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# Detection of complex organic molecules in a prestellar core: a new challenge for astrochemical models

A. Bacmann, V. Taquet, A. Faure, C. Kahane, C. Ceccarelli

► To cite this version:

A. Bacmann, V. Taquet, A. Faure, C. Kahane, C. Ceccarelli. Detection of complex organic molecules in a prestellar core: a new challenge for astrochemical models. Astronomy and Astrophysics - A&A, 2012, 541, 10.1051/0004-6361/201219207. insu-03612434

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Distributed under a Creative Commons Attribution 4.0 International License
LETTER TO THE EDITOR

# Detection of complex organic molecules in a prestellar core: a new challenge for astrochemical models*,**

A. Bacmann¹, V. Taquet¹, A. Faure¹, C. Kahane¹, and C. Ceccarelli¹

UJF-Grenoble 1 / CNRS-INSU, Institut de Planétologie et d'Astrophysique de Grenoble (IPAG) UMR 5274, 38041 Grenoble, France
e-mail: abacmann@ujf-grenoble.fr

Received 12 March 2012 / Accepted 29 March 2012

## ABSTRACT

**Context.** Complex organic molecules (COMs) have long been detected in the interstellar medium, especially in hot cores and in the hot corinos of low-mass protostars. Their formation routes however remain uncertain. Both warm gas-phase reactions and warm grain-surface reactions have been invoked to account for their presence in low-mass protostars. In this latter scheme, COMs result from radical-radical reactions on the grains as radicals become mobile when the nascent protostar warms up its surroundings and the resulting molecules are subsequently desorbed into the gas phase at higher temperatures.

**Aims.** Prestellar cores are the direct precursors of low-mass protostars and offer a unique opportunity to study the formation of COMs before the warm-up phase. Their very low temperatures (≤10 K) and the absence of any heating source or outflow exclude any efficient warm gas phase or warm dust chemistry, so that the presence of COMs in prestellar cores would have to originate from non-thermal chemical processes.

**Methods.** We used the IRAM 30 m telescope to look for four O-bearing COMs (acetaldehyde CH<sub>3</sub>CHO, dimethyl ether CH<sub>3</sub>OCH<sub>3</sub>, methyl formate CH<sub>3</sub>OCHO, and ketene CH<sub>2</sub>CO) in the prestellar core L1689B.

**Results.** We report the unambiguous detection of all four molecules in the cold gas phase of L1689B. These detections support the role played by non-thermal (possibly photolytic) processes in COM formation and desorption, though the presence of dimethyl ether is so far unexplained by current grain formation scenarios. The data show univocally that COM synthesis has already started at the prestellar stage and suggests at least part of the COMs detected in hot corinos have a prestellar origin.

**Key words.** astrochemistry – molecular processes – ISM: molecules – stars: formation

## 1. Introduction

Complex organic molecules (COMs) have long been detected in the hot cores of massive star-forming regions such as Sgr B2 (Cummins et al. 1986) or OMC-1 (Blake et al. 1987). The warm central regions of low-mass protostars have been shown to harbour the same organic molecules (van Dishoeck et al. 1995; Cazaux et al. 2003), which led them to be called “hot corinos” (Bottinelli et al. 2004). In these regions, O- and N- bearing species like methyl formate (CH<sub>3</sub>OCHO), dimethyl ether (CH<sub>3</sub>OCH<sub>3</sub>) or acetonitrile (CH<sub>3</sub>CN) are abundant, so that these molecules have been designated as “hot core tracers”. The mechanisms leading to COM formation in star-forming regions are however still debated. COMs were first thought to be the result of warm gas-phase chemistry between ice mantle constituents, sometimes called primary species or parent molecules, which are evaporated in the gas phase in the vicinity of the nascent protostar (Charnley et al. 1992). In this scheme, primary molecules are formed on the grain ice mantles via successive hydrogenation of atoms or simple molecules like CO. When they are injected into the gas phase via thermal desorption, they can react and form secondary species like, e.g., CH<sub>3</sub>OCHO. The gas-phase scenario has been however challenged by experimental results showing that gas-phase formation pathways were much less efficient

than originally thought (Horn et al. 2004, for methyl formate). Alternative scenarios have been proposed in which radicals are formed by ultraviolet (UV) irradiation of grain mantle molecules (e.g., Garrod & Herbst 2006) or trapped within the bulk of ice mantles of dust grains (Taquet et al. 2012). Following Garrod & Herbst (2006), the radicals can react on the grain mantles once the protostar starts heating its immediate surroundings and the grains reach a temperature of around 30–40 K, and form complex species on the grain, before they are being desorbed and released into the gas phase. Laboratory experiments however have shown the possibility to synthesize COMs such as methyl formate by irradiation (with UV photons or energetic particles) of methanol ices at temperatures as low as 10 K (Gerakines et al. 1996; Bennett & Kaiser 2007).

