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THE ASTROPHYSICAL JOURNAL, 479 307õ312, 1997 April 10
1997. The American Astronomical Society. All rights reserved. Printed in U.S.A.
(
HYDROGENATION OF INTERSTELLAR MOLECULES : A SURVEY FOR
METHYLENIMINE (CH 2 NH)
J. E. W. M. AND C. H.
DICKENS, IRVINE, DEVRIES
Five College Radio Astronomy Observatory, 619 Lederle GRC, University of Massachusetts, Amherst, MA 01003
dickens=fcrao1.phast.umass.edu, irvine=fcrao1.phast.umass.edu , devries=fcrao1.phast.umass.edu
AND
M. OHISHI
Nobeyama Radio Observatory, Nobeyama, Minamimaki, Minamisaku, Nagano 384­13, Japan ; ohishi=nro.nao.ac.jp
Received 1996 August 14 ; accepted 1996 November 4
ABSTRACT
Methylenimine has been convincingly detected for the ïrst time outside the Galactic center
(CH 2 NH)
as part of a study of the hydrogenation of interstellar molecules. We have observed transitions from
energy levels up to about 100 K above the ground state in the giant molecular clouds W51, Orion KL
and G34.3]0.15. In addition, was found at the ```` radical­ion peak îî on the quiescent ridge of
CH 2 NH
material in the Orion molecular cloud. The abundance ratio at the radical­ion peak
CH 2 NH/HCN
agrees with the predictions of recent gas­phase chemical models. This ratio is an order of magnitude
higher in the warmer cloud cores, suggesting additional production pathways for probably on
CH 2 NH,
interstellar grains.
Subject headings : ISM : abundances õ ISM : molecules
1. INTRODUCTION
At ïrst glance interstellar molecular clouds would seem
to provide an extremely reducing environment. Since
hydrogen is more than 1000 times as abundant as other
chemically reactive elements, one might expect the com­
pounds formed to be fully hydrogenated, e.g., CH 3 OH
(methanol), (ethanol), and (dimethyl
CH 3 CH 2 OH (CH 3 ) 2 O
ether). This is, however, manifestly not the case. Many
doubly and triply bonded molecules are found to have rea­
sonable abundances, particularly in the cold dark clouds.
Thus, in the classic dark cloud TMC­1 the abundances of
closed­shell molecular species with the linear heavy­element
backbone CCCN [HC 3 N (cyanoacetylene) : CH 2 CHCN
(vinyl (ethyl cyanide)] are in the ratio
cyanide) : CH 3 CH 2 CN
of 1 : 0.05 : \0.02 & Irvine Such a result is
(Minh 1991).
understandable from the kinetics of gas­phase reactions in
such an environment, since hydrogenation reactions often
have energy barriers that are difficult to overcome at the
temperatures in dark clouds (of order 10 K), even when such
reactions are energetically allowed Adams, &
(Herbst,
Smith In regions where active star formation is
1983).
taking place, with the resultant energy inputs to the inter­
stellar medium, more hydrogenated compounds are indeed
found. In the Orion ```` hot core,îî for example, the relative
abundances for are found
HC 3 N : CH 2 CHCN : CH 3 CH 2 CN
to be 1 : 1 : 7, almost the reverse of those in dark clouds (e.g.,
& Hjalmarson et al. The reasons
Irvine 1984 ; Blake 1987).
for this are thought to include the higher kinetic tem­
peratures in such regions, which permit a wider range of
gas­phase reactions to occur, and reactions on the surfaces
of solid particles (interstellar ```` grains îî) followed by the
release of such molecules back to the gas phase.
Consequently, the relative abundances for molecules that
are related by their degree of hydrogenation can provide
important information on the chemical processes taking
place in the interstellar medium, including the chemical role
of the interstellar grains. We are investigating such pro­
cesses by studying the abundances of species in the hydro­
genation series based on the cyanide radical ïve
(H n CN),
members of which are known to exist in the interstellar
medium : CN (n \ 0), HCN (n \ 1), (n \ 2),
H 2 CN CH 2 NH
(n \ 3), and (n \ 5). Among the closed­shell
CH 3 NH 2
species, HCN is well studied, but (methylenimine)
CH 2 NH
and (methylamine) have been securely detected
CH 3 NH 2
only in the Galactic center molecular cloud Sgr B2 (Godfrey
et al. et al.
