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A 2 MILLIMETER SPECTRAL LINE SURVEY OF THE STARBURST GALAXY NGC 253
S. MartI Ò Ð Òn and R. Mauersberger
Instituto de RadioastronomÐÒa Milime Òtrica ( IRAM), Avda. Divina Pastora 7, Local 20, Eí18012 Granada, Spain; martin@iram.es
J. MartI Ò Ð ÒníPintado
Departamento de AstrofisÐÒca Molecular e Infrarroja, Instituto de Estructura de la Materia, CSIC, Serrano 121, Eí28006 Madrid, Spain
C. Henkel
MaxíPlanckíInstitut f u Ør Radioastronomie, Auf dem Hu Øgel 69, Dí53121 Bonn, Germany
and
S. GarcI Ò Ð ÒaíBurillo
Observatorio de Madrid, Alfonso XII, 3, 28014 Madrid, Spain
Received 2005 November 28; accepted 2006 February 8
ABSTRACT
We present the first unbiased molecular line survey toward an extragalactic source, namely the nuclear region of
the starburst galaxy NGC 253. The scan covers the frequency band from 129.1 to 175.2 GHz, i.e., most of the 2 mm
atmospheric window. We identify 111 spectral features as transitions from 25 different molecular species. Eight of
which (three tentatively) are detected for the first time in the extragalactic interstellar medium. Among these newly
detected species, we detected the rare isotopomers 34 SO and HC 18 O + . Tentative detections of two deuterated species,
DNC and N 2 D + , are reported for the first time from a target beyond the Magellanic Clouds. In addition, three hyí
drogen recombination lines are identified, while no organic molecules larger than methanol are detected. Column
densities and rotation temperatures are calculated for all the species, including an upper limit to the ethanol abuní
dance. A comparison of the chemical composition of the nuclear environment of NGC 253 with those of selected
nearby galaxies demonstrates the chemical resemblance of IC 342 and NGC 4945 to that of NGC 253. On the other
hand, the chemistries characterizing NGC 253 and M82 are clearly different. We also present a comparison of the
chemical composition of NGC 253 with those observed in Galactic prototypical sources. The chemistry of NGC 253
shows a striking similarity with the chemistry observed toward the Galactic center molecular clouds, which are
thought to be dominated by lowívelocity shocks. This resemblance strongly suggests that the heating in the nuclear
environment of NGC 253 is dominated by the same mechanism as that in the central region of the Milky Way.
Subject headings: galaxies: abundances --- galaxies: individual (NGC 253) --- galaxies: ISM --- galaxies: nuclei ---
galaxies: starburst --- radio lines: galaxies --- surveys
1. INTRODUCTION
Our knowledge of the chemical composition of the interstellar
medium ( ISM) in the nuclei of external galaxies has so far been
restricted to only #30 molecular species (see x 4.1), which is still
far from the 129 molecules detected in the interstellar and cirí
cumstellar medium within our Galaxy ( Lovas 2004). The initial
detection of most of these 129 molecules has been the result of
unbiased frequency scans toward specific Galactic sources, such
as massive staríforming regions (e.g., Sgr B2(M), Cummins et al.
1986; OriMCí1, Lee et al. 2001), cold molecular clouds (e.g.,
TMCí1, Kaifu et al. 2004), and evolved stars (e.g., IRC +10216,
Cernicharo et al. 2000). Spectral line scans toward Galactic sources
focused mainly on the hot cores associated with massive starí
forming regions such as Sgr B2 and OrioníKL due to their complex
chemistry as well as the brightness of their lines. Extragalactic
molecular spectroscopy has been so far limited to selecting the
strongest features seen in the Galactic center and disk sources
and observing them toward well selected extragalactic targets.
The most ambitious census of the molecular content of an external
galaxy is the compilation of data toward the starburst galaxy
NGC 4945 by Wang et al. (2004). Their multiline study comí
bines observations of 19 molecular species previously known in
other extragalactic sources.
Up to now unbiased molecular line surveys of extragalactic
sources do not exist in spite of providing a powerful tool for deí
termining the physical parameters of the molecular ISM. Obserí
vations of a large number of molecular lines with similar angular
resolution allow us to describe the chemical complexity of the
source, which provides fundamental information on the physical
processes heating the medium in heavily obscured regions. In
addition, frequency scans lead to serendipitous detections of new
species.
The nuclear starburst of the galaxy NGC 253, the target of this
survey, is one of the brightest and most prolific (in terms of moí
lecular line detections) extragalactic molecular line sources.
The Sculptor galaxy NGC 253 is an almost edgeíon barred spiral
(i ® 72 # 78 # , Puche et al. 1991; Pence 1981) classified as type
SAB(s)c (de Vaucouleurs et al. 1991) or Sc(s) (Sandage &Tammann
1987). At a distance of #3 Mpc (e.g., Mouhcine et al. 2005),
NGC 253 is one of the nearest archetypes of nuclear starburst
galaxies (Rieke et al. 1980). Its nuclear region contains one of the
brightest extragalactic IRAS sources with a 100 #m flux of 1860 Jy
(Soifer et al. 1989). In fact, most of the overall IR luminosity of this
galaxy (L IR ® 2:1 ; 10 10
L# , SFR IR ® 3:6 M # yr #1 , Strickland
et al. 2004) stems from the regions of intense massive star forí
mation within its central few hundred parsecs. Violent massive
star formation is also revealed by the high supernova rate of
0.05--0.3 yr #1 in the nuclear region of this galaxy ( Mattila &
Meikle 2001; Ulvestad & Antonucci 1997).
The high nuclear star formation activity is driven by huge
amounts of molecular gas in the central few hundred parsecs of
450
The Astrophysical Journal Supplement Series, 164:450--476, 2006 June
# 2006. The American Astronomical Society. All rights reserved. Printed in U.S.A.

NGC253 [(1:3 2:6) ; 10 9 M # ; Canzian et al. 1988; Mauersberger
et al. 1996], likely powered by the bar structure first detected at
nearíinfrared wavelengths (Scoville et al. 1985; Forbes &Depoy
1992). Positionívelocity diagrams from highíresolution obserí
vations of the nuclear region of NGC 253 show that the molecular
material seems to follow orbital motions around the dynamií
cal center, which are interpreted as x 1 and x 2 orbits within the
context of a barlike potential (OH, Turner, 1985; CO, Canzian
et al. 1988; Das et al. 2001; Paglione et al. 2004; HCN, Paglione
et al. 1995; CS, Peng et al. 1996; SiO, H 13 CO + , GarcÐÒaíBurillo et al.
2000).
Here we present the first unbiased molecular line survey
carried out toward a source outside the Milky Way. Technically,
this survey was driven by the availability of two new wideíband
(1 GHz) receivers at 2 mm and new 1 GHzwide filter bank specí
trometers, together with the large collecting area of the IRAM
30 m telescope. Species such as NO, NS, SO 2 , H 2 S and H 2 CS,
detected for the first time toward a starburst environment, were
previously reported by MartÐÒn et al. (2003, 2005). Additional
new molecules together with a full analysis of the data are
presented.
2. OBSERVATIONS AND RESULTS
The first molecular frequency scan of an external galaxy was
carried out at 2 mmwith the 30 m IRAM telescope at Pico Veleta,
Spain, between 2001 and 2004. It covers #86% of the observable
2 mm atmospheric window, from 129.1 to 175.2 GHz (Fig. 1). At
these frequencies, the telescope beamwidth ranges from 19 00 to
14 00 . Figure 2 shows the size of the 30 m beam on top of an
interferometric map of SiO emission (GarcÐÒaíBurillo et al. 2000).
A Kíband image of NGC 253 (Engelbracht et al. 1998) is disí
played in gray scale to illustrate the size of the observed region
relative to the galaxy.
The observations were pointed at the position # J2000 ®
00 h 47 m 33: s 3, # J2000 ® #25 # 17 23 00 . This position is #6 00 southí
east of the dynamical center of NGC 253. A pointing accuracy of
#3 00 was achieved by measuring cross scans on nearby continí
uum sources every #2 hr. The observations were carried out in a
wobbleríswitching mode with a symmetrical beam throw of 4 0 in
azimuth and a switching frequency of 0.5 Hz.
Each of the two available SIS receivers working at 2 mm
with orthogonal polarizations provides a 1 GHz bandwidth. The
Fig. 1.--- Complete 2 mm spectral frequency scan toward the nuclear region of NGC 253.
Fig. 2.--- Grayíscale, Kíband image of NGC 253 (left panel; Engelbracht et al. 1998) and highíresolution SiO emission (right panel; GarcÐÒaíBurillo et al. 2000). Beam sizes
of the IRAM 30 m telescope for the extreme frequencies observed in the survey are shown as circles on top of the SiOmap. The continuous line shows the 19 00 beam at 129GHz
and the dashed line the 14 00 beam at 175GHz.A cross indicates the nominal position of the 2 mm scan in this work and a dot indicates the dynamical center as derived byGarcÐÒaí
Burillo et al. (2000).
2 mm LINE SURVEY OF NGC 253 451

receivers were tuned to adjacent frequencies, with an overlap
of 100 MHz, to cover an instantaneous bandwidth of 1.9 GHz
for each frequency setup. As spectrometers we used two 256 ;
4 MHz filter banks, providing a velocity resolution between
7 and 9 km s #1 at the observed frequencies.
The rejection of the image band (the upper sideband) was typí
ically #10 dB so that only the very few strongest lines were also
detected from the image band. The image sideband rejection was
calculated for each frequency setup ( Fig. 3, lower panel ) by
measuring the difference in power between a hot and a cold load,
first with full rejection of the upper sideband and then with full
rejection of the lower sideband by using a MartiníPuplett interí
ferometer. The image gain is then computed as the ratio between
these two measurements. We assume this gain ratio to be coní
stant throughout the whole 1 GHz frequency band covered by
each back end.
The spectra were calibrated with a standard dual load system.
The temperature scale of the spectra is in TMB obtained as
TMB ® (F eA =B eA )T #
A ; Å1�
where the forward efficiency (F eA ) is 0.93 at 2 mm. The beam efí
ficiency was calculated for each frequency with the Ruze function,
B eA ® 1:2# exp ­#(4#R#=k) 2
#; Å2�
using # ® 0:69, R# ® 0:07, and k in mm.
At the latitude of Pico Veleta Observatory, NGC 253 is at eleí
vations >20 # only for about 4 hr. A complete observing session
with about 120 minutes effective observing time was spent on each
frequency setup, out of which half the time was spent on source.
The integration time for each frequency setup is shown in the cení
tral panel of Figure 3.
The observing mode combined with the stability of the system
provided highíquality baselines. Only linear baselines were subí
tracted from the spectra. The rms of the residuals after subtractí
ing the fitted line profiles on each frequency setup is shown in the
upper panel of Figure 3. The values are mostly on the order of
2--4 mK in #8 km s #1 wide channels, increasing up to #6--
10 mK at the highest measured frequencies.
In the scanned 46 GHz wide band we have detected a total of
111 lines, the weakest of which have an intensity of #3 mK. The
spectral density of detected lines is 2.4 features GHz #1 . Since the
Doppler line width in NGC 253 is about 200 km s #1 , or 100 MHz,
we are not limited by line blending, but by a lack of sensitivity. For
targeted searches with longer integration time it may still be posí
sible to detect and identify lines at levels of #1 mK .
Figure 4 shows the full 2 mm spectral line scan of NGC 253
with the line identifications superposed on the spectra. In Table 1
we are listing the parameters derived from Gaussian fits to the
detected lines. We identify 25 molecular species and 3 hydrogen
recombination lines.
3. DATA ANALYSIS
3.1. Line Identification and Fitting
In order to identify and fit the observed features we have used
the rest frequencies from the molecular line catalogs of Lovas
(1992, 2004) and Pickett et al. (1998). The existing molecuí
lar line surveys toward the Galactic center sources Sgr B2(OH)
(Cummins et al. 1986; Turner 1989, 1991) and Sgr B2(N, M,
and NW) (Nummelin et al. 1998, 2000) were also used as addií
tional information to estimate the expected intensity of the fainter
identified lines not included in the Lovas (2004) catalog.
One of the main difficulties in identifying lines from singleí
dish observations of extragalactic sources stems from their large
width, typically #100 km s #1 . In addition, the emission lines
from the nuclear region of NGC 253 show two velocity comí
ponents, roughly at 180 and 280 km s #1 , with line widths of
#100 and 110 km s #1 , respectively (from CS; MartÐÒn et al.
2005). These components arise from the two main molecular
lobes separated by 10 00 and located on both sides of the nucleus.
As a consequence, some of the observed transitions will apí
pear partially or totally blended. Table 1 shows the results of the
Gaussian fits to all the molecular transitions identified toward
the NGC 253 nuclear region. In the third column of Table 1 we
have included a note to indicate whether the transition is affected
by blending. Appendix A describes the fitting procedures used for
the different types of blending found in our spectral line survey.
3.2. Column Density Determination
From the observed main beam brightness temperature of the
measured molecular lines one can estimate column densities and
rotation temperatures (T rot ) for each species (see x B1 for a deí
tailed discussion). All the necessary spectroscopic information
required to derive these parameters (i.e., A ul , #, g u , E u , and Z in
eq. [B6]) were extracted or derived from the JPL catalog ( Pickett
et al. 1998).
In order to derive T rot and to extrapolate the column densities
in the observed states to a total column density for a given molí
ecule, more than one transition has to be measured. In the case
Fig. 3.---Summary of the observational parameters of individual frequency
setups. Open and filled symbols refer to each of the 2 mm SIS receivers. Upper
panel: rms of the spectrum in 7--9 km s #1 wide channels in Tmb after subtracting
a linear baseline; Middle panel: Total integration time, which includes the time
spent on the reference position; Lower panel: Measured rejection of the image
(upper) sideband.
MARTI Ò Ð
ÒN ET AL.
452

Fig. 4.---2 mm spectral line survey composite toward the nuclear region of NGC 253. The resolution has been smoothed to four channels, which is equivalent to
velocity resolutions of 35--27 km s #1 . Line identifications are plotted on top of the spectra. Species labeled with i ® molecule correspond to lines observed from the
image band.