In contrast to warm star-forming regions such as hot cores or hot corinos, little research has been carried out to investigate the contents of cold gas. Studies of the dark cloud TMC-1 have revealed the presence of acetaldehyde (Matthews et al. 1985), ketene (Irvine et al. 1989), and methanol (Friberg et al. 1988). Evidence of COMs in prestellar cores is similarly scarce: Matthews et al. (1985) reported the detection of acetaldehyde in L183 (L134N), but searches for dimethyl ether (Friberg et al. 1988), ketene (Irvine et al. 1989) and methyl formate (Requena-Torres et al. 2007) only yielded upper limits. Despite their very low temperatures (~10 K), prestellar cores are however an important piece of the puzzle, as they represent the stage immediately before the gravitational collapse leading to protostar formation. They thus offer the opportunity to study the initial conditions of the star formation process. The relative simplicity

<footnote>* Based on observations carried out with the IRAM 30 m Telescope. IRAM is supported by INSU/CNRS (France), MPG (Germany), and IGN (Spain).</footnote>

<footnote>** Appendices are available in electronic form at http://www.aanda.org</footnote>
<figure><img src="image_3.png" /><figcaption>**Fig. 1.** CH<sub>3</sub>OCH<sub>3</sub> spectrum at 99.325 GHz in the prestellar core L1689B. Three lines corresponding to the AA, EE, EA, and AE symmetry substates (AE and EA being at the same frequency) are clearly detected at the >7σ level. The vertical dotted lines indicate the rest frequencies measured in the laboratory and listed in the CDMS database (Müller et al. 2001).</figcaption></figure>

of their structure (shallow temperature gradients, absence of in-
ner heating source, absence of shocks or outflows, environment
shielded from external radiation fields) makes them ideal targets
to study chemical processes in cold gas at the very early stages
of star formation.

We present here the first detection of the COMs dimethyl
ether and methyl formate in a prestellar core, as well as the de-
tection of the two other COMs acetaldehyde and ketene, in the
gas phase at a temperature of ~10 K.

**2. Observations**

The target of our observation is the prestellar core L1689B.
This source is a typical core characterised by low dust tem-
peratures around 11–12 K (Kirk et al. 2007), high central den-
sities of 10<sup>5</sup>–10<sup>6</sup> cm<sup>-2</sup>, and a high central column density of
~10<sup>23</sup> cm<sup>-2</sup> (Bacmann et al. 2000; Dapp & Basu 2009). It was
selected based on the assumption that the gas-phase abundances
of C-bearing molecules such as formaldehyde (Bacmann et al.
2003) or methanol (Bacmann et al., in prep.) were relatively
large with respect to other cores. Therefore, it can be hypoth-
esised that COMs are more likely to be detected in this object
than in more depleted ones.

Observations were carried out in March 2011 at the IRAM
30 m telescope. We used the E090 receiver tuned to a frequency
of 99.325 GHz (close to low energy transitions of CH3OCH3)
to which we connected both the high resolution autocorrelator
VESPA and the low resolution Fourier transform spectrometer
(FTS). The spectral resolution and bandwidth of VESPA were
20 kHz (corresponding to a velocity resolution of ~0.06 km s-1)
and 100 MHz, respectively. The FTS covered the whole 4 GHz
receiver band (from 97.6 to 101.6 GHz in this case) with a spec-
tral resolution of 192 kHz, corresponding to a velocity resolu-
tion of ~0.6 km s-1. The observations were carried out in the
frequency switching mode with a frequency switch of 7.5 MHz
to reduce standing waves. Pointing was checked every two hours
on a nearby quasar and was found to be accurate to within 1".
The weather conditions were very good and system tempera-
tures were typically 100–120 K despite the low elevation of the
source. We integrated on the dust peak of the core L1689B, at
α2000 = 16h34m48s30, δ2000 = -24°38′04.0″. The size of the
telescope beam at the observing frequency was ~25", and the
forward and beam efficiencies were ηfwd = 0.95 and ηbeam =
0.80, respectively.

The data were reduced using the IRAM CLASS/GILDAS
package<fnref n="1" />. The scans were coadded, deconvolved from the fre-
quency switching procedure, and a low order (typically 3) poly-
nomial was then fitted over line-free regions to correct for base-
line oscillations. The spectra were converted from T<sub>a</sub>\* to the T<sub>mb</sub>
scale, using the values of the efficiencies given above.