1973 ; Kaifu 1974 ; Turner 1991).
Methylenimine is a prolate asymmetric rotor with com­
ponents of the electric dipole moment along both the a and
b molecular axes, with magnitudes 1.325 and 1.53 D, respec­
tively. The nitrogen nucleus produces electric quadrupole
hyperïne structure in low­lying transitions. The rotational
spectrum to 300 GHz has been tabulated by Kirchho+,
Johnson, & Lovas has been shown to be a
(1973). CH 2 NH
product of UV irradiation of model icy interstellar grain
mantles and a precursor of other complex organics which
may be present in cometary ices et al. It is
(Bernstein 1995).
also of potential prebiotic interest, since it can serve as a
precursor to glycinenitrile and glycine (e.g., Dickerson
We report here on the results of a survey for
1978). CH 2 NH
in both giant molecular clouds and in cold dark clouds.
2. OBSERVATIONS
A survey for interstellar methylenimine in Galactic
molecular clouds was carried out in 1995 December at the
National Radio Astronomy Observatory 12 m radio tele­
scope at Kitt Peak, Arizona. We used dual channel SIS
mixer receivers in the 2 and 1 mm regions with system
temperatures in the ranges 400õ850 K and 450õ650 K,
respectively, depending on weather, elevation, and fre­
quency. For the back ends, we simultaneously observed
with the hybrid autocorrelator and ïlter banks, but we con­
centrated our analysis on the autocorrelator data and used
the ïlter bank data as a check for consistency. For the ```` hot
core îî type sources (W51, Orion KL, G34.3]0.15, and
W49), we set up the hybrid autocorrelator spectrometer
with a 600 MHz total bandwidth with 781 kHz resolution.
307

308 DICKENS ET AL. Vol. 479
For the quiescent region, Ori 3N, we used 150 MHz band­
width with 195 kHz resolution, and for the dark clouds
(L134N and TMC­1), we used 37.5 MHz bandwidth with 49
kHz resolution.
All data were obtained using position switching 30@ in
azimuth. The HPBWs at 2 and 1 mm are about 37A and 26A,
respectively. We checked the pointing as often as possible
using continuum observations of the planets. However,
because of the paucity of pointing sources and the fact that
half of the dish was directly illuminated by the Sun during
daytime observing, we have conservatively put a 30% sys­
tematic uncertainty into our error analysis. In we
Table 1
list the transition frequency, the transition designation in
the nomenclature for an asymmetric rotor, the
J K~1K‘1
intrinsic line strength (S), the energy of the upper level of the
transition above the ground state the dipole moment
(E u ),
(k), and the beam size and corrected main­beam efficiency
at each observed frequency.
(g B )
The ïrst deïnitive detections of were made in
CH 2 NH
the two regions of massive­star formation, W51 and Orion
KL, with several transitions observed in each source. Time
constraints allowed the detection of only a single transition
in the giant molecular cloud G34.3]0.15, where a tentative
detection of an additional line has been reported by
et al. Line parameters are given in
Macdonald (1996). Table
and sample spectra are shown in For the two
2, Figure 1.
sources of this type where multiple transitions were
observed, we have plotted ```` rotation diagrams,îî from which
rotation temperatures and total column densities may be
found (e.g., et al. et al. These
Johansson 1984 ; Blake 1987).
are illustrated in Figure 2.
In addition, was observed to be present in the
CH 2 NH
quiescent gas at the so­called radical­ion peak along the
Orion ridge (here labeled Ori 3N), where the kinetic tem­
perature is about 25 K At this location, its
(Bergin 1995).
production is presumably by gas­phase ion­molecule chem­
istry, which suggests that may be rather widely
CH 2 NH
present in molecular clouds.