Fig. 4.--- Continued
454

TABLE 1
Results from Gaussian Fits to the Observed Lines
Molecule
Transition
#
( MHz) Note
R TMB dv
(mK km s #1 )
V LSR
( km s #1 )
#v 1 = 2
( km s #1 )
TMB
(mK)
SiO 3--2 v = 0.......................................... 130268.6 1580 (50) 183 76 18
1170 (50) 260 91 12
C 2 S 11 10 --10 9 ........................................... 131551.9 640 (90) 230 151 4.0
HNCO 6 0,6 --5 0,5 ....................................... 131885.7 2400 (150) 182 72 32
2000 (180) 284 80 24
CH 2 NH 2 1,1 --1 1,0 ...................................... 133272.1 hf 420 (40) 187 83 b 4.8
450 (40) 283 83 b 5.1
OCS 11--10 .............................................. 133785.9 1220 (130) 212 200 5.8
SO 2 8 2,6 --8 1,7 ............................................ 134004.8 300 (70) 242 91 3.1
H36# ........................................................ 135286.0 b, s 2800 (. . .) 250 a 282 a 9.3
H 2 CS 4 1,4 --3 1,3 ......................................... 135297.8 b 1040 (90) 221 158 6.2
SO 2 5 1,5 --4 0,4 ............................................ 135696.0 740 (160) 245 a 140 a 4.2
34 SO 4 3 --3 2 ............................................... 135775.3 b 300 (60) 180 a 73 a 3.8
b 200 (70) 274 a 88 a 2.1
HC 3 N 15--14............................................ 136464.4 2410 (130) 183 77 29
2000 (300) 271 85 22
CH 3 CCH 8 k --7 k ........................................ 136728.0 3410 (110) 255 158 20
SO 4 3 --3 2 .................................................. 138178.5 1400 (200) 180 73 18
1100 (200) 274 88 12
13 CS 3--2 .................................................. 138739.3 550 (70) 188 70 b 7.4
390 (50) 274 70 b 5.2
NH 2 CN 7 1,7 --6 1,6 ...................................... 139032.0 580 (110) 280 a 155 a 3.5
H 2 CS 4 1,3 --3 1,2 ......................................... 139483.4 b 1340 (150) 221 176 7.2
NH 2 CN 7 0,7 --6 0,6 ...................................... 139842.1 490 (160) 280 a 155 a 3.0
NH 2 CN 7 0,7 --6 0,6 ...................................... 139954.4 m 520 (160) 279 140 3.5
C 2 S 10 11 --9 10 ............................................ 140180.7 470 (130) 250 85 5.2
SO 2 6 2,4 --6 1,5 ............................................ 140306.1 510 (60) 248 115 4.1
H 2 CO 2 1,2 --1 1,1 ......................................... 140839.5 4250 (140) 186 104 b 38
6050 (150) 288 104 b 55
C 2 S 11 11 --10 10 .......................................... 142501.7 430 (70) 192 85 a 4.7
CH 3 OH 3 1,3 --2 1,2 A+ ............................... 143865.7 1000 (130) 225 168 5.6
C 2 S 12 11 --11 10 .......................................... 144244.8 290 (120) 210 112 2.4
C 34 S 3--2 .................................................. 144617.1 3000 (300) 188 105 27
1500 (300) 284 86 16
cíC 3 H 2 3 1,2 --2 2,1 ....................................... 145089.5 b, s 1750 (. . .) 234 a 156 a 11
CH 3 OH 3--2 ............................................. 145103.2 b, m 9400 (200) 185 115 77
5760 (190) 294 96 57
HC 3 N 16--15............................................ 145560.9 b, s 2100 (. . .) 183 a 77 a 26
b, s 1700 (. . .) 271 a 85 a 19
H 2 CO 2 0,2 --1 0,1 ......................................... 145602.9 b 3400 (100) 177 90 b 35
b 3300 (100) 283 90 b 34
OCS 12--11 .............................................. 145946.8 580 (170) 212 a 210 a 2.6
CH 3 OH 3 1,2 --2 1,2 A-- ............................... 146368.3 1020 (150) 196 139 6.9
SO 2 4 2,2 --4 1,3 ............................................ 146605.5 800 (170) 241 148 5.1
CS 3--2 ..................................................... 146969.0 b 11900 (200) 185 100 111
b 13700 (200) 288 117 110
H35# ........................................................ 147046.8 b, s 2800 (. . .) 250 a 282 a 9.3
CH 3 CN 8 k --7 k ........................................... 147174.5 b, m 2040 (180) 258 165 12
HOCO + 7 0,7 --6 0,6 ...................................... 149675.8 410 (110) 281 68 5.6
NO 3
2
1
2 # ” ............................................ 150176.5 hf 3400 (500) 244 158 10
cíC 3 H 2 2 2,0 --1 1,1 ....................................... 150436.5 b, s 600 (. . .) 234 a 156 a 3.6
H 2 CO 2 1,1 --1 1,0 ......................................... 150498.3 b 5300 (100) 180 99 b 50
b 5100 (100) 286 99 b 48
NO 3
2
1
2 # # ............................................ 150546.5 hf, b, s 3400 (. . .) 244 a 158 a 10
cíC 3 H 2 4 0,4 --3 1,3 ....................................... 150820.6 b 1150 (. . .) 234 b 156 b 7.0
cíC 3 H 2 4 1,4 --3 0,3 ....................................... 150851.9 b 3380 (110) 234 b 156 b 20
cíC 3 H 2 5 1,4 --5 0,5 ....................................... 151343.8 b, s 165 (. . .) 234 a 156 a 1.0
SO 2 2 2,0 --2 1,1 ............................................ 151378.6 b 590 (170) 225 158 3.5
DNC 2--1 ................................................. 152609.7 490 (160) 225 61 7.5
CH 3 CCH 9 k --8 k ........................................ 153817.2 b 4300 (200) 271 157 26
HNCO 7 0,7 --6 0,6 ....................................... 153865.0 2660 (150) 183 72 34
b, s 2240 (. . .) 284 a 80 a 26
N 2 D + 2--1 ................................................. 154217.0 570 (170) 219 120 4.5
HNCO 7 1,6 --6 1,5 ....................................... 154414.7 540 (110) 283 a 80 a 6.4
HC 3 N 17--16............................................ 154657.3 1640 (120) 191 72 b 21
1400 (110) 272 72 b 18
455

that only one transition was detected in the 2 mm scan, an educated
guess of the rotation temperature was made (see Appendix C).
Table 2 shows the derived total column densities and rotation
temperatures for the detected species in NGC 253. Sourceíaveraged
abundances are also given, where a sourceíaveraged molecular
hydrogen column density N (H 2 ) ® 6:7 ; 10 22 cm #2 has been así
sumed (see x 4.4.1).
3.3. Uncertainty in Column Densities
and Rotation Temperatures
The main source of uncertainty, when deriving the column
densities and rotation temperatures, arises from the uncertainty
of the extent of the emitting regions. The sourceíaveraged brightí
ness temperature (T B ) can be estimated from the measured main
beam brightness temperature as
T B ® # #1 bf TMB ; Å3�
where, in the approximation of a Gaussian source distribution
of size # s observed with a Gaussian beam of size # b , the beam
filling factor # bf ® # 2
s /(# 2
s ” # 2
b ) accounts for the dilution effect
due to the coupling between the source and the telescope beam.
In order to correct for beam dilution in the nuclear region of
NGC 253, Mauersberger et al. (2003) smoothed the highíresolution
interferometric CS map from Peng et al. (1996) to convert the
three observed transitions of CS by Mauersberger &Henkel (1989)
TABLE 1--- Continued
Molecule
Transition
#
( MHz) Note
R TMB dv
(mK km s #1 )
V LSR
( km s #1 )
#v 1=2
( km s #1 )
TMB
(mK)
cíC 3 H 2 3 2,2 --2 1,1 ............................................... 155518.2 820 (180) 205 118 6.6
CH 3 OH 6 0,6 --6 #1,6 E........................................ 157048.6 540 (120) 275 103 b 5.0
CH 3 OH 5 0,5 --5 #1,5 E........................................ 157178.9 b 960 (150) 242 103 b 8.8
CH 3 OH J 0;k J #1;k E ....................................... 157270.7 b, m 2110 (160) 183 103 b 19
J = 1. . .4; k = J ............................................... 5700 (200) 290 103 b 52
SO 3 4 --2 3 .......................................................... 158971.7 700 (200) 180 a 92 7.2
b 1000 (200) 274 a 84 11
NH 2 CN 8 0,8 --7 0,7 .............................................. 159814.6 750 (180) 283 b 155 b 4.5
NH 2 CN 8 0,8 --7 0,7 .............................................. 159942.7 m 700 (200) 283 b 158 b 4.1
H34# ................................................................ 160211.5 2800 (200) 250 282 9.3
NH 2 CN 8 1,7 --7 1,6 .............................................. 161000.3 b 1430 (170) 280 155 a 8.6
NS 7
2
5
2 e ........................................................ 161297.2 hf 1800 (500) 199 224 3.1
NS 7
2
5
2
f ........................................................ 161697.2 hf 2400 (400) 249 278 3.4
HC 3 N 18--17.................................................... 163753.4 1500 (400) 185 a 73 20
700 (300) 272 a 98 6.6
CH 3 OH J 1;k#1 J 0;k E...................................... 165050.1 b, m 1700 (400) 169 72 22
J = 1. . .3; k = J ............................................... 4700 (400) 273 100 45
SO 2 5 2,4 --5 1,5 .................................................... 165144.6 b, s 640 (. . .) 248 a 115 a 5.2
CH 3 OH 4 1,3 --4 0,4 E .......................................... 165190.5 b 900 (200) 291 90 a 9.1
SO 2 7 1,7 --6 0,6 .................................................... 165225.4 b, s 900 (. . .) 248 a 115 a 7.4
CH 3 CN 9 k --8 k .................................................. 165568.9 b, m 1200 (300) 258 a 166 a 7.0
H 2 S 1 1,0 --1 0,1 .................................................... 168762.7 1520 (180) 180 100 14
2250 (80) 275 88 24
H 2 CS 5 1,5 --4 1,4 ................................................. 169113.5 970 (200) 270 165 a 5.5
CH 3 OH 3 2,3 --2 1,3 .............................................. 170060.6 2400 (400) 248 178 12
HC 18 O + 2--1..................................................... 170322.7 500 (300) 206 46 11
500 (200) 285 49 10
CH 3 CCH 10 k --9 k .............................................. 170905.6 5100 (500) 275 102 47
HOCO + 8 0,8 --7 0,7 .............................................. 171055.9 b 1700 (300) 295 66 25
SO 4 4 --3 3 .......................................................... 172181.4 1300 (300) 180 a 87 14
900 (300) 274 a 74 11
H 13 CN 2 k --1 k ................................................... 172677.9 hf 3200 (200) 181 74 40
2400 (200) 279 66 34
HC 3 N 19--18.................................................... 172849.3 2800 (400) 184 74 35
1800 (400) 269 75 22
H 13 CO + 2--1..................................................... 173506.7 860 (90) 170 60 b 14
1450 (100) 269 60 b 23
SiO 4--3 v = 0.................................................. 173688.3 1070 (110) 184 a 76 a 13
1000 (120) 270 a 91 a 10
cíC 3 H 3 1,2 --2 1,1 ................................................ 174086.1 hf 900 (200) 180 a 41 b 6.5
1600 (500) 280 a 40 12
HN 13 C 2--1 ...................................................... 174179.4 1050 (170) 271 60 16
C 2 H 2--1........................................................... 174663.2 hf 30000 (6000) 162 108 82
27000 (6000) 282 104 88
Notes.---See Appendix A for details: (b) blended line; (s) synthetic Gaussian to isolate a given feature (see text for details on the parameters
used for deriving these profiles for each species); (hf ) hyperfine structure (frequency and intensity refers to the main component of the group);
(m) multitransition line (frequency refers to the main component of the group).
a Parameter fixed in the Gaussian fit.
b Parameters forced to have the same value in the Gaussian fit.
MARTI Ò Ð
ÒN ET AL.
456 Vol. 164

