. 2002b). In this environment, surface chemistry can occur but the newlyformed species stay in the solid phase and thus remain
unobservable at millimeter wavelengths, except for a small
fraction which may be removed back in the gas phase by
non-thermal desorption processes such as cosmic-ray spot
heating.
Second, the photodissociation of molecules in the upper layers of protoplanetary disks by the ultraviolet radiation from the central star and from the ambient interstellar
16
Thi et al.: Molecules in protoplanetary disks
Table 9. Beam averaged column density of various molecules using the disk sizes listed in Table 7.
N (cm−2 )
LkCa15
TW Hya
HD 163296
MWC 480
Species
COa
13
CO
HCO+
HCO+
H13 CO+
DCO+
CN
HCN
H13 CN
HNC
DCN
CS
Transition
CO J = 3 → 2
J =3→2
J =4→3
H13 CO+ J = 4 → 3
J =4→3
J =5→4
J = 3 72 → 2 52
J =4→3
J =4→3
J =4→3
J =5→4
J =7→6
Rdisk =450 AU
1.9(16)
3.6(14)
3.3(11)
...
<1.5(11)
<2.9(11)
1.5(13)
1.8(12)
...
...
...
5.1(12)
Rdisk =165 AU
3.2(16)
5.5(14)
4.4(12)
1.2(13)
2.0(11)
4.4(11)
6.6(13)
9.2(12)
<4.8(11)
<1.4(12)
<4.0(10)
...
Rdisk =680 AU
3.5(16)
5.9(14)
9.4(11)
...
...
...
1.5(13)
<1.0(12)
...
...
...
...
Rdisk =400 AU
6.9(16)
1.2(15)
1.0(12)
...
...
...
1.5(13)
<1.2(12)
...
...
...
...
H2 CO
H2 CO
H2 CO
H2 CO
H2 CO
J = 212 → 111
J = 303 → 202
J = 322 → 221
J = 312 → 211
J = 515 → 414
5.1(12)
1.7(12)
<1.4(13)
7.1(11)
2.4(12)
...
...
...
...
<8.0(11)
<2.1(12)
<2.6(12)
<9.4(12)
<1.5(12)
<1.1(12)
<2.7(13)
...
...
<9.4(11)
<1.5(12)
CH3 OH
CH3 OH
CH3 OH
CH3 OH
J = 2K → 1K
J = 42 → 31 E+
J = 5K → 4K
J = 7K → 6K
<7.1(13)
<4.3(14)
<2.4(13)
...
...
...
...
<1.1(13)
<2.8(13)
...
...
...
...
<6.9(13)
...
...
N2 H+
H 2 D+
J =4→3
J = 110 → 111
<1.4(12)
<8.7(11)
<1.0(13)
<4.4(12)
...
...
<1.5(12)
<1.0(12)
SO
J = 88 → 77
...
<4.3(12)
...
...
12
13
Note.– The excitation temperature is assumed to be 25 K for all lines; uncertainties in the column densities are of the order of
30%, not including uncertainties in the disk size.
a(b) means a× 10b .
a 12
CO column density derived from 13 CO intensity assuming [12 C]/[13 C]=60.
medium can limit the lifetime of molecules. The ultraviolet flux from the central star can reach 104 times the
interstellar flux (Glassgold, Feigelson & Montmerle 2000).
Aikawa et al. (2002) and van Zadelhoff et al. (2003) have
modeled the chemistry in disks, taking these mechanisms
into account. Their models show that molecules are abundant in the intermediate height regions of disks, consistent with the derived temperature range (20–40 K) for
the emitting gas. According to the flaring disk model, this
intermediate region is located just below the warm upper
layer (T ≃ 100 K).
The molecular abundance distributions predicted by
the above chemical models including photodissociation
and freeze-out have been put into a 2-D radiative transfer code to compute the level populations using statistical equilibrium rather than LTE and to take the optical
depth effects properly into account. The resulting integrated fluxes can be compared directly with observations.
As shown by Aikawa et al. (2002) they differ by factors of
a few up to an order of magnitude, which indicates that
such models are to first order consistent with the data.
5.2. CN/HCN abundance ratio
Table 11 includes the CN/HCN abundance ratios derived for the disks. Compared with IRAS 16293-2422, the
CN/HCN ratio is more than two orders of magnitude
larger, and even compared with galactic PDRs, all disk
ratios are higher. The disk ratios may be overestimated
due to underestimate of HCN optical depth effects, but
the high values are a strong indication that photodissociation processes play a role in the upper layers of the
disks.