We did not detect toward the cold dark clouds
CH 2 NH
TMC­1 and L134N, which have been the subjects of numer­
ous chemical studies (e.g., Ohishi, & Kaifu
Irvine, 1991 ;
Irvine, & Kaifu et al. and refer­
Ohishi, 1992 ; Pratap 1994 ;
ences therein). This is likely to be a matter of low excitation
rather than low abundance, since the lowest energy tran­
sition that could be probed was for which the
1 11 õ0 00 ,
upper level is about 10 K above the ground state and for
which the critical density is about 106 cm~3. This line may
simply not be excited at the temperatures and densities that
TABLE 1
TRANSITION PARAMETERS FOR CH 2 NH
Frequency E u k Beam Size g B
(MHz) J K~ { 1K‘ 1
[ J K~ _ 1K‘ _ 1
A S (K) (D) (arcsec) (%)
172267.113 . . . . 2 1,1 õ2 0,2 2.460 17.5 1.530 34 66
225554.692 . . . . 1 1,1 õ0 0,0 1.000 10.8 1.530 26 53
245125.974 . . . . 4 1,4 õ3 1,3 3.749 37.3 1.325 24 51
250161.865 . . . . 7 1,6 õ7 0,7 5.906 97.2 1.530 24 50
251421.379 . . . . 6 0,6 õ5 1,5 3.053 64.1 1.530 24 50
255840.431 . . . . 4 2,3 õ3 2,2 3.000 62.2 1.325 23 49
NOTEõS is the intrinsic line strength ; is the upper level energy above the ground state k
E u
is the dipole moment along either the a or the b molecular axis and is the corrected
g B
main­beam efficiency (NRAO 12 m observing manual).
TABLE 2
OBSERVED LINE PARAMETERS FOR CH 2 NH
T R * (rms) V LSR *v Resolution
Source J K~ 1K‘ { 1
@ [ J K~ _ 1K‘ _ 1
A (mK) (km s~1) (km s~1) (kHz) Remarks
Orion KL . . . . . . 2 1,1 õ2 0,2 238 (50) 7.2 7.3 781
1 1,1 õ0 0,0 219 (10) 8.7 5.9 781 Blended
4 1,4 õ3 1,3 399 (13) 9.5 6.9 781 Blended
7 1,6 õ7 0,7 216 (30) 7.3 11.6 781
6 0,6 õ5 1,5 459 (38) 7.6 7.1 781 Blended
W51 . . . . . . . . . 1 1,1 õ0 0,0 121 (11) 55.6 10.0 781 Blended
4 1,4 õ3 1,3 140 (13) 59.3 9.5 781
6 0,6 õ5 1,5 57 (11) 56.4 7.5 781
4 2,3 õ3 2,2 29 (9) 59.9 6.9 781
G34.3]0.15 . . . . . 4 1,4 õ3 1,3 29 (8) 58.9 5.4 781
W49 . . . . . . . . . 1 1,1 õ0 0,0 . . . (10) . . . . . . 781
Ori 3N . . . . . . . . . 2 1,1 õ2 0,2 96 (30) 10.1 3.3 195
1 1,1 õ0 0,0 39 (12) 9.1 2.87 195
TMC­1 . . . . . . . . . 1 1,1 õ0 0,0 . . . (20) . . . . . . 49
L134N . . . . . . . . 1 1,1 õ0 0,0 . . . (25) . . . . . . 49
NOTEõLine parameters were derived by Gaussian ïtting. is the beam­chopperõcorrected antenna tem­
T R *
perature, and we use to estimate in the main beam ; the errors (in parentheses) are 1 p. Source
T MB \T R */g B T R
positions are as follows (epoch \B1950.0). Orion KL : d \[05¡24@23A ; W51
a \ 05h32m47s.0, a \ 19h21m26s.3,
d \ 14¡24@36A ; G34.3]0.15 : d \ 01¡11@13A ; W49 d \ 09¡01@17A ; Ori 3N : a \
a \ 18h50m46s.2, a \ 19h07m49s.8,
d \[05¡20@50A TMC­1 : d \ 25¡42@45A ; L134N d \[02¡43@31A.