to a common angular resolution of 32 00 . Since we know the intení
sities of the CS lines with the original resolution (T MB ) and those
corrected to a 32 00 beam (T 32
MB ) we can derive the beam filling
factor. From the expression # 32
bf TMB ® # bf T 32 0
MB
we estimate the
extent of the emitting region to be # s ® 23 00 , 21 00 , and 19 00 for the
CS J ® 2 1, J ® 3 2, and J ® 5 4 transitions, respectively. In
our analysis we will then consider an equivalent source size of
20 00 to convert our main beam brightness temperature to sourceí
averaged brightness temperature.
The other main source of uncertainty stems from the assumpí
tion that all the observed species arise from the same volume.
The only way to confirm this would be highíresolution observaí
tions of different transitions of several species. The interferoí
metric map of the J ® 2 1 line of CS by Peng et al. (1996)
shows that most of the emission of this molecule is concentrated
within a radius of #20 00 . The CS emission is seen to be clumped
into four large main molecular cloud complexes symmetrically
distributed with respect to the dynamical center and roughly
aligned with the molecular bar. A recent highíresolution mulí
tiline study carried out toward IC 342 ( Meier & Turner 2005)
shows clear differences in the distribution of several important
molecular tracers. Toward NGC 253, however, highíresolution
maps of molecular transitions requiring different excitation coní
ditions such as HCN, CO, SiO, H 13 CO + , and NH 3 ( Paglione
et al. 1995, 2004; GarcÐÒaíBurillo et al. 2000; Ott et al. 2005)
show very similar distributions, in good agreement with that of
CS. Therefore, it is plausible that for NGC 253, as long as the bulk
of the emission arises from the area confined by the starburst,
the extent of the emission of different molecules is similar.
We now address the question of how an error in the assumed
source size will affect our determination of physical parameters.
As derived from equation ( B1) (see x B1), the calculated total
column density is affected by the filling factor through T B . The
upper panel in Figure 5 shows the variation of the inverse of the
filling factor (i.e., the conversion factor from TMB to T B from
eq. [3]) as a function of source size, normalized to the value for a
source of 20 00 . This value represents, in a first approximation, the
factor by which we should multiply our derived column densities
if the source extent would be different from that assumed. We
have considered the extreme cases of the beam sizes at 129 and
175 GHz. The gray shaded area corresponds to source sizes
between 10 00 and 30 00 , where we expect the source size to be
TABLE 2
Physical Parameters Derived for All Detected Species
Molecule
Velocity
( km s #1 )
N a
(cm #2 )
T rot
( K)
[X ]/[H 2 ] b
(;10 #9 )
SiO ...................... 180 5.0(1.0) ; 10 12 7.4(0.7) 0.07
270 3.6(0.9) ; 10 12 8.7(1.2) 0.05
C 2 S ...................... 1.4(1.2) ; 10 13 24(9) 0.2
NH 2 CN................ 1.2(0.5) ; 10 13 67(13) 0.2
CH 3 CN ................ 2.0(0.6) ; 10 13 9.6(0.7) 0.3
cíC 3 H .................. 130 8.1(1.6) ; 10 12 12 0.1
280 1.5(0.3) ; 10 13 12 0.2
cíC 3 H 2 ................. 3.0(6.0) ; 10 13 9(8) 0.4
NS c ...................... 4.3(0.6) ; 10 13 7.2(1.0) 0.6
HC 3 N................... 180 1.8(0.5) ; 10 13 33(4) 0.3
270 1.9(0.6) ; 10 13 24(3) 0.3
HOCO + ................ 2.6(1.8) ; 10 13 12 0.4
H 2 CS ................... 4.4(4.5) ; 10 13 11(4) 0.6
SO 2 ...................... 5.0(2.0) ; 10 13 15(2) 0.7
H 2 S ...................... 180 2.3(1.1) ; 10 13 12 0.3
275 3.5(1.4) ; 10 13 12 0.5
CH 2 NH................ 190 2.5(1.1) ; 10 13 12 0.4
300 2.8(0.9) ; 10 13 12 0.4
HNC c .................. 270 7.2(1.2) ; 10 13 12 1
SO ....................... 180 4.5(3.3) ; 10 13 40(24) 0.7
270 3.1(2.8) ; 10 13 34(24) 0.5
HNCO ................. 180 5.7(2.7) ; 10 13 23(6) 0.8
280 4.8(2.5) ; 10 13 23 0.7
HCO +c ................. 170 3.9(0.8) ; 10 13 12 0.6
270 6.4(1.2) ; 10 13 12 1
H 2 CO................... 180 6.9(0.9) ; 10 13 27(2) 1
285 9.0(1.7) ; 10 13 34(4) 1
OCS d ................... 2.5(0.3) ; 10 14 17(2) 4
HCN c .................. 180 1.8(0.3) ; 10 14 12 2
280 1.4(0.2) ; 10 14 12 2
CS c,d .................... 180 2.0(0.2) ; 10 14 9.7(0.4) 3
280 1.4(0.2) ; 10 14 10.0(0.2) 2
CH 3 CCH ............. 4.3(0.2) ; 10 14 63(20) 6
CH 3 OH................ 8.3(0.3) ; 10 14 11.6(0.2) 12
C 2 H...................... 160 6.8(1.1) ; 10 14 12 10
280 5.6(0.9) ; 10 14 12 8
NO d ..................... 4.0(0.4) ; 10 15 6(2) 60
Note.---Sourceíaveraged column densities, rotation temperature, and abuní
dances relative to H 2
a Errors derived from the rotation diagrams without beam filling factor coní
siderations (see x 3.3).
b Assumed N ( H 2 ) ® 6:7 ; 10 22 cm #2 . See Table 6.
c Calculated from the observed 13 C isotope transition with 12 C/ 13 C # 40
from Henkel et al. (1993).
d Using additional transitions measured by MartÐÒn et al. (2003, 2005).
Fig. 5.---Upper panel: Variation of the filling factor normalized to that
corresponding to a source extent of 20 00 as a function of source size. The solid
line corresponds to the beam at 129 GHz and the dashed line to the beam at 175
GHz. The shaded area highlights values with source sizes between 10 and 30 00 .
Lower panel: Effect of the assumed source size on the rotation temperature
determined from rotation diagrams relative to the assumed T rot ® 15 K for
# s ® 20 00 with two transitions measured at 129 and 175 GHz, where the involved
levels are separated by 60 (solid line), 30 (dashed line) and 15 K (dotted line).
2 mm LINE SURVEY OF NGC 253 457
No. 2, 2006

confined. In the case of a source size larger than 20 00 , column
densities would be overestimated by only #20%, while for sizes
down to 10 00 , the column density would be underestimated by
less than a factor of 2. Only in the case of an extremely clumped
and compact emission where source sizes were well below 10 00 ,
column densities would be dramatically affected. We have also
considered the effect of source size on the derived rotation
temperature from rotation diagrams. We will assume a case in
which we derive T rot ® 15 K (for a 20 00 source) by observing two
transitions of a given species at 129 and 175 GHz, respectively.
In the lower panel in Figure 5 we have plotted the relative change
in the derived T rot for those cases in which the transitions would
have upper levels with an energy difference of 60, 30, and 15 K.
It is obvious that the more affected T rot is that derived from the
data with the smaller dynamical range in energies. From Figure 5
we see that the largest expected uncertainty in the determination
of T rot is smaller than 20%.
4. DISCUSSION
4.1. Extragalactic Molecular Census
Table 3 lists the detection of all the species known to date in
the extragalactic ISM in chronological order. During the last
three decades, 29 molecules and 11 rare isotopic substitutions
have been detected outside the Galaxy. As a result of the freí
quency scan presented in this paper (see also MartÐÒn et al. 2003,
2005), 8 new molecules and 2 isotopes (denoted by asterisks in
Table 3) have been added to the census of known extragalactic
species, which represents a #30% increase in the total number of
identified molecules. While H 2 S and H 2 CSwere detected toward
the nearby Large Magellanic Cloud ( HeikkilaØ et al. 1999), their
identification in our survey (see MartÐÒn et al. 2005) represents the
first detections toward a starburst galaxy. Note that CH 2 NH,
HOCO + , and C 3 H ( Table 3) are tentatively detected and need
confirmation by observations of additional transitions. In Table 3
the known extragalactic species not detected in the 2 mm band
are indicated by dagger symbols.
4.2. Deuterated Molecules in NGC 253
The observation of deuterated molecules in external galaxies
is interesting for several reasons. The interstellar D/H ratio deí
pends on the degree of processing of the gas, since Dwas mainly
produced soon after the big bang, while nuclear processes during
stellar evolution rather destroy than form deuterium. A low D/H
ratio may therefore indicate that the observed gas component has
been recycled in stars, and a higher ratio could be a hint for recent
infall of poorly processed material. It has been observed that D is
strongly enriched ( by a factor of >10 4 compared to the cosmic
D/H ratio) in several molecules observable at centimeter and
millimeter wavelengths. From model calculations, such a high
degree of fractionation is expected if the molecules are being
formed in a relatively cool (T kin < 30 K) environment (e.g.,
TABLE 3
Known Extragalactic Molecules (See x 4.1)
Molecule
Publication
Year Telescope Reference
OHy ............................. 1971 OVRO 1
H 2 CO........................... 1974 64 m 2
COy ............................. 1975 11 m 3
13 COy .......................... 1975 11 m 4
H 2 Oy............................ 1977 100 m 5
HCNy .......................... 1977 11 m 6
H 2 y .............................. 1978 2.3 m 7
NH 3 y............................ 1979 100 m 8
HCO +
y......................... 1979 7 m 9
CHy ............................. 1980 64 m 10
CS................................ 1985 7 m 11
C 3 H 2 ............................ 1986 43 m 12
CH +
y............................ 1987 1.4 m 13
CH 3 OH........................ 1987 30 m 14
CNy ............................. 1988 30 m 15
C 2 H.............................. 1988 30 m 15
HNCy .......................... 1988 30 m 15
HC 3 N........................... 1988 30 m 15, 16
HNCO ......................... 1989 30 m 17, 18
C 34 S............................. 1989 30 m 19
C 18 Oy .......................... 1991 11 m / 30 m 20
C 17 Oy .......................... 1991 11 m / 30 m 20
SO ............................... 1991 15 m 21
N 2 H +
y .......................... 1991 30 m 22
SiO .............................. 1991 30 m 22
H 13 CO + ........................ 1991 30 m 22
HN 13 C ......................... 1991 30 m 22
H 13 CNy ....................... 1991 30 m 22
CH 3 CCH ..................... 1991 30 m 23
CH 3 CN ........................ 1991 30 m 23
13 CS............................. 1993 30 m 24
OCS............................. 1995 30 m 25
HCOy .......................... 1995 12 m 26
DCO + .......................... 1996 15 m 27
DCN ............................ 1996 15 m 27
HC 15 Ny ....................... 1999 15 m 28
H 2 S .............................. 1999 15 m 29
H 2 CS ........................... 1999 15 m / 30 m 29, 30
CO +
y............................ 2000 30 m 31
SO 2 # ............................ 2003 30 m 32
NO # ............................. 2003 30 m 32
NS # .............................. 2003 30 m 32
34 SO # ........................... 2003 30 m 32
HOC +
y......................... 2004 30 m 33
C 2 S # ............................. 2006 30 m 34
TABLE 3--- Continued
Molecule
Publication
Year Telescope Reference
CH 2 NH # ......................... 2006 30 m 34
NH 2 CN # ......................... 2006 30 m 34
HOCO +# ......................... 2006 30 m 34
C 3 H # ............................... 2006 30 m 34
HC 18 O +# ......................... 2006 30 m 34
Notes.---Table adapted and updated from Mauersberger & Henkel (1993).
Asterisks ( # ) denote molecular species not previously identified outside the Milky
Way. Dagger symbols (y) denote species not measured in this survey. Telescopes:
(OVRO) Owens Valley interferometer; (64 m) Parkes 64 m; (11 m) NRAO Kitt
Peak 11 m; (100 m) Effelsberg 100 m; (2.3 m) Steward observatory 2.3 m; (7 m)
AT&TBell Labs 7 m; (43 m) NRAOGreenBank 43 m; (1.4 m) ESO 1.4 mCoude Ò;
(30 m) IRAM Pico Veleta 30 m; (15 m) ESO SEST 15 m;.
References.---(1) Weliachew 1971; (2) Gardner &Whiteoak 1974; (3) Rickard
et al. 1975; (4) Solomon & de Zafra 1975; (5) Churchwell et al. 1977; (6) Rickard
et al. 1977; (7) Thompson et al. 1978; (8) Martin & Ho 1979; (9) Stark & Wolff
1979; (10) Whiteoak et al. 1980; (11) Henkel & Bally 1985; (12) Seaquist & Bell
1986; (13) Magain &Gillet 1987; (14) Henkel et al. 1987; (15) Henkel et al. 1988;
(16) Mauersberger et al. 1990; (17) NguyeníQíRieu et al. 1989; (18) Nguyení
QíRieu et al. 1991; (19) Mauersberger & Henkel 1989; (20) Sage et al. 1991;
(21) Johansson 1991; (22) Mauersberger & Henkel 1991; (23) Mauersberger et al.
1991; (24) Henkel et al. 1993; (25) Mauersberger et al. 1995; (26) Sage & Ziurys
1995; (27) Chin et al. 1996; (28) Chin et al. 1999; (29) HeikkilaØ et al. 1999;
(30) MartÐÒn et al. 2005; (31) Fuente et al. 2000; (32) MartÐÒn et al. 2003; (33) Usero
et al. 2004; (34) this work.
MARTI Ò Ð
ÒN ET AL.
458 Vol. 164