CN is particularly enhanced by photochemistry since
it is produced by radical reactions involving atomic C
and N in the upper layers as well as photodissociation
of HCN. Moreover, CN cannot be easily photodissociated
itself since very high energy photons (<1000 Å, > 12.4 eV)
are required to destroy the radical (van Dishoeck 1987).
The CN/HCN ratio appears to be higher in disks around
Herbig Ae stars than around T Tauri stars, although our
high ratio of > 11 for MWC 480 disagrees with the ratio
of ∼4 by Qi (2001) with OVRO. The disagreement be-
Thi et al.: Molecules in protoplanetary disks
17
Table 10. Beam-averaged molecular abundances with respect to H2 for the adopted disk sizes.
X = N/N (H2 )
a
b
c
Species
LkCa15
TW Hya
CO
HCO+
H13 CO+
DCO+
CN
HCN
H13 CN
HNC
DCN
CS
H2 CO
CH3 OH
N2 H+
H 2 D+
SO
3.4(-07)
5.6(-12)
<2.6(-12)
<2.31(-12)
2.4(-10)
3.1(-11)
...
...
...
<8.5(-11)
4.1(-11)
<3.7(-10)
<2.3(-11)
<1.5(-11)
...
5.7(-08)
2.2(-11)c
3.6(-13)
7.8(-13)
1.2(-10)
1.6(-11)
<8.4(-13)
<2.6(-12)
<7.1(-14)
...
<1.4(-12)
<1.9(-11)
<1.8(-11)
<7.8(-12)
<4.1(-11)
HD 163296
MWC 480
3.1(-07)
7.8(-12)
...
...
1.3(-10)
<9.1(-12)
...
...
...
...
<1.0(-11)
<1.5(-10)
...
...
...
6.9(-07)
1.0(-10)
...
...
1.4(-10)
<1.1(-11)
...
...
...
...
<1.4(-11)
<2.0(-09)
<1.5(-11)
<1.0(-11)
...
DM
This worka
9.6(-06)
7.4(-10)
<3.6(-11)
...
9.0(-09)
4.9(-10)
...
1.5(-10)
...
2.4(-10)
2.4(-10)
...
<5.0(-09)
...
...
Tau
Dutrey et al.
1.5(-05)
7.4(-10)
<3.6(-11)
...
3.2(-09)
5.5(-10)
...
2.4(-10)
...
3.3(-10)
5.0(-10)
...
<2.0(-10)
...
...
IRAS 16293-2422b
4.0(-05)
1.4(-09)
2.4(-11)
1.3(-11)
8.0(-11)
1.1(-09)
1.8(-11)
6.9(-11)
1.3(-11)
3.0(-09)
7.0(-10)
3.5(-10)
...
...
4.4(-09)
Re-analysis of data from Dutrey et al. (1997), see text
Outer envelope abundances from Schöier et al. (2002)
Value inferred from H13 CO+
Table 11. Relative molecular abundances
Source
LkCa 15
TW Hya
DM Taua
DM Taub
HD 163296
MWC 480
IRAS 16293-2422c
TMC-1d
Orion Bare
IC 63f
N (HCO+ )
N (CO)
1.6(-5)
3.8(-4)1
1.4(-4)2
5.3(-5)
7.6(-5)
2.7(-5)
1.5(-5)
3.6(-5)
1.0(-4)
2.0(-5)
2.7(-5)
N (CN)
N (DCO+ )
N (HCN)
7.9
7.1
N (HCO+ )
<0.411
0.035
5.8
18.4
>12.4
>11.7
0.07
1.5
3.8
0.7
References. (a) Dutrey et al. (1997); (b) from Table 7; (c) Schöier et al. (2002); (d) Ohishi et al. (1992); (e) Hogerheijde et al.
(1995); (f) Jansen et al. (1995).
Notes. 1 Value derived from H13 CO+ , assuming [12 C]/[13 C]=60. 2 Value derived from HCO+ .
tween the results of Qi (2001) and ours can be ascribed to
the fact that Qi (2001) had H13 CN data available to constrain the optical depth of the HCN line. Also, their 1–0
lines are less sensitive to the adopted disk density structure. In general the fluxes from MWC 480 for transitions
which have high critical densities are lower than those for
T Tauri stars, whereas the CO fluxes (with lower critical
densities) are higher. This may imply that the level populations are subthermal for the disks around Herbig Ae
stars.