05h32m51s.0, a \ 04h38m16s.6, a \ 15h51m32s.0,

No. 1, 1997 HYDROGENATION OF INTERSTELLAR MOLECULES 309
FIG. 1.õNew detections (dashed line) of at 225 GHz left) and 245 GHz right) for the sources listed in the upper left­hand
CH 2 NH (1 11 õ0 00 (4 14 õ3 13
corner of each panel. The center frequencies are 225554.69 MHz for the three left­hand panels and 245125.98 MHz for the three right­hand panels.
FIG. 2.õRotation diagrams for W51 (top) and Orion KL (bottom). The
rotation temperature, is related to the inverse of the slope of the
T rot ,
best­ït line, and the column density is obtained from the intercept at
(cf. Johansson et al. 1984 ; Blake et al. 1987). Both plots are
E upper \ 0
presented on the same scale to point out the di+erence in the rotation
temperatures between the two sources.
are typical of dark clouds : T B 10 K and n(H 2 ) B 104õ105
cm~3. We did attempt to see whether this transition could
be seen in absorption toward 3C 111, which lies behind a
portion of the Taurus cloud region (cf. Moore, &
Marscher,
Bania but the continuum signal was too weak.
1993),
We shall now discuss the individual sources where detec­
tions were made. Note that in the course of this survey, a
number of other emission lines from a variety of species
were detected, many for the ïrst time astronomically. These
will be discussed in a future paper et al.
(Dickens 1997).
3. RESULTS
3.1. W 51
This region of massive­star formation at a distance of
about 7 kpc has been extensively studied at centimeter,
millimeter, submillimeter, and infrared wavelengths (e.g.,
et al. Avery, & Watson
Bieging 1975 ; Ja+e 1989 ; Bell, 1993 ;
et al. The direction toward W51 includes
Rudolph 1990).
dense molecular cloud material, maser sources, and the two
compact H II regions W51e1 and W51e2 &Downes
(Genzel
Our beam was centered close to W51e1.
1977).
Upper­state energies of the four transitions of CH 2 NH
that we detected range between about 10 and 65 K above
the ground state. The data are well ïtted by a standard

310 DICKENS ET AL. Vol. 479
rotation diagram, consistent with the emission being opti­
cally thin and the energy­level populations being described
by a Boltzmann distribution The line
(Fig. 2). 1 11 õ0 00
appears to be blended with an ethanol line, as is discussed
later in connection with the more complicated spectrum
observed toward Orion KL, but the other transitions are
isolated. The emission is characterized by a between 56
V LSR
and 59 km s~1 and a FWHM line width of about 9 km s~1
We ïnd a total column density of
(Table 2). N(CH 2 NH)\ 8
(^0.4) ] 1013 cm~2 and a rotation temperature T rot \ 22
K. The derived rotation temperature is much lower
^ 3
than the kinetic temperature of the hot gas in the direction
observed ; for example, Harris, & Genzel ïnd a
Ja+e, (1987)
kinetic temperature greater than 70 K from observations of
the CO (J \ 7õ6) transition, and Das, & Genzel
Ho, (1983)
ïnd gas at a temperature B100 K from ammonia data
toward W51e1 and W51e2. However, there is also extended
gas and dust emission characterized by lower temperatures
and densities (e.g., T B 35 K from far­IR dust emission ;
et al. Becklin, & Hildebrand
Rudolph 1990 ; Ja+e, 1984 ;
et al. Since the critical densities for thermali­
Harvey 1986).
zation of the observed transitions are fairly high
CH 2 NH
(B106 cm~3), it would not be surprising if our observed
lines are subthermally excited so that the rotation tem­
perature is rather low. This would seem to imply that the
bulk of our observed emission is from a region more
extended than the hot cores surrounding the compact H II
regions.