Millar 2002 and references therein). The fact that a high enrichí
ment of deuterated molecules is also observed in hot molecular
cores can be explained by evaporation of molecules from grain
mantles (e.g., Pineau des Fore Óts et al. 1989). While a value of
(3 # 0:4) ; 10 #5 has been estimated for the primordial D/H ratio
( Burles 2002 and references therein), a lower ratio of (1:4 #
0:24) ; 10 #5 is found in the Galactic neighborhood, and variaí
tions are shown on larger scales ( VidalíMadjar 2002). In molecí
ular clouds close to the center of the Milky Way, Jacq et al.
(1999) find DCN and DCO + enhancements that are compatible
with D/H an order of magnitude smaller than toward the local iní
terstellar medium. This indicates that the gas in the Galactic cení
ter region is highly processed. The enhancement of deuterated
molecules in the Large Magellanic Cloud, on the other hand, is
compatible with the local value of the D/H ratio (Chin et al. 1996).
No deuterated molecules have been detected so far toward the
central regions of other galaxies. Mauersberger et al. (1995) give
an upper limit for the DCN/HCN abundance ratio of 4 ; 10 #3
toward NGC 253.
Our survey includes the J ® 2 1 lines of DCO + , DCN, DNC,
and N 2 D + . All other 2 mm lines of deuterated molecules are exí
pected to be much weaker. Table 4 shows the frequencies of these
lines, the rms noise and limits or values of the integrated intení
sities. DCO + and DCN are not detected. There is a tentative feaí
ture at the frequency of DNC and a more significant feature at the
line frequency of N 2 D + . At the frequencies of DNC J ® 2 1 no
line from another molecule is expected to be strong enough to
explain the observed emission. Close to the frequencies of N 2 D +
J ® 2 1 is the transition frequency of OC 34 S J ® 13 12 (at
154.242 GHz). However, from our detections of OCS and our
estimates of the 32 S/ 34 S ratio (MartÐÒn et al. 2005), this emission is
expected to be an order of magnitude weaker than what is obí
served. Another potentially blending line is the 7 1;6 6 1;5 ; v 6 ® 1
transition of vibrationally excited HNCO (at 154.228 GHz). This
line has been detected toward the Galactic hot core G10.47+0.03
with an intensity of about 1/5 that of the corresponding HNCO
ground state transition, compatible with a kinetic temperature of
about 380 K (Wyrowski et al. 1999). The feature we are observí
ing here is about as strong as the HNCO 7 1;6 6 1;5 ground vibraí
tional state transition ( Table 1). If the feature at 154.228 GHz
comes from vibrationally excited HNCO, this would either indií
cate a higher kinetic temperature or maser emission. Both would
be remarkable in view of the large volume observed in our beam
toward NGC 253. While all this suggests that we have detected,
for the first time, one or two deuterated molecules beyond the
Magellanic Clouds, we will nevertheless regard the measured
features as tentative.
Table 4 gives limits to the [DX]/[HX] column density ratios
for HCO + , HCN, and HNC based on our limits for the J ® 2 1
transitions of the deuterated isotopomers, the integrated intení
sities of the 13 C bearing isotopomers measured in this survey,
and the 12 C/ 13 C ratio of 40 obtained by Henkel et al. (1993). We
have assumed that the excitation conditions of the deuterated and
main species are similar. Intensities were corrected to account for
beam dilution assuming a source size of 20 00 (see also x 3.3). For
N 2 D + , we have only the N 2 H + J ® 1 0 line for a comparison
( Mauersberger & Henkel 1991). The value for the N 2 D + / N 2 H +
column density ratio given in Table 4 assumes optically thin emisí
sion, making the same assumptions on the excitation as they have
been made for DCN in the Appendix of Mauersberger et al.
(1995).
We obtain limits to the [DX]/[HX] values of 4 ; 10 #3 for
HCO + , of 1 ; 10 #3 for HCN, and 11 ; 10 #3 for HNC. These limí
its are 2 orders of magnitude above the local interstellar D/H raí
tio and 1 order of magnitude larger than the [DCN]/[HCN] ratios
estimated for Sgr B2 (Jacq et al. 1999), but 1 order of magnitude
lower than the values determined by Chin et al. (1996) in the
Large Magellanic Cloud. The limits are consistent with the high
kinetic temperatures determined toward the molecular interstellar
medium of NGC 253. We can exclude the scenario found in many
Galactic hot cores, where recent grain mantle evaporation leads to
high D/H ratios that were preserved from a time when the temí
perature of the gas was much lower than it is now.
4.3. Alcohols in NGC 253
In the nuclear region of the Milky Way widespread large abuní
dances of the methanol (CH 3 OH) and ethanol (C 2 H 5 OH) alcohols
have been reported (RequenaíTorres et al. 2005; MartÐÒníPintado
et al. 2001). This fact has motivated a careful search for the presí
ence of C 2 H 5 OH in the 2 mm scan. Table 5 shows the frequencies
and energies of its 14 brightest 2 mm transitions assuming T rot ®
12 K , as that derived for methanol (see Table 2). We have comí
puted the rms of the noise for 20 km s #1 velocity resolution
whenever the transition appeared not to be contaminated by other
spectral features. From this rms, we estimate a 3 # detection
limit to the integrated intensity of the transitions. With the limit
to the detection of the 152 GHz transition we derive an upper limit
TABLE 4
Parameters of Selected Deuterated Species
Molecule
Transition
J J 0
#
( MHz)
rms a
(mK)
R TMB dv b
( K km s #1 ) [ DX]/[HX ]
DCO + ................ 2--1 144077.3 1.9 <0.30 <4 ; 10 #3c
DCN .................. 2--1 144827.9 1.1 <0.18 <1 ; 10 #3c
DNC .................. 2--1 152609.7 3.1 0.49 #1 ; 10 #2c
N 2 D + .................. 2--1 154217.0 2.8 0.57 #1.4 ; 10 #3d
Note.---See x 4.2 for details.
a rms computed for 20 km s #1 wide channels.
b 3 # limit to the integrated intensity assuming a line width of 150 km s #1 .
DNC and N 2 D + detections are tentative.
c Using the 13 C bearing species measured in this survey and assuming
12 C/ 13 C # 40 ( Henkel et al. 1993).
d Using the J ® 1 0 N 2 H + transition from Mauersberger & Henkel (1991).
TABLE 5
Limits to the Detection of C 2 H 5 OH
#
( MHz)
E u
(cm #1 )
rms a
(mK)
R TMB dv b
( K km s #1 )
129665.7.................................. 16.6 3.0 <1.440
131103.0.................................. 16.6 1.8 <0.864
131502.7.................................. 20.1 C 2 S
132935.1.................................. 6.9 Baseline
133323.4.................................. 16.5 CH 2 NH
142285.0.................................. 25.8 1.8 <0.864
144734.0.................................. 9.3 C 34 S
147427.4.................................. 21.0 2.9 <1.392
152656.8.................................. 9.3 3.4 <1.632
159781.8.................................. 12.2 NH 2 CN
160699.0.................................. 31.4 i = H 2 S
161609.0.................................. 26.1 NS
173391.3.................................. 12.2 3.6 <1.728
174232.9.................................. 15.7 cíC 3 H
a rms computed for 20 km s #1 wide channels. If the transition is blended,
the species is indicated.
b 3 # limit to the integrated intensity assuming an emission width of
150 km s #1 .
2 mm LINE SURVEY OF NGC 253 459
No. 2, 2006

to the ethanol column density of N (C 2 H 5 OH) <1:8 ; 10 14 cm #2 .
This limit to the column density can be improved if we add up all
the transitions that appear not to be blended. We estimate a total
column density limit of N (C 2 H 5 OH) < 9:6 ; 10 13 cm #2 .
The methanolítoíethanol abundance ratio derived for NGC
253 is #10. This lower limit is still a factor of 2 below the ratios
derived for the GC clouds by RequenaíTorres et al. (2005). Thus,
further highísensitivity observations are required to establish if
the methanolítoíethanol abundance ratio is similar to that found
in the GC clouds, as expected if alcohols are also ejected from
icy grain mantles by shocks in this galaxy (Charnley et al. 1995).
4.4. The Nuclear Environment in NGC 253
The detection of new molecules in the extragalactic ISM, far
from being a mere census of known species, provides a complete
description of the chemical complexity within the nuclear ení
vironment of galaxies and a tool to establish the origin of this
rich chemistry and the heating of the molecular gas. In the folí
lowing subsections we will compare the observed abundances in
NGC 253 with those available for five external galaxies (x 4.4.1)
and five prototypical sources within the Milky Way (x 4.4.2).
The comparison focuses mainly on the abundance ratios beí
tween the different molecular species compared to those obí
served toward NGC 253. This will minimize the systematic
error in the absolute fractional abundances due to the H 2 column
density determinations. The relative abundances between moleí
cules will provide us with a reliable chemical differentiation giving
us insights into the dominant chemical mechanism acting in each
source.
4.4.1. Comparison with Other Galaxies
We have selected the nearby galaxies NGC 4945, M82, and IC
342, where most of the known molecular species have also been
observed, and Maffei 2 and NGC 6946, for which fewer species
have been detected.
Table 6 presents the integrated intensities of the 13 CO J ®
2 1 transition that has been used to estimate the H 2 sourceí
averaged column density. We used the conversion N ( H 2 ) ® 3:3 ;
10 20 cm #2 I ( 13 CO 2 1 ) (Mauersberger et al. 2003).
Table 7 shows the derived fractional abundances relative to H 2
toward the selected galaxies. Because of the frequency coverage
of our survey, not all the species listed in Table 3 could be meaí
sured. Therefore, in Table 7, we have also included data of adí
ditional key molecular species from other studies, namely NH 3 ,
N 2 H + , and CN. H 2 and CO are not included as they are tracers of
the total mass of the molecular material content. Species such as
OH and H 2 O have also been detected toward NGC 253 and many
other galaxies. However, H 2 O exhibits maser emission that does
not allow us to determine directly the column densities, while
OH shows extremely complex behavior involving a mixture of
absorption, thermal and maser emission (e.g., Turner 1985; Ho
et al. 1987; Frayer et al. 1998; Bradford et al. 1999; Henkel et al.
2004; Goicoechea et al. 2005). Unfortunately, there is not enough
data available on CH, CH + , HCO, CO + , and HOC + for this
comparison.
In order to achieve the highest homogeneity in the derived
abundances, we have collected all the line profile parameters
from the available observations in the literature. From these valí
ues we have made our own estimates of the column densities for
each molecule and each source with the assumptions used for
NGC 253.
We have assumed a source size # s ® 20 00 for all these sources
(see, e.g., x 3.3 for NGC 253 and Wang et al. 2004 for NGC
4945). This is a fairly simple assumption for some of the targets
and may lead to considerable errors in the column density deterí
mination. As this factor is also included in the sourceíaveraged
column density determination of the molecular hydrogen, the erí
ror in the resulting fractional abundance will, however, cancel
out to first order.
The top panel in Figure 6 shows a graphical representation
of the data in Table 7. Molecular species are ordered based on
increasing abundances in NGC 253. We have also considered
the logarithmic difference of abundances in each selected galaxy
with respect to the abundances derived in this work for NGC 253
as #X ® log 10 (X /X NGC 253 ). This variable, #X , is plotted in the
bottom panel of Figure 6. The difference in the Yíaxis between
two species provides a measure of their abundance ratio relative
to that in NGC 253.
M82 shows relative abundances quite similar to those of NGC
253 in many but not all of the observed species. Evident chemí
ical differences between both nuclear starbursts were already
analyzed (Mauersberger &Henkel 1993; Takano et al. 1995, 2002;
Wang et al. 2004) and have been interpreted as a difference in
the evolutionary stage of both starbursts. Compared to that in NGC
253, the nuclear starburst in M82 represents an evolved state where
photodissociation regions ( PDRs) dominate the heating of the
molecular material (GarcÐÒaíBurillo et al. 2002). From our comí
parison, the high relative abundance of C 3 H 2 , a common PDR
tracer, is particularly outstanding. On the other hand, we note reí
markably low abundances of CH 3 OH and SiO, which are thought
to be ejected from grain mantles into the gas phase (Charnley et al.
1995; MartÐÒníPintado et al. 1992; RequenaíTorres et al. 2005).
This suggests that shocks do not play the dominant role within
the nuclear region of M82. Molecules such as HNCO, apparí
ently associated with shock chemistry ( Zinchenko et al. 2000),
also show low relative abundances compared to those in NGC
253. Furthermore, the nuclear environment of M82 is characterized
by systematically lower relative abundances of CH 3 CN, NH 3 ,
and N 2 H + . It is also interesting to note how other species used
as PDR tracers, namely C 34 S and C 2 H ( Meier & Turner 2005),
do not show high relative abundances in M82 compared to the
other observed species. Moreover, they are a factor of #2--3 less
abundant than most other the species. If we compare the M82
relative abundances with those of the Orion Bar in x 4.4.2, we find
that all the species with low abundances in M82 show the same
behavior in the Orion Bar. This comparison strongly supports the
idea of M82 as an evolved nuclear starburst largely dominated by
PDRs.
In NGC 4945, relative abundances of all the detected moí
lecular species resemble within a factor of 3 those of NGC 253.
TABLE 6
Molecular Hydrogen Sourceíaveraged Column Density
Source
Transition
J J
R TMB dv
( K km s #1 )
# beam
(arcsec)
N H 2
(cm #2 )
NGC 253......... 13 CO 2--1 52.4 a 34 6.7 ; 10 22
NGC 4945....... 13 CO 2--1 81.2 b 23 6.4 ; 10 22
M82 ................. 13 CO 2--1 29.4 a 34 3.8 ; 10 22
IC 342 ............. 13 CO 2--1 17.2 a 34 2.2 ; 10 22
Maffei 2........... 13 CO 2--1 13.9 a 34 1.8 ; 10 22
NGC 6946....... 13 CO 2--1 22.2 c 21 1.5 ; 10 22
Note.---A 20 00 source size was assumed for all targets.
a 10 m HHT by Mauersberger et al. (2003).
b 15 m SEST by Wang et al. (2004).
c 30 m IRAM by Israel & Baas (2001).
MARTI Ò Ð
ÒN ET AL.
460 Vol. 164