Van Zadelhoff et al. (2003) have investigated the effects
of different UV radiation fields on the disk chemistry, focusing on T Tauri stars with and without excess UV emis-
sion. CN is clearly enhanced in the upper disk layers for
radiation fields without any excess UV emission owing to
its reduced photodissociation. When convolved with the
JCMT beams, however, the differences are difficult to discern: the HCN emission is nearly identical for the different radiation fields, whereas the CN emission varies by
only a factor of a few. Since T Tauri stars like TW Hya
have been observed to have excess UV emission (Costa
et al. 2000) —probably originating from a hot boundary
layer between the accretion disk and the star—, the difference in the CN/HCN chemistry with the Herbig Ae stars
may be smaller than thought on the basis of just the stellar spectra. Bergin et al. (2003) suggest that strong Lyα
18
Thi et al.: Molecules in protoplanetary disks
emission dominates the photodissociation rather than an
enhanced continuum flux. Since CN cannot be photodissociated by Lyα radiation but HCN can (Bergin et al. 2003,
van Zadelhoff et al. 2003), the CN/HCN ratio is naturally
enhanced.
Other chemical factors can also affect the CN/HCN ratio. Radicals such as CN are mainly destroyed by atomic
oxygen in the gas-phase and therefore a lower oxygen
abundance can increase the CN/HCN ratio. Since atomic
oxygen is a major coolant for the gas, a lower abundance will also maintain a higher mean kinetic temperature. Alternatively, the dust temperature could be in
the regime that HCN is frozen out but CN not because
the two molecules have very different desorption energies
(Edes (CN)=1510 K and Edes (HCN)=4170 K; Aikawa et
al. 1997).
Yet an alternative explanation for high CN abundances
is production by X-ray photons emitted from the active
atmosphere of T Tauri stars (e.g., Aikawa & Herbst 1999a;
Lepp & Dalgarno 1996). TW Hya is a particularly strong
X-ray emitter, with a measured X-ray flux 10 times higher
than the mean X-ray flux observed toward other T Tauri
stars (Kastner et al. 2002). In addition, H2 v = 1 → 0 S(1),
another diagnostic line of energetic events, has been observed toward this object (Weintraub, Kastner, & Bary
2000). TW Hya may however constitute a special case
since neither LkCa 15 nor DM Tau seems to show enhanced X-ray emission, yet they have a similar CN/HCN
ratio. Further observations of molecules in disks around
strong X-ray emitting pre-main-sequence stars are warranted to better constrain the contribution of X-rays on
the chemistry in disks.
5.3. HCO+ /CO abundance ratio
Table 11 compares the HCO+ /CO ratios found in the
disks with those found in a protostellar region (IRAS
16293-2422), a dark cloud (TMC-1) and two galactic
photon-dominated regions (PDRs) (Orion Bar and IC 63).
Within a factor of two, all values are very similar, except
for the TW Hya disk. It should be noted, however, that
except for TW Hya, the ratios in disks have been derived
from the optically thick HCO+ line and may therefore
be underestimates. Indeed, the ratio obtained using the
main HCO+ isotope for TW Hya is closer to that of the
other objects. Observations of H13 CO+ for all disks are
warranted to make definitive conclusions.
HCO+ is produced mainly by the gas phase reaction
+
H+
3 + CO → HCO + H2 . Its formation is increased by
enhanced ionization (e.g., by X-ray ionization to form H+
3
in addition to cosmic rays) and by enhanced depletion
(which also enhances H+
3 , see e.g., Rawlings et al. 1992).
The fact that all HCO+ abundances in disks are higher
than those found in normal clouds (after correction for
HCO+ optical depths) suggests that these processes may
play a large role in the intermediate warm disk layer where
both molecules are thought to exist. In this context it is
interesting to note that TW Hya has the largest depletion of CO and is also the most active X-ray emitter (see
below).
The derived HCO+ abundances provide a lower limit
to the ionization fraction in disks. The typical values of
10−11 − 10−10 are high enough for the magnetorotational
instability to occur and thus provide a source of turbulence
and mixing in the disk (e.g., Nomura 2002). Ceccarelli et
al. (2004) used their H2 D+ observations toward DM Tau
and TW Hya to derive an electron abundance of (4–7).
+
10−10 , assuming that H+
3 and H2 D are the most abundant ions. These values refer to the midplane of those disks
where the depletion of CO enhances the H2 D+ /H+
3 ratio
(see § 5.5), whereas our values apply to the intermediate
layer where most of the HCO+ emission arises.
5.4. H2 CO/CH3 OH
Lines of H2 CO have been detected toward only one source,
LkCa 15. CH3 OH is not detected in the single-dish observations of any disks, although it is seen in the OVRO
interferometer data toward LkCa 15 by Qi (2001), who
derived a CH3 OH column density of 7–20 × 1014 cm−2
compared with our upper limit of 9.4 × 1014 cm−2 . Our
upper limits for H2 CO and CH3 OH are derived from the
spectra with the lowest r.m.s. Methanol has been detected
from the class 0 protostar L1157 by Goldsmith, Langer &
Velusamy (1999), where its emission has been ascribed to a
circumstellar disk. However, this object is much younger
than those studied here and presumably has a different
physical structure and chemical history.