In line with our goal of comparing the abundance of
with that of the less hydrogenated species HCN,
CH 2 NH
we may derive an HCN column density from the measure­
ment of H13CN by et al. scaled to our
Goldsmith (1981),
position (about one beamwidth away) using an H13CN
map by Scaling by an isotopic abundance
Howe (1996).
ratio 12C/13C \ 70, and using the same rotation tem­
perature as that measured for we obtain
CH 2 NH,
N(HCN)\ 3.3 ] 1015 cm~2 ; et al.îs
Goldsmith (1981)
measurement for HC15N would give approximately the
same value. Comparing this result with our CH 2 NH
column density, we ïnd an abundance ratio
N(CH 2 NH)/N(HCN)\ 0.025.
3.2. Orion
Toward Orion KL we have detected ïve transitions of
and although with varying signal­
CH 2 NH (Table 2 Fig. 1),
to­noise ratio. Moreover, at the sensitivity reached in this
chemically rich, relatively nearby source, the spectra are
very crowded with emission features. A number of these
features are seen for the ïrst time astronomically, and there
are also some unidentiïed lines (to be presented in Dickens
et al. This line crowding results in several of our
1997).
lines being blended with other molecular emission,
CH 2 NH
thus complicating the rotation diagram analysis. The line
parameters listed in have been obtained from
Table 2
Gaussian ïts to the emission features, subtracting
CH 2 NH
out other blended emission, as will now be discussed
separately for each transition.
The line is blended with two transitions of
1 11 õ0 00
ethanol the stronger of which is also present
(CH 3 CH 2 OH),
in the corresponding W51 spectrum. These three Orion fea­
tures are rather well separated by the Gaussian ïts,
however.
The line of has broad wings which
4 14 õ3 13 CH 2 NH
appear to result from blends with other lines o+set by about
[6 MHz and ]4 MHz. Fitting this triplet produces a
line with a and FWHM that agree with those
CH 2 NH V LSR
of the other transitions of this molecule.
The line shows a broad blueshifted
6 06 õ5 15 CH 2 NH
wing, which, however, can be assigned to SO 2 (v 2 \ 1,
Fortunately there is another vibrationally
10 37 õ10 28 ).
excited line of similar energy
SO 2 (v 2 \ 1, 13 1,13 õ12 0,12 )
above the ground state that is observed in the same
bandpass and that is relatively unblended. We have ïtted
the latter line and subtracted this result from the blended
feature, producing the line parameters given
CH 2 NH­SO 2
in Table 2.
The two remaining lines observed, and
CH 2 NH 2 11 õ2 02
are unblended within the discrimination of our
7 16 õ7 07 ,
signal­to­noise ratio.
The rotation diagram ïtted to the observed Orion KL
lines of yields a rotation temperature
CH 2 NH T rot \ 107
K and a column density
^ 55 N(CH 2 NH)B 6
(^1.8) ] 1014 cm~2. Together with the observed V LSR B 8
km s~1 and line width, *v B 7 km s~1, the data seem most
compatible with emission arising from the
CH 2 NH
```` compact ridge îî source identiïed by et al.
Johansson
and further studied (e.g.) by et al. This is
(1984) Blake (1987).
consistent with the tentative detection of the tran­
3 13 õ2 02
sition of by et al. whose higher
CH 2 NH Sutton (1995),
angular resolution data also indicated emission from the
compact ridge. There does not seem to be identiïable
emission associated with the Orion hot core or
CH 2 NH
plateau source. If the emission does indeed come
CH 2 NH
from the compact ridge rather than the quiescent extended
ridge, it is difficult to obtain a column density ratio relative
to HCN, since the hyperïne structure of the HCN J \ 1õ0
line blends together emission from these two regions within
our beam. If we simply adopt the HCN column density for
the extended ridge material in our line of sight obtained in
the Onsala and the OVRO surveys et al.