The nondetection of SiO indicates a clear underabundance, simí
ilar to what is seen toward M82. On the other hand, unlike in
M82, CH 3 OH shows an abundance similar to that in NGC 253.
This resemblance between the abundances in NGC 4945 and
NGC 253, altogether with excitation considerations and a detailed
study of species such as CN,HCN, and HNC, has been interpreted
as an intermediate evolutionary state, between NGC 253 and M82
(Wang et al. 2004). From our comparison, the nuclear starburst of
NGC 4945 shows much closer resemblance to NGC 253 than to
M82.
IC 342 shows high relative abundances of HC 3 N and H 13 CO + ,
while molecules such as C 2 H and N 2 H + have lower abundances
as compared with NGC 253. Many of the available observations
for IC 342 (marked by an asterisk in Table 7) were obtained near
the D position as labeled by Downes et al. (1992), #15 00 north of
the main nuclear star formation region. The abundance of C 2 H
given in Table 7 has been taken from Meier & Turner (2005) toí
ward the D position. Its relative abundance is low compared to
that in NGC 253. Even taking the C 2 H abundance at the position
where this molecule peaks (i.e., a value an order of magnitude
higher toward the central position, Meier & Turner 2005), the
relative abundances of this species with respect to the rest of
molecules would be similar to that in NGC 253. Highíresolution
maps from Meier & Turner (2005) clearly differentiate a region
dominated by UV radiation from the central nuclear cluster and
the region where the chemistry is dominated by shocks produced
by collisions of clouds in a barred potential. Given the position
toward which many of the spectra were taken, away from the
dynamical center, the molecular emission of IC 342 should
resemble more that of NGC 253 than that of M82. In fact IC 342
does not show the relatively low HNCO abundance seen in M82
but does show a similar low relative abundance of C 2 H. Still,
there is a large number of molecules not yet observed that are
critical to fully evaluate the chemistry of IC 342, such as SiO, SO,
OCS, and CH 3 OH.
As far as Maffei 2 and NGC 6946 are concerned, the available
molecular information is not sufficient to establish a clear chemí
ical differentiation from the other sources. What can be noted is a
remarkably high relative abundance of HNCO in Maffei 2 comí
pared to the other galaxies and a low relative abundance of H 2 CO
in NGC 6946. Both galaxies also show low relative abundances
of N 2 H + .
In order to quantify this comparison, Table 8 shows the mean
and rms of #X for each source for two cases. The first only uses
the species detected in both sources. The number of molecules
included in the calculations is also indicated. The second also
uses upper limits to the abundances derived for the comparií
son sources. To take these limits into account, we followed two
TABLE 7
Molecular Fractional Abundances toward NGC 253 and Other Galaxies
Molecule NGC 253 NGC 4945 M82 IC 342 Maffei 2 NGC 6946
HN 13 C .................................. #10.6 #9.5 #9.5 . . . . . . . . .
H 13 CO + ................................. #10.4 #10.0 #9.9 #9.5 a . . . . . .
H 13 CN .................................. #9.9 #9.7 <#9.9 <#9.4 a . . . . . .
SiO ....................................... #9.9 <#9.9 <#9.9 <#9.3 a <#9.1 . . .
CH 3 CN ................................. #9.5 . . . <#9.7 . . . . . . . . .
C 34 S...................................... #9.4 #9.4 #9.3 #9.5 a
. . . <#9.0
cíC 3 H 2 .................................. #9.3 #8.9 #8.1 . . . . . . . . .
HC 3 N.................................... #9.2 #8.8 #8.7 #8.6 . . . . . .
N 2 H + ..................................... #9.2 #9.2 #9.4 #9.6 a
#9.7 #9.6
HNC ..................................... #9.0 #8.6 #8.8 #8.7 #8.7 #8.8
SO ........................................ #8.9 #8.7 <#8.5 <#8.7 . . . . . .
HCO + ................................... #8.8 #8.4 #8.4 #8.8 #8.6 #8.5
HNCO .................................. #8.8 #8.4 <#8.8 #8.6 #8.3 . . .
H 2 CO.................................... #8.6 #8.1 #8.2 #9.4 #8.8 #9.1
OCS...................................... #8.4 <#7.5 #7.9 <#8.6 . . . . . .
CN ........................................ #8.3 #7.7 #8.2 #8.3 a . . . . . .
HCN ..................................... #8.3 #8.3 #8.4 #8.5 a
#8.3 #8.5
CS......................................... #8.2 #8.4 #8.2 #8.4 a
#8.7 #8.7
CH 3 CCH .............................. #8.2 #8.0 #7.7 . . . . . . . . .
CH 3 OH................................. #7.9 #7.4 <#8.3 <#7.8 #7.7 #7.9
C 2 H....................................... #7.7 #7.3 #7.6 <#8.4 . . . . . .
NH 3 ...................................... #7.2 . . . #8.4 #7.3 #7.1 . . .
a Observations taken #15 00 away from the central position (see x 4.4.1 for details).
References.---For NGC 253: ( NS, NO) this paper, MartÐÒn et al. 2003; (OCS) MartÐÒn et al. 2005; ( NH 3 ) Mauersberger et al. 2003;
(CN ) Henkel et al. 1988; ( N 2 H + ) Mauersberger & Henkel 1991. For NGC 4945: (SiO) Wang et al. 2004, Henkel et al. 1994. For
M82: (CS, C 34 S) Mauersberger & Henkel 1989; ( N 2 H + , SiO, H 13 CO + ) Mauersberger & Henkel 1991; (C 3 H 2 , CH 3 CN, CH 3 CCH )
Mauersberger et al. 1991; (OCS) Mauersberger et al. 1995; ( H 2 CO, CH 3 OH ) Hu Ø ttemeister et al. 1997; (CN, HC 3 N, C 2 H ) Henkel
et al. 1988; ( HCO + , HCN, H 13 CN ) Henkel et al. 1998; (SO) Petuchowski & Bennett 1992; ( HCNO) NguyeníQíRieu et al. 1991;
( HCN, HCO + , H 13 CO + ) NguyeníQíRieu et al. 1992; ( HCN, HCO + ) Wild 1990; ( NH 3 ) Wess et al. 2001; ( HNC ) Hu Ø ttemeister et al.
1995. For IC 342: (CS) Mauersberger & Henkel 1989; ( N 2 H + , SiO, H 13 CO + ) Mauersberger & Henkel 1991; (C 34 S, OCS, HC 3 N,
H 2 CO) Mauersberger et al. 1995; (CN, HNC, HC 3 N )Henkel et al. 1988; ( HNC ) Hu Ø ttemeister et al. 1995; ( H 2 CO, CH 3 OH )
Hu Ø ttemeister et al. 1997; (SO) Petuchowski & Bennett 1992; ( HCN, HCO + ) Wild 1990; ( HCNO) NguyeníQíRieu et al. 1991;
( HCN, HCO + , H 13 CO + ) NguyeníQíRieu et al. 1992; ( HCN, H 13 CN ) Henkel et al. 1998; (C 2 H ) Meier & Turner 2005; ( NH 3 )
Mauersberger et al. 2003. For Maffei 2: (CS) Mauersberger & Henkel 1989; (SiO, N 2 H + ) Mauersberger & Henkel 1991; ( H 2 CO,
CH 3 OH ) Hu Ø ttemeister et al. 1997; ( HCNO) NguyeníQíRieu et al. 1991; ( HCO + , HCN ) NguyeníQíRieu et al. 1992; ( NH 3 ) Henkel
et al. 2000, Mauersberger et al. 2003. For NGC 6946: ( H 2 CO, CH 3 OH ) Hu Ø ttemeister et al. 1997; (CS, C 34 S) Mauersberger &
Henkel 1989; ( HCO + HCN ) NguyeníQíRieu et al. 1992; ( N 2 H + ) Mauersberger & Henkel 1991.
2 mm LINE SURVEY OF NGC 253 461
No. 2, 2006

Fig. 6.---Molecular abundances (X ) of selected extragalactic sources compared to those of NGC 253. Arrows represent upper limits. A solid line connects the
abundances in NGC 253. Top panel: Logarithmic fractional abundances relative to H 2 as presented in Table 7. Bottom panel: Logarithmic abundances for each
species normalized to that measured in NGC 253.

criteria: upper limits higher than the abundances measured toward
NGC 253 were rejected as they do not provide useful information
for the comparison; for the calculations, the other observed limits,
lower than the abundances in NGC 253, were divided by 2 as an
estimated true value. The second assumption may overestimate
some abundances, but any other constraint might result in biased
guesses. While the mean value, #X , only gives an the overall abí
solute difference in the abundances of each source as compared
with NGC 253, the rms of #X provides us with a fairly reliable
measurement of differences in the abundance pattern that defines
the chemistry of each object. M82, as pointed out before, clearly
differs chemically from NGC 253, showing the highest rms valí
ues. NGC 4945 and IC 342 are much closer to NGC 253.
4.4.2. Comparison with Prototype Galactic Sources
We have selected five Galactic sources, namely Sgr B2(N),
Sgr B2(OH), TMCí1, L134N, and the Orion Bar, which are
considered to be prototypes of their respective chemistry within
the Galaxy. All these sources have been the target of multitraní
sition studies and some of them of frequency surveys. Table 9
shows the abundances of all the molecules observed in the NGC
253 survey compared to those of these Galactic sources. In order
to achieve consistency among ratios, we have tried, when availí
able, to use values from line surveys or multiline studies for each
source.
In the same way as in the extragalactic comparison, the top
panel in Figure 7 shows a graphical representation of the data in
Table 9, where molecular species are ordered based on increasí
ing NGC 253 abundances. The bottom panel in Figure 7 shows
the logarithmic abundance of each molecule normalized to the
abundance measured in NGC 253.
Sgr B2(N) is believed to represent the prototype of a hot moí
lecular core chemistry associated with massive star formation near
the nucleus of the Milky Way. A large number of molecular speí
cies have been identified toward this source, which is the brightest
molecular line emitter within the Galaxy (Friedel et al. 2004 and
references therein). There are clear differences between the abuní
dances found toward Sgr B2(N) and those in NGC 253. A set of
molecules, namely CH 3 CN, HC 3 N, NS, H 2 CS, SO 2 , CH 2 NH, SO,
CH 3 OH, and NO, show abundances relative to the other species
#2--3 orders of magnitude larger than what we observe toward
TABLE 8
Statistical Comparison of Abundances of Selected Galaxies
with Those in NGC 253
Without Limits With Limits
Source #X rms(#X ) Number #X rms(#X ) Number
NGC 4945...... 0.30 0.29 18 0.27 0.32 19
M82 ................ 0.19 0.51 15 0.05 0.54 19
IC 342 ............ 0.08 0.34 12 0.00 0.42 13
Maffei 2.......... 0.02 0.31 9 0.02 0.31 9
NGC 6946...... #0.20 0.30 7 #0.20 0.30 7
Notes.---#X ® ( log 10 X /X NGC 253 ). See last paragraph in x 4.4.1 for details.
TABLE 9
Comparison of NGC 253 Fractional Molecular Abundances with Those in Selected Galactic Sources
Molecule NGC 253 Sgr B2(N ) Sgr B2(OH ) TMCí1 L134N Orion Bar
HN 13 C .......................................... #10.6 #11.0 . . . . . . . . . . . .
H 13 CO + ......................................... #10.4 #11.4 . . . . . . . . . #10.3
SiO ............................................... #9.9 #10.7 . . . <#11.6 <#11.4 #10.3
NH 2 CN......................................... #9.7 #10.1 #10.0 . . . . . . . . .
C 2 S ............................................... #9.7 . . . #9.6 #8.1 #9.2 . . .
CH 3 CN ......................................... #9.5 #6.7 #9.4 #9.0 <#9.0 <#10.3
cíC 3 H ........................................... #9.5 #10.5 <#10.9 #9.3 . . . . . .
HOCO + ......................................... #9.4 #10.5 #9.7 . . . . . . . . .
C 34 S.............................................. #9.4 #10.2 . . . . . . . . . #9.0
cíC 3 H 2 .......................................... #9.3 #10.5 #9.8 #8.0 #8.7 #9.7
HC 3 N............................................ #9.2 #7.5 #9.0 #8.2 #9.7 . . .
NS ................................................ #9.2 #7.0 . . . #9.1 #9.5 . . .
H 2 CS ............................................ #9.2 #6.8 #8.7 #8.5 #9.2 . . .
SO 2 ............................................... #9.1 #6.6 #8.7 <#9.0 #8.4 #9.9
CH 2 NH......................................... #9.1 #7.0 #9.2 . . . . . . . . .
H 2 S ............................................... #9.1 #9.9 . . . <#9.3 #9.1 #8.2
HNC ............................................. #9.0 . . . . . . #7.7 #8.2 #9.0
SO ................................................ #8.9 #6.9 #8.7 #8.3 #7.7 #8.0
HCO + ........................................... #8.8 . . . . . . #8.1 #8.1 #8.5
HNCO .......................................... #8.8 #9.2 #8.4 #9.7 . . . <#10.8
H 2 CO............................................ #8.6 #9.3 #8.6 #7.7 #7.7 #8.2
OCS.............................................. #8.4 #8.6 #8.3 #8.7 #8.7 . . .
HCN ............................................. #8.3 . . . . . . #7.7 #8.4 #8.3
CS................................................. #8.2 . . . . . . #8.0 #9.0 #7.6
CH 3 CCH ...................................... #8.2 #8.4 #8.8 #8.2 <#8.9 . . .
CH 3 OH......................................... #7.9 #5.8 #7.3 #8.7 #8.5 #9.0
C 2 H............................................... #7.7 #9.7 . . . #7.1 <#7.3 #8.7
NO................................................ #7.2 #6.0 . . . <#7.5 #7.2 #8.6
References.---For NGC 253: ( NS, NO) this paper, MartÐÒn et al. (2003); (OCS) MartÐÒn et al. (2005). For Sgr B2N: ( H 2 S) Nummelin
et al. (2000), Minh et al. (1991). N ( H 2 ) ® 3:0 ; 10 24 cm #2 from Nummelin et al. (2000). For Sgr B2(OH ): Cummins et al. (1986), Turner (1991).
N ( H 2 ) ® 5 ; 10 23 cm #2 (see text for details on this estimate). For TMCí1 and L134N: ( NS) Ohishi et al. (1992), McGonagle et al. (1994); (SiO)
Ziurys et al. (1989). N ( H 2 ) ® 1:0 ; 10 22 cm #2 from Ohishi et al. (1992). For Orion Bar: (C 3 H 2 ) Jansen et al. (1995), Fuente et al. (2003); (SiO)
Schilke et al. (2001). N ( H 2 ) ® 6:5 ; 10 22 cm #2 from Jansen et al. (1995).
2 mm LINE SURVEY OF NGC 253 463

Fig. 7.---Molecular abundances (X ) of selected Galactic sources compared to those of NGC 253. Arrows represent upper limits. A solid line connects the
abundances in NGC 253. Top panel: Logarithmic fractional abundances relative to H 2 as shown in Table 9. Bottom panel: Logarithmic abundances for each species
normalized to that measured in NGC 253.