The H2 CO/CH3 OH abundance ratio of > 0.15 for
LkCa 15 is consistent with values found for embedded
YSOs (see Table 8 of Schöier et al. 2002 for IRAS
16293–2422 and van der Tak et al. 2000 for the case
of massive protostars). Heating of the disk, whether by
ultraviolet- or X-rays, should lead to strong ice evaporation and thus to enhanced gas-phase abundances for grainsurface products. Both CH3 OH and H2 CO have been detected in icy mantles, with the CH3 OH abundance varying strongly from source to source (Dartois et al. 1999,
Keane et al. 2001, Pontoppidan et al. 2003). For the few
sources for which both species have been seen, the solid
H2 CO/CH3 OH ratio varies from 0.1–1. Thus, the observed ratio in LkCa 15 could be consistent with grain
surface formation of both species. Since their absolute
abundances are much lower than typical ice mantle abundances of 10−6 with respect to H2 , this indicates that most
of the CH3 OH and H2 CO, if present, is frozen onto grains.
Deeper searches for both species in disks are warranted.
Protoplanetary disks are places where comets may
form and their volatile composition may provide constraints on their formation models. The CH3 OH abundance is known to vary significantly between comets. For
example, comet C/1999 H1 (Lee) shows a CO/CH3 OH
ratio around 1 whereas Hale-Bopp and Hyakutake have
ratios of 10 and 14 respectively (Biver et al. 2000).
Thi et al.: Molecules in protoplanetary disks
Comet Lee probably belongs to the so-called “methanolrich comets” group (Bockelée-Morvan, Brooke & Crovisier
1998; Davies et al. 1993). In addition, the measured CO
abundance is ∼1.8± 0.2% compared to H2 O, 5 times less
than found in Hale-Bopp. Alternatively, Mumma et al.
(2001) propose that Comet Lee has been heated sufficiently after its formation for CO to evaporate but not
CH3 OH, so that CH3 OH abundance is not enhanced
but rather CO is depleted. Mumma et al. (2001) notice
that the CH3 OH/H2 O and CO/H2 O ratios vary strongly
among comets coming from the giant-planets regions. The
picture is not complete since CO can be converted to
CO2 , whose abundance is high in interstellar ices (e.g.,
Ehrenfreund & Charnley 2000) but less well known in
comets (10% in comet 22P/Kopff, Crovisier et al. 1999).
Long period comets were probably formed in the
Jupiter-Saturn region (around 5–20 AU), whereas our
data are only sensitive to distances of more than 50 AU.
It would therefore be more relevant to compare the composition of protoplanetary disks to that of Kuiper Belt
Objects, which were formed beyond 50 AU in the solar
nebula. The chemical composition of Kuiper Belt Objects
is not well known (see Jewitt & Luu 2000), although observations show that comet nuclei and Kuiper Belt Objects
have different surface compositions (Luu & Jewitt 2002;
Jewitt 2002). The nature of Centaur objects is better understood. It is believed that Centaur objects were formed
beyond 50 AU and recently entered the planetary zone
with orbits crossing those of the outer planets. The best
studied Centaur object, 5145 Pholus, shows the presence
of CH3 OH although the exact amount is not well constrained (Cruikshank et al. 1998).
19
of deuterium fractionation can serve as a tracer of the
temperature history of the gas. The deuterium fractionation can be further enhanced by grain-surface formation
(Tielens 1983), although not for DCO+ /HCO+ . Recent
chemical models succeed in explaining the high fractionation observed here and in dark cloud cores (Rodgers &
Millar 1996; Roberts & Millar 2000, Tiné et al. 2000), but
only if significant freeze-out is included (Roberts et al.
2002, 2003). Our observed values are also close to those
found in disk models which include a realistic 2D temperature and density profile with freeze-out (Aikawa et al.
2002).
Table 2 in van Dishoeck et al. (2003) compares the
D/H ratio found in disks to typical values for the D/H
ratio in different protostellar and cometary environments.
The value found in disks is somewhat higher than that
in the low-mass protostellar envelope of IRAS 16293–
2422, but comparable to that seen in dark cloud cores.