(Johansson 1984 ;
et al. we ïnd
Blake 1987), N(CH 2 NH)/N(HCN)B 0.3,
where the HCN value has been obtained from data for the
13C and 15N isotopic lines. The HCN abundance is well
known to be enhanced in the hot, dense material of the hot
core and plateau sources. Using the value of N(HCN) from
the surveys for the hot core would reduce the abundance
ratio of interest to We con­
N(CH 2 NH)/N(HCN)B 0.03.
clude that for the Orion compact ridge,
is in the range 0.03õ0.3. We note that
N(CH 2 NH)/N(HCN)
the present value derived for is an order of
N(CH 2 NH)
magnitude greater than the upper limit obtained by Turner
This could simply be a result of beam dilution, if the
(1991).
bulk of our observed emission is indeed from a relatively
hot, compact source.
We also succeeded in detecting toward the rela­
CH 2 NH
tively cold but dense location 3@ north along the Orion ridge
referred to as the ```` radical­ion peak îî or Ori 3N, where
antenna temperatures peak for low­energy transitions of
many positive ions and radicals & Thaddeus
(Turner 1977).
This region has been studied in the emission of more than
30 transitions from more than 20 molecular species in a
survey at FCRAO (Ungerechts et al. 1995, 1997 ; Bergin
These survey data allow the temperature, density,
1995).
and chemical abundances to be well determined from a
consistent data set. The radical­ion peak has a kinetic tem­
perature of K, a density of cm~3, and
T k B 25 n(H 2 ) B 106
an HCN column density N(HCN)\ 1 ] 1015 cm~2 (Bergin
Since it lacks the embedded luminous sources that
1995).

No. 1, 1997 HYDROGENATION OF INTERSTELLAR MOLECULES 311
are present in our other source regions, it should provide
the most direct comparison to the results of gas­phase ion­
molecule chemical models.
We detected two transitions toward the radical­
CH 2 NH
ion peak but the line has considerably
(Table 2), 1 11 õ0 00
better signal­to­noise ratio. Using this line and an assumed
rotation temperature of 20 K, we ïnd N(CH 2 NH)\ 4
(^1) ] 1012 cm~2, so that N(CH 2 NH)/N(HCN)B 0.004.
Using value for the column density of CO
Berginîs (1995)
and a CO abundance of yields a frac­
CO/H 2 \ 8 ] 10~5
tional abundance relative to of
H 2 f (CH 2 NH)\ 6 ] 10~11.
3.3. Other Sources
We detected only one transition of toward the
CH 2 NH
massive­star formation region G34.3]0.15. An additional
transition, has been recently reported in an
10 1,9 õ10 0,10 ,
unbiased spectral survey of G34.3]0.15 by et
Macdonald
al. who also observed H13CN and HC15N (J \ 4õ3).
(1996),
The ïtted line parameters for our detection agree well with
Macdonaldîs line width and rest velocity for and
CH 2 NH
the two isotopic species of HCN. Using a rotation tem­
perature equal to the excitation temperature assumed for
H13CN by Macdonald K), we obtain
(T ex \ 41
cm~2. When compared to
N(CH 2 NH)\ 1(^0.3) ] 1013
the column density of H13CN, with a ratio of 12C/13C \ 70,
we ïnd N(CH 2 NH)/N(HCN)B 0.01.
We did not detect emission from in the ```` hot
CH 2 NH
core îî source W49. This source is situated at a distance of
about 14 kpc and therefore contains a complex super­
position of many di+erent regions within the observing
beam, including several compact H II regions surrounded
by molecular clouds and the most luminous water maser in
the Galaxy This makes estimates of a molec­
(Nyman 1983).
ular column density very difficult. Assuming the same line
width and excitation temperature that found
Nyman (1983)
for HCN, i.e., B20 km s~1 and 50 K, yields an upper limit
to the column density of less than 5 ] 1013 cm~2,
CH 2 NH
and N(CH 2 NH)/N(HCN)\ 1.
We also failed to detect this molecule in the dark clouds
TMC­1 and L134N. If we assume a line width of B0.5 km
s~1 and an excitation temperature of 5 K, both of which are
typical for millimeter­wavelength transitions of high dipole
moment molecules in dark clouds, we obtain upper­limit
column densities of 1.3 and 1.4 ] 1012 cm~2, respectively.
et al. report N(HCN)B 2 ] 1014 cm~2
Ohishi (1992)
in TMC­1, suggesting that N(CH 2 NH)/N(HCN)\ 0.006.