NGC 253. Almost all the other species observed toward Sgr
B2(N) have smaller abundances than in NGC 253, but their relaí
tive abundances are, within an order of magnitude, similar to those
in NGC 253. Exceptions are HOCO + and C 2 H, which show lower
relative abundances by factors of #30 and 100, respectively, in the
Sgr B2(N) hot core. Moreover, the rotation temperatures derived
from the observed molecules in NGC 253 are low compared to the
typical temperatures of >70 K observed toward hot cores. Even if
hot cores associated with massive protostars are present in the
nuclear environment of NGC 253, they clearly do not dominate
the molecular emission. While the angular size of the Sgr B2 moí
lecular cloud complex is #18 0 (Scoville et al. 1975), the emission
of some large complex molecules toward Sgr B2(N) appears coní
centrated within a region of #5 00 , corresponding to 0.2 pc at a disí
tance of 8.5 kpc ( Kerr & LyndeníBell 1986). At the distance of
NGC 253, such a source would have a diameter of <0B02, and
because of beam dilution, its detection is certainly out of reach
for the 30 m telescope. A large number of such sources would
be needed in order to be able to observe traces of their complex
chemistry.
The Orion Bar has been selected as the prototype of photoí
dissociation regions ( PDRs). A large fraction of the molecular
gas in the Milky Way and in external galaxies is expected to be
affected by PDRs ( Hollenbach & Tielens 1997). In PDRs, moí
lecular emission originates from the surface layers of interstellar
molecular clouds exposed to intense faríultraviolet photons from
nearby young OB stars. The nondetection of molecules such as
CH 3 CN and HNCO toward the Orion PDR clearly indicates a low
relative abundance of these species with respect to other molí
ecules, which is in contrast with the observed NGC 253 abuní
dances. Similarly, the relative abundances of SiO, SO 2 , CH 3 OH,
C 2 H and NO appear to be #1--2 orders of magnitude lower in the
Orion bar than in NGC 253. High abundances of hydrocarbon
radicals, such as C 2 H and C 3 H 2 , are observed toward PDRs
( Fosse Ò et al. 2000; Fuente et al. 1993, 2003; Pety et al. 2005). In
this context, it is surprising to see how C 2 H and C 3 H 2 , comí
monly used as PDR tracers, show relatively low abundances in
the Orion PDR when compared not only with NGC 253 but also
with the prototypical Galactic sources except the Sgr B2(N) hot
core. The observations of the Orion Bar have been taken from
the position of maximum H 2 column density (Jansen et al. 1995;
Fuente et al. 1996), which does not match the position of highest
C 3 H 2 abundance (Fuente et al. 2003). However, the contribution
of the position of the C 3 H 2 emission peak to the total column dení
sity is smaller than that of the H 2 column density peak. Therefore,
it is expected that the molecular abundance of the H 2 peak domí
inates the emission of PDRs, also in external galaxies. Anyhow,
even if we consider the position of the larger C 3 H 2 abundance (up
to 1 order of magnitude; Fuente et al. 2003), the relative abuní
dance would still be close to that found in NGC 253. From this
comparison as well as that to M82 (see x 4.4.1) it is clear that the
chemistry of the nuclear environment of NGC 253 is not domií
nated by photodissociation.
Quiescent cold dark clouds are represented by TMC1 and
L134N. The molecular composition at sites of lowímass star forí
mation is dominated by gasíphase ionímolecule chemistry due
to a lack of embedded luminous sources. In our comparison,
both dark cloud complexes have in common low relative abuní
dances of SiO and, to a lesser degree, of CH 3 OH. The remaining
species behave quite differently in each source. SO 2 , H 2 S, HNCO,
and NO show low abundances in TMCí1, while C 2 S and C 3 H 2
appear to have abundances larger than in NGC 253 by up to
1 order of magnitude. On the other hand, L134N shows a high
SO abundance and relatively low abundances of HC 3 N and
CH 3 CCH. The abundance pattern defined by these two dark clouds
clearly does not follow the relative abundances found toward
NGC 253. In addition, even the rotation temperatures derived
for NGC 253, which in general are smaller than T kin , are higher
than the typical kinetic temperature (T kin # 10 K) measured toí
ward dark clouds.
Sgr B2(OH), located at the southern end of the Sgr B2 moí
lecular envelope, is free from the emission stemming from the
hot cores in this cloud. This position in the Sgr B giant molecular
cloud is taken as the prototype of Galactic center (GC) molecular
cloud complexes ( MartÐÒníPintado et al. 1997). The H 2 column
density of this source has been estimated to be roughly N ( H 2 ) ®
5 ; 10 23 ( Nummelin et al. 2000). The relative abundances of all
the species measured toward this Galactic center molecular cloud
complex appear to closely follow the abundance pattern in NGC
253 within a factor #5. The only species clearly underabundant in
Sgr B2(OH) when compared to NGC 253 is cíC 3 H, but this molí
ecule is only tentatively detected toward NGC 253. Thus, the raí
tios between any of the observed species are similar to those in
NGC 253 within 1 order of magnitude.
In analogy to our statistical comparison of molecular abuní
dances in prototypical nearby galaxies shown in Table 8, Table 10
presents a comparison between NGC 253 and the selected Galací
tic sources. To calculate the parameters, the available abundances
in Table 9 have been taken into account following the criteria
explained at the end of x 4.4.1. Given the uncertainty of its detecí
tion, cíC 3 H was not included. The computed rms of the logarithí
mic differences shows the envelope of the Sgr B2 complex as the
one most closely resembling NGC 253.
The agreement in relative abundances between NGC 253 and
a typical Galactic center cloud is a remarkable and clear indií
cation that the heating and chemistry of the bulk of the ISM is
dominated by the same processes. The molecular envelope around
the star forming region Sgr B2 is formed by a relatively dense
and warm material (n H 2 ® 2 ; 10 5 cm #3 and T kin ® 40 60 K;
de Vicente et al. 1997). The large abundances of NH 3 , SiO, and
C 2 H 5 OH observed in the gas phase of the Galactic center moí
lecular clouds are claimed to be the result of grain chemistry and
subsequent ejection onto the gas phase due to the disruption of
grains by lowívelocity shocks (v < 20 km s #1 , Flower et al.
1995; MartÐÒníPintado et al. 1997, 2001). In the Galactic center,
the possible origin for these shocks is so far unclear. It has been
claimed that they could be produced by the largeíscale shocks
associated with cloudícloud collisions due to the orbital motion
of molecular clouds in a barred potential ( Hasegawa et al. 1994;
Sato et al. 2000) and the interaction of expanding bubbles due to
supernova events and /or strong stellar winds from WolfíRayet
TABLE 10
Statistical Comparison of Abundances of Selected Galactic Sources
with Those in NGC 253
Without Limits With Limits
Source #X rms(#X ) Number #X rms(#X ) Number
Sgr B2(N )........ 0.40 1.50 22 0.40 1.50 22
Sgr B2(OH ) ..... 0.07 0.40 16 #0.04 0.58 17
TMCí1.............. 0.47 0.66 18 0.26 0.85 21
L134N .............. 0.18 0.58 16 #0.00 0.76 18
Orion Bar ......... #0.07 0.71 14 #0.28 0.85 17
Note.---See Table 8 and x 4.4.2 for details.
2 mm LINE SURVEY OF NGC 253 465

stars (Sofue 1990; de Vicente et al. 1997; MartÐÒníPintado et al.
1999) associated with a burst of star formation that occurred10 7 yr
ago ( RodrÐÒguezíFerna Òndez et al. 2004; RodrÐÒguezíFerna Òndez &
MartÐÒníPintado 2005).
5. CONCLUSIONS
1. We present the first unbiased molecular line survey of an
extragalactic source. The survey covers the 2 mm atmospheric
window from 129.1 to 175.2 GHz toward the inner 200 pc of
NGC 253. A total of 111 features are identified as transitions of
25 different molecular species, eight of which (three tentatively)
have been detected for the first time outside the Milky Way. The
rare isotopes, 34 SO and HC 18 O + , were also detected for the first
time in an extragalactic source. In addition, three hydrogen reí
combination lines and, tentatively, two deuterated species are
identified, N 2 D + , and DNC. The origin of the observed features
of both N 2 D + , and DNC is still unclear and deserves further
investigation. If real, these would represent the first deuterated
molecules observed beyond the Magellanic Clouds. Column
densities and rotation temperatures have been determined for
each species under local thermodynamic equilibrium ( LTE)
conditions.
2. As a result of this survey and by adding existing data from
three molecules, namely NH 3 , N 2 H + , and CN, we obtain the
most complete description to date of the chemical complexity
within the nuclear few hundred parsecs of a starburst galaxy. We
have performed a comparison between the chemistry of NGC
253 and those of five other outstanding nearby galaxies. This
comparison clearly shows the strong chemical differentiation beí
tween nuclei of galaxies. Most prominent differences are obí
served between NGC 253 and M82. This can be interpreted in
terms of a more evolved stage of the nuclear starburst in M82.
The chemistry of NGC 4945, although claimed to be at an interí
mediate stage of evolution between NGC 253 and M82, clearly
resembles more that of NGC 253 except for the lack of SiO.
Similarly, the position within the nuclear region of IC 342, where
most of the molecular observations have been made, show relative
abundances close to those in NGC 253. This position is far from
the main region of star formation where PDRs are claimed to
dominate the chemistry.
As far as Maffei 2 and NGC 6946 are concerned, the available
molecular observations do not yet allow us to carry out a simií
larly detailed comparison.
3. A comparison of the molecular abundances of NGC 253
with selected Galactic sources shows a striking similarity beí
tween NGC 253 and Sgr B2(OH). The latter source, located in the
envelope of Sgr B2, is commonly taken as the prototype of the
molecular cloud complexes in the Galactic center region. This
indicates that the chemistry of the nuclear molecular environí
ment in NGC 253 and the Galactic center molecular clouds are
driven by similar chemical processes. If this is the case, the chemí
istry and the heating of the nuclear molecular material in NGC 253
is dominated by largeíscale lowívelocity shocks.
4. The comparison with Galactic sources also shows a close
similarity between the abundances within the nuclear starburst of
M82 and those of the Orion Bar, the prototype of a photoní
dominated region. This resemblance fully supports the idea of
the chemistry in M82 being dominated by PDRs.
S. M. and R. M. were supported by the Programas de Acciones
Integradas between Spain and Germany. J. M.P. has been partially
supported by the Spanish Ministerio de Educacio Ò n y Ciencia
under projects ESP 2004í00665, AYA 2002í10113íE and AYA
2003í02785íE. S. G. B. acknowledges financial support from
AYA 2003í07584 and ESP 2003í04957. Wewould like to thank
the referee, whose comments helped to significantly improve the
paper.
APPENDIX A
GAUSSIAN FITTING TO BLENDED LINES
Table 1 includes a label in the third column indicating if and how the respective spectral feature is affected by blending. In the
following, the encountered different cases are described:
Blended (b).---Some observed features are composed of several transitions from different species. If the line is only partially
blended, it is still possible to separate the individual components by fitting different Gaussian profiles. For these fits, sometimes it was
necessary to fix the velocity and /or line width.
Synthetic (s).---The blending may prevent us from making a reliable multiple Gaussian fitting to an observed feature. In this case,
when we have observed several transitions of one of the species contributing to the blended feature, it is possible to use these to
estimate the species' contribution to the blended feature by interpolating or extrapolating the linear fit to the rotation diagrams (see
x B1). This is equivalent to using equation ( B6) to estimate the expected integrated intensity of a line from the derived column density
and rotation temperature. Using the estimated line intensity and the line width and radial velocity of the other observed transitions, we
subtract the expected synthetic Gaussian profile from the spectrum. The estimated integrated line intensities using this method have
been plotted in the corresponding rotation diagrams with open symbols ( Fig. 8). The transitions from other species contributing to the
blended feature have been fitted to the residual spectrum after the subtraction of the synthetic profile.
Hyperfine (hf ).---Some species show transitions that consist of a group of unresolved hyperfine components. The relative
intensities of these lines are fixed by their spectroscopic parameters, assuming LTE and optically thin line emission. Thus, the fitting
has been carried out with a comb of Gaussian profiles at the rest frequencies of the hyperfine components with the same widths and
with fixed line intensities relative to the main component, which is taken as a free parameter. The derived parameters shown in Table 1
correspond to the integrated intensity of all the hyperfine components but the frequency, line intensity and velocity refer to the main
hyperfine component.
Multitransition (m).---Some features of species such as CH 3 OH and CH 3 CN consist of groups of overlapping transitions, involving
different energy levels. In these cases, a single Gaussian profile has been fitted to the observed spectral feature. The total integrated
line intensity of all the blended transitions has been used to derive rotation diagrams as described in x B2.
MARTI Ò Ð
ÒN ET AL.
466

Fig. 8.---Rotation diagrams of detected species. When differentiated, velocity components are represented as filled squares and solid lines (180 km s #1 ) and filled
triangles and dashed lines (270 km s #1 ). Unfilled markers represent blended observed transitions whose intensities have been estimated by inter or extrapolation from
the diagrams (see Appendix A for details). The NS, NO, and OCS diagrams include additional transitions observed by MartÐÒn et al. (2003, 2005).