DCO+ /HCO+ has not been observed in comets, but the
D/H ratios derived from DCN/HCN in pristine material in jets originating from below the comet surface is
found to be similar to that seen for DCO+ /HCO+ in the
TW Hya disk (Blake et al. 1999). Alternatively, Rodgers
& Charnley (2002) propose that the DCN and HCN seen
in these cometary jets are the photodestruction products
of large organic molecules or dust grains. In either case,
the D/H ratio of pristine icy material in comets is high.
The similarity suggests that either the gas is kept at low
temperatures as it is transported from pre-stellar clouds
to disks and eventually to comet-forming regions, or, alternatively, that the DCO+ /HCO+ ratio has been reset
by the chemistry in disks. Comparison with D/H ratios of
molecules which likely enter the disks as ices are needed
to distinguish these scenarios.
5.5. D/H ratio in a circumstellar disk
The D/H ratio in the TW Hya disk of 0.035 ±0.01 has
been derived from the H13 CO+ and DCO+ column density
ratios, assuming an isotopic ratio [12 C]/[13 C] of 60 (van
Dishoeck et al. 2003). Hints of H13 CN and DCN features
are seen in the TW Hya spectra, but neither of them is
definitely detected. Searches for other deuterated species
in the LkCa 15 disk, in particular DCN and HDO, are
reported by Kessler et al. (2003).
It should be kept in mind that our observations provide
only a value of the D/H ratio averaged over the entire disk.
Models of the DCO+ /HCO+ abundance ratio show that
it remains rather constant throughout the disk down to
a radius of 100 AU, however (Aikawa et al. 2002). All
values are significantly higher than the elemental [D]/[H]
abundance of 1.5×10−5 (Pettini & Bowen 2001; Moos et
al. 2002).
Theoretically, the amount of deuterium fractionation
in molecules depends on the gas kinetic temperature,
which drives the isotopic exchange reactions, and on the
cosmic ray ionization rate (Aikawa & Herbst 1999b). Also,
the abundance is enhanced if CO is significantly depleted
onto grains (Brown & Millar 1989). Thus, the amount
6. Conclusions
We surveyed low- and high-J transitions of simple organic
molecules in two classical T Tauri and two Herbig Ae stars.
Analysis of line ratios indicates that the emission comes
from a dense (nH = 106 –108 cm−3 ) and moderately warm
region (T ≃ 20–40 K). Detailed fits to the 12 CO 3–2 emission line profiles provide independent estimates of the sizes
of the disks.
Emission from the ion HCO+ and the radical CN are
particularly strong, indicating an active gas-phase chemistry in the surface layers of disks which is affected by UV
radiation from the central stars. H2 CO is detected in one
source but CH3 OH is not observed in any object in our
sample. In one source (TW Hya) the detection of DCO+
allows to constrain the DCO+ /HCO+ ratio to ∼0.035, a
value that is higher than that found in the envelopes of
low-mass protostars but comparable to that observed in
cold dark cores, where fractionation due to low temperature chemistry and CO freeze-out is important.
This work demonstrates that organic chemistry in
disks around low- and intermediate-mass pre-mainsequence stars can now be studied observationally. The
20
Thi et al.: Molecules in protoplanetary disks
detection of molecular species in disks is hampered, however, by the small sizes of disks compared with the actual
beams of single-dish telescopes. Moreover, because the total amount of material is small (few × 10−2 M⊙ ), the
observations are limited to the most abundant species. It
is likely that the chemistry in more tenuous disks, in which
the ultraviolet radiation can penetrate through the entire
disk, is different from that for our objects (e.g., Kamp &
Bertoldi 2000, Kamp et al. 2003). Although the outer disks
can be resolved by current millimeter interferometers, integration times are too long to do molecular line surveys
and the inner tens of AU are still out of reach. The detection of more complex and much less abundant molecules
in protoplanetary disks at different stages of evolution
awaits the availability of the Atacama Large Millimeter
Array (ALMA). Complementary infrared observations of
solid-state species along the line of sight of edge-on protoplanetary disks will help to constrain quantitatively the
level of depletion in the mid-plane of disks.
Acknowledgements. WFT thanks PPARC for a Postdoctoral
grant to UCL. This work was supported by a Spinoza grant
from the Netherlands Organization for Scientific Research
to EvD and a postdoctoral grant (614.041.005) to WFT.
We thank Remo Tilanus, Fred Baas, Michiel Hogerheijde,
Kirsten Knudsen-Kraiberg, Annemieke Boonman, and Peter
Papadopoulos, who have performed some of the JCMT observations in service; Geoff Blake, Charlie Qi and Jackie Kessler
for communicating their OVRO results prior to publication;
and Yuri Aikawa for fruitful discussions on disk models. We
acknowledge the IRAM staff at Granada for carrying part of
the observations in service mode.
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