For L134N, determined that
Swade (1987)
N(H13CN)B 8 ] 1011 near our observing position. Taking
13C/12C \ 70, we obtain N(CH 2 NH)/N(HCN)\ 0.025.
4. DISCUSSION
lists the column densities of and the
Table 3 CH 2 NH
abundance ratio determined as described in
CH 2 NH/HCN,
the previous section. We can compare our results with the
gas­phase chemical models of Bettens, &Herbst
Lee, (1996).
TABLE 3
ABUNDANCE OF CH 2 NH
N(CH 2 NH) T k
Source (]1013 cm~2) N(CH 2 NH)/N(HCN) (K)
Orion KL . . . . . 60 ^ 18 0.03õ0.3 100õ200a
Ori 3N . . . . . . . . 0.4 ^ 0.1 0.004 25b
W51 . . . . . . . . . . 8 ^ 0.4 0.025 35õ100c
G34.3]0.15 . . . . 1 ^ 0.3 0.01 B100d
W49 . . . . . . . . . . \5 \1 B50e
TMC­1 . . . . . . . . \0.13 \0.006 10
L134N . . . . . . . . . \0.14 \0.025 10
Sgr B2f . . . . . . . . 86 0.10 100õ300
NOTEõErrors are 1 p, and upper limits are 3 p. Values of HCN column
density are discussed in the text, and is kinetic temperature.
T k
et al.
a Blake 1987.
b Bergin 1995.
c See text.
et al.
d Macdonald 1996.
e Nyman 1983.
f All data are from et al.
Sutton 1991.
For the range of densities (103õ105 cm~3), temperatures
(10õ50 K), and cloud ages (105 yr to steady state) con­
sidered, et al. ïnd typical values of
Lee (1996)
This agrees well with
N(CH 2 NH)/N(HCN)B 0.002õ0.004.
the value obtained by us for Ori 3N, and their fractional
abundance of with respect to at ```` early îî times
CH 2 NH H 2
seems to agree with our value (6 ] 10~11) as well. This
suggests that gas­phase models adequately account for the
abundance in cold, dense regions with no sign of
CH 2 NH
star formation. This is supported by the upper limit we
report in TMC­1, and our limit for L134N is not inconsis­
tent with this result. However, our ïndings for the GMC
sources suggest that the abundance is enhanced in
CH 2 NH
warmer regions includes representative values of
(Table 3
kinetic temperature, for reference). data in the
T k , CH 2 NH
```` hot core îî source Sgr B2 from et al. give
Sutton (1990)
further evidence of enhancement in warmer sources (Table
et al. note that is a product of
3). Bernstein (1995) CH 2 NH
UV irradiation of model interstellar icy grain mantles.
Therefore, we posit that in these sources, where massive­star
formation is occurring, we observe enhancements of
due to formation on grain mantles and subsequent
CH 2 NH
release back into the gas phase of the ISM. We plan to test
this type of enhanced hydrogenation further in warmer
regions by observations of CH 3 NH 2 .
The National Radio Astronomy Observatory is a facility
of the National Science Foundation, operated under coo­
perative agreement by Associated Universities, Inc. This
research was supported in part by NASA grant NAGW­436
and NSF grant AST 94­20159. We are grateful to the sta+ at
the NRAO 12 m telescope for their assistance, to H. H. Lee
and E. Herbst for helpful communications, and to S. Yama­
moto for his participation in earlier searches for CH 2 NH.