Fig. 8.--- Continued

APPENDIX B
ROTATION DIAGRAMS
B1. BASICS
Assuming optically thin emission for the observed molecular transitions, the column density in the upper level can be derived as
N u ®
8#k# 2
hc 3 A ul
1#
J # (T BG )
J # (T ex )
# # #1
Z T B dv; ÅB1�
where A ul , the Einstein coefficient for a transition connecting the levels u and l, can be written as A ul ® (64# 4 # 3 /3hc 3 )S# 2
ul /g u , and
J # (T ) ® (h#/k)­exp(h#/kT ) # 1# #1 , with T BG denoting the temperature of the background, T ex is the excitation temperature of the
transition, and g u is the degeneracy of the upper level.
In the RayleighíJeans approximation (h#TkT ), and assuming T ex 3T BG , equation ( B1) can be reduced to
N u ®
8#k# 2
hc 3 A ul
W ; ÅB2�
where the integrated brightness temperature
R T B dv is denoted as W. If LTE is assumed, the population distribution of all the levels
can be represented by a single temperature, the rotation temperature (T rot ), given by
N u
g u
® N
Z
e #E u =kT rot ; ÅB3�
where N is the total column density, E u is the energy of the upper level, and Z is the partition function calculated as
Z ® X 8i
g i e #E i =kT rot : ÅB4�
Fig. 8.--- Continued
2 mm LINE SURVEY OF NGC 253 469

Using equation ( B2) in equation ( B3), we obtain
N ®
8#k# 2 Z
hc 3 A ul g u
We E u =kT rot ; ÅB5�
or the equivalent logarithmic expression,
log
8#k# 2
hc 3 A ul g u
W ® log
N
Z #
log eE u
kT rot
; ÅB6�
which constitutes the basic equation for the rotation diagrams (see Blake et al. 1987, Turner 1991, and Goldsmith & Langer 1999 for a
complete description of this method). Plotting N u /g u for each transition against the energy of its upper level in a logarithmic scale, the
population distribution can be fitted by a straight line that provides the total column density divided by the partition function as well as T rot .
B2. BLENDED LINES
Whenever a molecular feature is composed of a number of unresolved blended transitions of a given molecular species, rotation
diagrams cannot be directly used to determine the physical properties of the molecular emission. The spectroscopic parameters of the
blended individual transitions within an observed feature allow us to estimate the contribution of each transition to the total integrated
intensity, assuming LTE and optically thin emission. Thus, using equation ( B5) we can express the relative integrated intensity
between two transitions with Einstein coefficients A 0 ul and A 00 ul , energies of the upper level E 0
u and E 00
u , and degeneracies g 0
u and g 00
u as
W 0
W 00 ® # 00 2 A ul g 0
u
# 0 2 A 00
ul g 00 u
e (E u #E 0 u )=k T 0
rot : ÅB7�
This allows us to compute the relative intensity of transition for a given T 0
rot . We will differentiate between the temperature T 0
rot we assume
for estimating the contribution of each transition to the blended line, and the rotational temperature T rot derived from the rotation diagram.
If we consider the case of a molecule with hyperfine structure where all the transitions within a feature have the same upper level
energy, then we can make the assumptions of E 0
u# E 00
u and #
0# # 00 . Therefore, the contribution of each transition will not depend on
the assumed T rot and will only depend on the spectroscopic parameters of each transition in the form
W 0
W 00
# # HF
®
A 0 ul g u
A 00
ul g 00 u
: ÅB8�
When the contribution of each hyperfine transition to the total line is computed and plotted on the rotation diagrams, they will all
lie in the same point in the rotation diagram (see NO and NS diagrams).
In the general case where E 0
u 6® E 00
u , the calculated contribution of each line for a given T 0
rot will lie on a straight line when included
into the rotation diagram. The line traced by these points will have the slope log 10 e/kT 0
rot .
The LTE approximation assumes that all the transitions will lie in a straight line in the rotation diagram, thus we can modify the
assumed T 0
rot used to calculate the contributions within the blended lines to obtain the best fit to a straight line to all data in the rotation
diagram. At this point we will obtain the situation in which T 0
rot# T rot .
In this paper we have several examples where this method has been applied, i.e., the CH 3 CN rotation diagram, where only two
multipleítransition lines are plotted, the CH 3 CCH diagram, where three blended lines are perfectly aligned to a straight line, and the
CH 3 OH diagram, where a mixture of multitransition lines and single lines are plotted altogether ( Fig. 8).
We have to take special care when dealing with the symmetric top rotors such as CH 3 CN and CH 3 CCH. In this case, the rotation
temperature is not sufficient to describe the relative population of the levels given that both T kin and T rot are required to evaluate the partition
function as described by Turner (1991). Thus, T rot would characterize the population of J levels within a K ladder, while T kin would
describe the population between K ladders. Our observations do not resolve the transitions of different K ladders. We assume both
temperatures to be the same (T kin ® T rot ) so that the population of all levels is described by a single rotation temperature. Given that T rot is a
lower limit to T kin , with typically T rot TT kin , this assumption may result in an overestimation of the column density derived from rotation
diagrams as the contribution to the total integrated intensity of the K # 1 levels may have been underestimated. On the other hand, the
evaluation of the partition function will also be underestimated. Thus, the overall effect in the total column density determination is not
expected to be more than a factor of 2--3, while the derived rotation temperature will only be marginally affected.
APPENDIX C
NOTES ON INDIVIDUAL MOLECULES
In the following, we will discuss the fitting procedure of the identified molecular species in detail.
C1. CARBON MONOSULFIDE: CS
Only the J ® 3 2 transition of carbon monosulfide at 147 GHz is observed in the survey. It appears to be partially contaminated by
the H35# recombination line. The estimated Gaussian profile derived from the observed H34# (see x C26) has been subtracted to
account for its contribution.
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Two rare isotopes, 13 CS and C 34 S at 144.6 and 138.8 GHz, respectively, are also observed in the spectral scan. MartÐÒn et al. (2005)
used these observations of CS and its isotopes, as well as additional CS lines, to derive isotope ratios of 32 S/ 34 S # 8 # 2 and
34 S/ 33 S > 9. From the observation of additional CS transitions at 3, 1.5, and 1.2 mm, MartÐÒn et al. (2005) derive a T rot # 10 K.
C2. NITRIC OXIDE: NO
One rotational transition of nitric oxide (150.2 and 150.5 GHZ) has been identified in our survey. MartÐÒn et al. (2003) confirmed this
identification by observing additional transitions at 250 GHz. These transitions have also been included in the rotation diagram in
Figure 8. The topology of the NO energy levels is similar to that of NS. Transitions connecting upper and lower levels are labeled # +
and # # .
It was only possible to fit the # + hyperfine group of the transition at 150.2 GHz, given that the # # group is blended with the more intense
H 2 CO line (see x C14). As shown in Table 1, the # # group of transitions has been estimated by fixing width and radial velocity of each
component to the values derived from the fitting to the # + transition and the line intensities determined by their relative line strengths.
C3. NITRIC SULFIDE: NS
One rotational transition of nitric sulfide (161.2 and 161.7 GHz) has been detected in our survey. MartÐÒn et al. (2003) confirmed the
identification by detecting other transitions at 115 and 207 GHz. The rotation diagram in Figure 8 includes these additional transitions.
This molecule presents #ídoubling and hyperfine splitting. Because of the #ídoubling, each rotational level is divided into two
levels with opposite parity. Transitions connecting the upper levels are denoted as e and those connecting the lower levels as f (see
Table 1). We are able to resolve the split rotational levels due to #ídoubling but the intrinsic line widths prevent us from resolving the
hyperfine structure.
C4. SULFUR MONOXIDE: SO
We detect three transitions of sulfur monoxide (138.2, 159, and 172.2 GHz). The velocity components at 180 and 275 km s #1 are
clearly identified. As seen in the rotation diagrams ( Fig. 8), the 275 km s #1 component of the 3 4 --2 3 transition shows a much higher
intensity than expected if we assume both components to have the same rotation temperature. This is likely due to the contamination
of this component by the 8 1,8 --7 1,7 transition of NH 2 CN. The uncertainty in the contribution of this line due to the noise strongly affects
the fitting of the highívelocity component of this SO transition.
One of the isotopic substitutions of sulfur monoxide, 34 SO (135.8 GHz) is detected blended with a SO 2 transition. This faint feature has
been fitted with a double Gaussian, with the velocity and width of each component fixed to the values derived from the same transition,
4 3 --3 2 , of the main isotope. We calculate a ratio 32 SO/ 34 SO ® 5:1 # 1:2, using both velocity components of the 4 3 #3 2 transition of both
isotopes. This ratio is consistent within the errors with 32 S/ 34 S ® 8 # 2 as derived by MartÐÒn et al. (2005) from CS data.
C5. SILICON MONOXIDE: SIO
Two transitions of silicon monoxide in the ground vibrational state are observed at 130.3 and 173.7 GHz (J ® 3 2 and J ® 4 3,
respectively). It is possible to identify two velocity components from the J ® 3 2 transition at 180 and 260 km s #1 . Velocity and
width of the fainter and more noisy J ® 4 3 transition were fixed to the values obtained from the fit to the J ® 3 2 line. The rotation
diagram has been derived for each velocity component ( Fig. 8, solid and dashed lines).
C6. HYDROGEN SULFIDE: H 2 S
Only one line of hydrogen sulfide at 168.8 GHz, with velocity components at 180 and 275 km s #1 , is detected in the 2 mm survey
(see also MartÐÒn et al. 2005). Physical parameters in Table 2 have been derived by assuming a rotation temperature of 12 K similar to
that derived for many other species (see Table 2).
C7. HYDROCYANIC ACID: HCN
The J ® 2 1 transition of HCN at 177.2 GHz lies outside the covered frequency range, but the 2--1 line of its H 13 CN isotope at
172.6 GHz was detected. A fit to this transition allows us to estimate the abundance of HCN if we assume a rotation temperature of
12 K and the 12 C/ 13 C ratio of 40 derived by Henkel et al. (1993).
C8. OXOMETHYLIUM: HCO +
The only 2 mm transition of the main species of oxomethylium, with energy levels low enough to be observable in the ISM of NGC
253, at a frequency of 178.3 GHz, lies a few GHz above the observed frequency range. Nevertheless, one of its isotopes, H 13 CO + , has
been detected in the J ® 2 1 transition at 173.5 GHz, near the upper end of the spectral scan. Two velocity components can be fitted
to the observed profile. The HCO + column density shown in Table 2 has been calculated from the integrated intensity of the measured
H 13 CO + line with an excitation temperature of 12 K and an isotopic ratio of 12 C/ 13 C # 40 ( Henkel et al. 1993).
The J ® 2 1 transition of HC 18 O + is also identified at 170.3 GHz. The lowívelocity component shows a higher intensity than
expected from the observed H 13 CO + profile as well as a slightly higher velocity. If we compare both profiles, we derive a H 13 CO + /
HC 18 O + ratio of 1:7 # 1:1 and 2:9 # 1:2 for the lowí and highívelocity components, respectively. The value derived for the highívelocity
2 mm LINE SURVEY OF NGC 253 471
No. 2, 2006

component agrees within the errors with the expected value of #3.7 if we assume the isotopic ratios 12 C/ 13 C ® 40 ( Henkel et al.
1993) and 16 O/ 18 O ® 150 ( Harrison et al. 1999).
C9. HYDROISOCYANIC ACID: HNC
Similar to HCN, the J ® 2 1 transition of HNC at 181.3 GHz lies outside the frequency range of the survey. We identify a feature at
174.1 GHz as the J ® 2 1 transition of HN 13 C. The detection is uncertain and thus tentative as the line is strongly blended with the
also tentatively identified 3 1;2 2 1;1 transition of cíC 3 H (see x C18).
C10. ETHYNYL: C 2 H
The J ® 2 1 group of hyperfine transitions of C 2 H, observed at 174.7 GHz near the upper frequency cutoff of the survey, is the
second brightest feature observed in the 2 mm window after CS. Even though the hyperfine structure is unresolved, it is still possible
to differentiate the two velocity components at 160 and 280 km s #1 .
C11. THIOXOETHENYLIDENE: C 2 S
The survey shows four faint lines of thioxoethenylidene (131.6, 140.2, 142.5, and 144.2 GHz). The low intensity of the observed
transitions makes the fitting uncertain (see Table 1). Only the 11 10 --10 9 and 10 11 --9 10 transitions at 131.6 and 140.2 GHz are reliably
detected, while the transitions at 142.5 and 144.2 GHz are tentative.
C12. CARBON OXIDE SULFIDE: OCS
Two transitions of carbon oxide sulfide (133.8 and 145.9 GHz) are detected in the survey. In the rotation diagram in Figure 8 these
observations are complemented by the 3 mm OCS transitions presented by MartÐÒn et al. (2005). The J ® 13 12 transition of OCS at
158.1 GHz might have been detected close to a SO 2 line, but the signalítoínoise ratio is not high enough for a reliable Gaussian fit.
C13. SULFUR DIOXIDE: SO 2
Five transitions of sulfur dioxide (134.0, 135.7, 140.3, 146.6 and 151.4 GHz) have been identified. The upper energy levels of the
transitions range from 12 to 43 K providing a good estimate of the rotation temperature.
The 5 1,5 --4 0,4 transition at 135.7 GHz is blended with a 34 SO line (see x C4). Velocity and width of the line were fixed to perform a
two Gaussian fit for the SO 2 and 34 SO transitions.
The 2 2,0 --2 1,1 line at 151.4 GHz is contaminated by emission of a much fainter cíC 3 H 2 line, which, according to its estimated
intensity (see x C21), would account for #15% of the observed feature. The contribution of the cíC 3 H 2 strongly depends on the
assumed T rot . If we consider the uncertainty in the T rot derived from cíC 3 H 2 and estimate the contribution of this transition to the
observed SO 2 line, the effect on the derived SO 2 parameters would be of #1 K in T rot and #5% in the column density.
The 5 2,4 --5 1,5 and 7 1,7 --6 0,6 transitions are blended with a CH 3 OH group of lines at 165.1 GHz. We have estimated their contribution
to the observed feature by assuming a T rot ® 15 K as derived from the rotation diagram. For the derived excitation conditions of SO 2 ,
its 3 2,2 --3 1,3 line at 158.1 GHz is expected to have an intensity of #5 mK, which is below the noise level at this frequency.
C14. FORMALDEHYDE: H 2 CO
Two ortho (140.8 and 150.5 GHz) and one para (145.6 GHz) J ® 2 1 transitions of formaldehyde are detected. In each of the
profiles, the two velocity components at 180 and 285 km s #1 are clearly separated.
The 2 0,2 --1 0,1 transition at 145.6 GHz and the 2 1,1 --1 1,0 at 150.4 GHz are blended with an HC 3 N and a cíC 3 H 2 line, respectively.
Before performing the Gaussian fits to the H 2 CO features, the estimated contributions from HC 3 N (see x C19) and cíC 3 H 2 (see x C21)
were subtracted from the spectra.
Rotation diagrams are plotted for each velocity component ( Fig. 8) where rotation temperatures of 27 and 34 K are derived for the
180 and 285 km s #1 components, respectively. The 285 km s #1 component of the 2 1,2 --1 1,1 transition appears to have a significantly higher
intensity than expected from the rotation diagram. The large intensity of the 2 1,2 --1 1,1 line of H 2 CO could be due to the contamination by an
unidentified line. If this measurement is not taken into account, the resulting temperature derived for the 285 km s #1 component would be
28 # 2 K, while the estimated column density would be lower by #30%, closer to the parameters derived from the 180 km s #1 component.
C15. THIOFORMALDEHYDE: H 2 CS
Three transitions, all belonging to orthoíthioformaldehyde (135.3, 139.5, and 169.1 GHz), are tentatively detected. The 4 1,4 --3 1,3
transition is blended with the H36# recombination line. A Gaussian profile similar to that of the observed H34# line at 160 GHz (see
x C26) has been subtracted. We assume that the residual emission is due to H 2 CS. We also find the 4 1,3 --3 1,2 line blended with the CS
J ® 3 2 emission from the image band. With the known image sideband rejection at this frequency ( Fig. 3), we can estimate the
contribution of the CS line measured at 146.9 GHz (x C1). The resulting feature after the subtraction of the CS line is fitted by a
Gaussian profile. There is a part of this feature that is not properly fitted by emission from CS and H 2 CS. This may be caused by a
slight difference in the observing position between this spectrum and that containing CS in the signal band, which would cause an
MARTI Ò Ð
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472 Vol. 164

appreciable change in the shape of the subtracted line. Contamination by a third species is unlikely. The 5 1,5 --4 1,4 line at 169.1 GHz is
not blended, but the relatively high noise level at this frequency and the presence of nearby lines makes the fit uncertain.
A rotation temperature of 11 K is determined from the rotation diagram ( Fig. 8) and is consistent with those derived from
other molecules. With this rotation temperature, the 4 0,4 --3 0,3 transition of paraíH 2 CS at 137.3 GHz is expected to have an intensity of
#6 mK. The nondetection of this line could be explained by an excitation temperature of >20 K.
C16. ISOCYANIC ACID: HNCO
The two brightest lines of isocyanic acid in the covered frequency band are clearly detected (131.9 and 153.9 GHz). The 280 km s #1
component of the 7 0,7 --6 0,6 transitions at 153.9 GHz is blended with a group of CH 3 CCH lines. Thus, the contribution of this component
has been calculated by assuming the 180 and 280 km s #1 velocity components have the same excitation temperature. This assumption is
supported by the good fit of the CH 3 CCH line to the residual feature (see x C25). The line observed at 154.4 GHz is tentatively identified as
the 280 km s #1 velocity component of the HNCO 7 1.6 --6 1,5 transition. The 7 1,7 --6 1,6 transition at 153.3 GHz, which should have a similar
intensity, is not detected. Due to the uncertainty of this identification, this line has not been used in the analysis of the HNCO excitation
conditions. Including this transition in the rotation diagram (Fig. 8) would result in a derived T rot # 29 K, slightly higher than that derived
for the 180 km s #1 velocity component, and a column density #16% lower than that in Table 2.
C17. PORTMANTEAU CARBON DIOXIDE: HOCO +
We tentatively identified two transition (149.6 and 171.0 GHz) of portmanteau carbon dioxide, a quasilinear molecular ion and
nearly symmetric top. These are the brightest observable transitions in the 2 mm range. Only the highívelocity component of both
lines are detected. The large intensity observed in the 8 1,8 --7 1,7 transitions, expected to be similar or lower than the 7 0,7 --6 0,6 , is mainly
due to the uncertainties of the baseline substraction. A rotational temperature of 12 K have been assumed to estimate an average
column density with both detected transitions. Further observations of other HOCO + transitions are needed to confirm the detection
and constrain the excitation temperature of this molecule.
C18. CYCLOPROPENYLIDYNE: cíC 3 H
The 3 1,2 --2 1,1 group of transitions at 174 GHz is the only feature of cíC 3 H identified in the survey. Contamination by the feature by
the J ® 2 1 transition of HN 13 Cmakes the fitting uncertain. Thus, the velocity component were fixed to V LSR ® 180 and 280 km s #1
as derived for most species and the width of the lower velocity component to #v 1 = 2 ® 41 km s #1 similar to that from the fitting to the
highívelocity one. The 4 1,4 --3 1,3 transitions at 172.6 GHz appear to be blended with a H 13 CN line. At 132.9 GHz, the 3 1,3 --2 1,2 group
of transitions, which is expected to be similar in intensity to those at 174 GHz, is likely present blended to a transition of CH 2 NH and
an also likely transition of CH 3 OH, but the noise level at these frequencies do not allow us to separate these lines. Given the considered
source size of 20 00 the 133 GHz line would be 20% fainter than that at 174 GHz, so beam dilution cannot account for the nondetection
of these lines. Therefore, its detection, similar to the HN 13 C line to which it is blended (x C9) is highly tentative.
C19. CYANOACETYLENE: HC 3 N
Four out of the five transitions of cyanoacetylene detected in the observed frequency range (136.5, 145.6, 154.7, 163.8, and 172.8 GHz)
appear not to be blended with any other molecular line. Two velocity components have been fitted to each line. The J ® 16 15 transition
at 145.6 GHz is blended with an H 2 CO line. Since we have observed more HC 3 N than H 2 CO lines to accurately determine the physical
parameters, we have estimated the relative contribution of each of the velocity components of HC 3 N. For the HC 3 N J ® 18 17 line,
where the velocity components are not clearly differentiated, the velocity of each component has been fixed. The J ® 19 18 transition
seems to lie above the T rot # 20 K derived from the other observed transitions. This may indicate that the lower energy transitions would
not be optically thin as blending with any other molecular line seems unlikely. Multiline observations by Mauersberger et al. (1990; toward
a position 9 00 north of our position) also noticed an excess of the higher transition lines of HC 3 N. Their comparison with model
computations show that two molecular components are required to explain the observed line intensities, one with n(H 2 ) # 10 4 cm #3 and
T kin # 60 K and a second with n(H 2 ) # 5 ; 10 5 cm #3 and T kin # 150 K. The line intensity of the J ® 17 16 transition measured in this
work is about half that measured by Mauersberger et al. (1990), which indicates the strong effect of a pointing offset in the total integrated
intensity. However, as far as the physical parameters are concerned, the use of the J ® 19 18 line, although causing a change in the
determined rotation temperature of #50%, affects the total column density by less than a factor of 2.
C20. METHANIMINE: CH 2 NH
One transition of methanimine is clearly identified in the line scan. The observed profile of the 2 1,1 --1 1,0 transition at 133 GHz is
heavily affected by nonlinear baselines, and the fit shown in Table 1 is uncertain. The 1 1,0 --1 0,1 transition at 167 GHz, expected to have
an intensity twice as large as that at 133 GHz, is not reliably detected due to the lack of sensitivity at this frequency. Although the
detection seems to be clear, further observations of transitions in other frequency bands are needed to confirm this detection.
C21. CYCLOPROPENYLIDENE: cíC 3 H 2
Three transitions of cyclopropenylidene (155.5 and 150.8 GHz) are clearly detected. However, only one transition, the 3 2,2 --2 1,1
line at 155.5 GHz, appears not to be blended with any other line.
2 mm LINE SURVEY OF NGC 253 473
No. 2, 2006

The 4 0,4 --3 1,3 and 4 1,4 --3 0,3 lines, both at 150.8 GHz, appear to be blended. A double Gaussian profile was fitted by assuming equal
width and velocity and a line intensity ratio determined by their spectroscopical parameters, assuming optically thin emission.
The intensities of the other blended transitions (145.1 and 150.4 GHz) were estimated for a rotation temperature of 9 K as derived
from the rotation diagram ( Fig. 8). Given the small dynamic range in energies of the unblended transitions used in the rotation
diagram, the estimated T rot is affected by a large uncertainty (see Table 2) and, therefore, the estimates for the blended transitions are
also quite uncertain. The 5 1,4 --5 0,5 transition at 151.3 GHz with an estimated intensity of #1 mK, well below the noise level of the
spectra, has been calculated with the aim of estimating its contribution to the faint SO 2 transition at that frequency (x C13).
C22. CYANAMIDE: NH 2 CN
Only one line of cyanamide is clearly detected in the survey, the 8 1,7 --7 1,6 transition at 161.0 GHz, which appears to be partially
blended with image sideband emission from H 2 S. The emission of H 2 S in the signal band has been corrected by the image sideband
rejection and subtracted from the spectrum before fitting the NH 2 CN line. The other observed lines are close to the noise level, and
therefore all fitted parameters are strongly affected by the baseline. The 8 1,8 --7 1,7 transition at 158.8 GHz, seems to be blended with a
SO (x C4) but the noise level does not permit a fit. The high rotation temperature of 63 K derived from the rotation diagram ( Fig. 8)
depends mainly on the fainter and therefore less reliably fitted transitions. Amuch lower rotation temperature would lower the column
density by a factor of about 2.
C23. METHYL CYANIDE: CH 3 CN
Two spectral features are identified as methyl cyanide transitions (147.2 and 165.6 GHz). Each of the J J transitions of this
symmetric top consist of a number of overlapped K components (K ® 0 : : : J 1). Single Gaussian profiles were fitted to the obí
served lines. The contribution of the K ® 0 and 1 ladders, for the low rotation temperatures derived from other high dipole moment
molecules, account for #98% of the line profile. Therefore, the K > 1 contribution to the total line intensity has not been taken into
account. We derive a T rot ® 10 K from the rotation diagram in Fig. 8, which has been obtained as explained in x B2.
C24. METHANOL: CH 3 OH
A total of nine transitions or groups of transitions of methanol is detected in the line survey at 143.9, 145.1, 146.4, 157, 165, and 170 GHz.
The transitions at 145.1 GHz are blended with the C 3 H 2 3 1,2 --2 2,1 line, whose contribution has been estimated (see x C21) and
subtracted from the observed feature. In the same way, the lines at 165.0 GHz are slightly blended with SO 2 transitions. The emission
of these SO 2 lines, subtracted from the spectra (see x C13), mainly affects the fit of the 4 1,3 --4 0,4 transition. The 6
#1 --5 0 transition of
CH 3 OH expected to be detected at 132.8 GHz is heavily affected by baseline instabilities. The rotation diagram derived from the
observed transitions has been plotted using both the singleí and multitransition lines as described in x B2. The derived rotation
temperature is 12 K.
C25. METHYL ACETYLENE: CH 3 CCH
Three methyl acetylene features are observed (136.7, 153.8, and 170.9 GHz). As in the case of CH 3 CN, each of these corresponds to
a #J ® 1 transition consisting of a number of unresolved K components (with K ® 0 : : : J 1). Single Gaussian profiles have been
fitted to the observed profiles.
The best fit to the rotation diagram as described in x B2 results in a rotation temperature of #62 K. Given the high rotation
temperature derived, only contributions of transitions with K ® 0 : : : 4, which have upper level energies <160 K, have been taken
into account in the fit. For a temperature of 60 K, transitions with K ® 4 contribute #5% to the observed profiles, while higher K lines
represent less than 1% of the integrated intensities. Neglecting the contribution of these higher energy transitions (i.e., assuming the
Fig. 9.---Recombination line fluxes of measured H# lines in the range from 1 to 160 GHz as a function of frequency. A bestífit to the line fluxes in the form
S L / # 2 is shown as fitted by Puxley et al. (1997).
MARTI Ò Ð
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474 Vol. 164

whole component is dominated by the K ® 0 transitions) would result in an overestimation of the column densities by up to a factor
of 4.
C26. HYDROGEN RECOMBINATION LINES
Four H# recombination lines are present within the frequency range covered by our survey. Only one of them, H34#, is not blended
and can be fitted. The H33# line lies close to the 2--1 line of C 2 H at the upper end of the frequency range where the rms is similar to the
expected intensity of the line. The H35# and H36# lines are blended with CS 3--2 and H 2 CS 4 1,4 --3 1,3 lines, respectively. If the
relation of the line flux S L / # 2
L applies at these frequencies ( Puxley et al. 1997), the lines should have flux densities 20% and 30%
below that measured for the H34# line. We have considered them to have a similar line profile to that of the measured H34# line given
the uncertainty in the fitted intensity.
To convert the measured H34# line temperature into flux densities, we use the conversion factor for the 30 m telescope S/T MB ®
4:95 Jy K #1 . Following Puxley et al. (1997), a source to beam coupling factor (# 2
b ” # 2
s )/# 2
b has been applied for a source size # s ® 7 00 .
We calculate an integrated line flux of (1:08 # 0:08) ; 10 #19 Wm #2 . As seen in Figure 9, this value closely follows the best fit to the
relation between the integrated line flux versus frequency derived by Puxley et al. (1997). In this figure the flux of the H34# line is
plotted together with the recombination lines observed in millimeter and centimeter wavelengths ( Puxley et al. 1997 and references
therein).
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