REFERENCES
M. B., Avery, L. W., &Watson, J. K. G. 1993, ApJS, 86,
Bell, 211
E. A. 1995, Ph.D. thesis, Univ.
Bergin, Massachusetts
M. P., Sandford, S. A., Allamandola, L. J., Chang, S., & Schar­
Bernstein,
berg, M. A. 1995, ApJ, 454, 327
J. 1975, in H II Regions and Related Topics, ed. T. L. Wilson &
Bieging,
D. Downes (Berlin : Springer), 443
G. A., Sutton, E. C., Masson, C. R., & Phillips, T. G. 1987, ApJ, 315,
Blake,
621
J. E., Irvine, W. M., DeVries, C. H., & Ohishi, M. 1997, in
Dickens,
preparation
R. E. 1978, Sci. Am., September,
Dickerson, 62
R., &Downes, D. 1977, A&AS, 30,
Genzel, 145
P. D., Brown, R. D., Robinson, B. J., & Sinclair, M. W. 1973,
Godfrey,
Astrophys. Lett., 13, 119
P. F., Langer, W. D., Ellde# r, J., Irvine, W. M., & Kollberg, E.
Goldsmith,
1981, ApJ, 249, 524

312 DICKENS ET AL.
P. M., Joy, M., Lester, D. F., &Wilking, B. A. 1986, ApJ, 300,
Harvey, 737
E., Adams, N. G., & Smith, D. 1983, ApJ, 269,
Herbst, 329
P. T. P., Das, A., &Genzel, R. 1983, ApJ, 266,
Ho, 596
J. 1996, private
Howe, communication
W. M. &Hjalmarson, 1984, Origins Life, 14,
Irvine, A# . 15
W. M., Ohishi, M., &Kaifu, N. 1991, Icarus, 91,
Irvine, 2
D. T., Becklin, E. E., &Hildebrand, R. H. 1984, ApJ, 279,
Ja+e, L51
D. T., Genzel, R., Harris, A. I., Lugten, J. B., Stacey, G. J., & Stutzki,
Ja+e,
J. 1989, ApJ, 344, 265
D. T., Harris, A. I., &Genzel, R. 1987, ApJ, 316,
Ja+e, 231
L. E. B., et al. 1984, A&A, 130,
Johansson, 227
N., Morimoto, M., Nagane, K., Akabane, K., Iguchi, T., & Takagi,
Kaifu,
K. 1974, ApJ, 191, L135
W. H., Johnson, D. R., & Lovas, F. J. 1973, J. Phys. Chem. Ref.
Kirchho+,
Data, 2, 1
H.­H., Bettens, R. P. A., &Herbst, E. 1996, A&AS, 119,
Lee, 111
G. H., Gibb, A. G., Habing, R. J., &Millar, T. J. 1996, A&A, in
Macdonald,
press
A. P., Moore, E. M., & Bania, T. M. 1993, ApJ, 419,
Marscher, L101
Y. C., & Irvine, W. M. 1991, Ap&SS, 175,
Minh, 165
L.­A. 1983, A&A, 120,
Nyman, 307
M., Irvine, W. M., & Kaifu, N. 1992, in Astrochemistry of Cosmic
Ohishi,
Phenomena, ed. P. D. Singh (Dordrecht : Kluwer), 171
P., Irvine, W. M., Schloerb, F. P., Snell, R. L., Bergin, E. A.,
Pratap,
Miralles, M. P., Dickens, J. E., &McGonagle, D. 1994, in ASP Conf. Ser.
65, Clouds, Cores and Low Mass Stars, ed. D. P. Clemens & R. Bar­
vainis (San Francisco : ASP), 25
A., Welch, W. J., Palmer, P., &Dubrulle, B. 1990, ApJ, 363,
Rudolph, 528
E. C., Peng, R., Danchi, W. C., Jaminet, P. A., Sandell, G., &
Sutton,
Russel, P. G. 1995, ApJS, 97, 455
D. A. 1987, Ph.D. thesis, Univ.
Swade, Massachusetts
B. E. 1991, ApJS, 76,
Turner, 617
B. E., & Thaddeus, P. 1977, ApJ, 211,
Turner, 755
H., Bergin, E. A., Goldsmith, P. F., Irvine, W. M., Schloerb,
Ungerechts,
F. P., & Snell, R. L. 1997, ApJ, submitted
1995, in The Physics and Chemistry of Interstellar Molecular
õõõ.
Clouds, ed. G. Winnewisser &G. C. Pelz (Berlin : Springer), 258