**3. Results**

The high spectral resolution spectrum obtained close to
99.325 GHz is shown in Fig. 1. The rms noise in the spec-
tra is ~7 mK. The frequencies and spectroscopic parameters of
the detected species are listed in Table A.1. Three spectral lines
are clearly detected (to the 7-12σ level), corresponding to the
414-303 transitions of the four symmetry substates (AA-, EE-,
AE-, and EA-) of CH3OCH3. The EA- and AE- transitions are
only 2 kHz apart and therefore are indistinguishable in our spec-
trum. The lines are very narrow, ≈0.4 km s⁻<fnref n="1" /> wide, as expected
for lines originating from cold gas at around 10 K. These spec-
tra show the unambiguous presence in the cold gas of dimethyl
ether, which is detected for the first time in a cold dark core.

The low-resolution spectrum covers the 97.6 GHz to 101.6 GHz region. Within this bandwidth, we also detect six lines of CH<sub>3</sub>OCHO (three lines of the A- and three lines of the E- substates), which had so far never been detected in a prestellar core. In addition, two lines of CH<sub>2</sub>CO (one ortho line and one para line) and two lines of CH<sub>3</sub>CHO (one of the A- and one of the E-substates) are also detected. Figures 2–4 present the spectra of CH<sub>3</sub>OCHO, CH<sub>3</sub>CHO, and CH<sub>2</sub>CO, respectively. Again, the detection of more than one line of each species at the measured catalog frequencies (see figures) is clear evidence of these three species in the gas phase of L1689B.

At these low temperatures, the different substates have to be
considered separately. Since we only have one line for each sub-
state (except for methyl formate), the excitation conditions are
poorly constrained and estimates of the column densities can
only be given by assuming an excitation temperature. In the case
of methyl formate, all detected lines have similar upper level en-
ergies and Einstein coefficients, so that having more than one
line per substate does not bring supplementary constraints for
the excitation. As these data are a subset of a larger ongoing
observing project targetting several lines for each molecule and
substate, we postpone the detailed analysis<fnref n="2" /> of the column den-
sities and abundances to a future paper and give here only pre-
liminary estimates of the column densities. At densities around
10⁵–10⁶ cm⁻³ (similar to the critical density of 10⁶ cm⁻³), and
gas temperatures of 7–10 K (Redman et al. 2004), it is reason-
able to assume excitation temperatures of around 5 K. Because
this value is uncertain, the column densities are also given for
more extreme values of the excitation temperature, 4 K and 8 K.

The total column densities of the species were derived from
the integrated line intensities assuming local thermodynamic
equilibrium (see Appendix B for more details). The derived col-
umn densities are given in Table 1. The difference between the
column density values at T<sub>ex</sub> = 4 K and T<sub>ex</sub> = 8 K varies be-
tween a factor of two for CH<sub>3</sub>OCH<sub>3</sub> and an order of magni-
tude for CH<sub>3</sub>OCHO. It is reasonable to expect the same exci-
tation temperature for the A- and E- substates of CH<sub>3</sub>CHO and
CH<sub>3</sub>OCHO, as these have similar upper level energies, Einstein
coefficients, frequencies, and upper level degeneracies. We find

<footnote n="1">¹ http://www.iram.fr/IRAMFR/GILDAS</footnote>

<footnote n="2">² Note that collisional rates for methyl formate have been obtained by Faure et al. (2011).</footnote>
<figure><img src="image_4.png" /></figure>

**Fig. 2.** CH<sub>3</sub>OCHO spectra of the 8<sub>17</sub>−7<sub>11</sub> transitions at 100.48 GHz (A and E substates, *bottom panel*), of the 9<sub>09</sub>−8<sub>08</sub> transitions at 100.68 GHz (A and E substates, *middle panel*), and of the 9<sub>19</sub>−8<sub>18</sub> transitions at 100.08 GHz (A and E substates, *top panel*) in L1689B. The vertical dotted lines indicate the rest frequencies measured in the laboratory and listed in the JPL database (Pickett et al. 1998). The CH<sub>3</sub>OCHO spectrum in the top panel also shows a HC<sub>3</sub>N line. The negative spectral feature between both CH<sub>3</sub>OCHO lines results from the frequency switching procedure and the folding of the ortho-CH<sub>2</sub>CO line shown in Fig. 4.

<figure><img src="image_5.png" /></figure>

**Fig. 3.** CH<sub>3</sub>CHO spectra of the 5<sub>14</sub>−4<sub>13</sub> transitions at 98.9009 GHz (A substate, *left*) and at 98.8633 GHz (E substate, *right*) in L1689B. The vertical dotted lines indicate the rest frequencies measured in the laboratory and listed in the JPL database (Pickett et al. 1998).

that both A- and E- substates are equally abundant in those two molecules. For CH<sub>3</sub>OCH<sub>3</sub>, we find the abundance ratio of the EE- to AA- substates to be equal to the statistical ratio of two, although no conclusion can be drawn about the AE- and

<figure><img src="image_6.png" /></figure>

**Fig. 4.** CH<sub>2</sub>CO spectra of the 5<sub>05</sub>−4<sub>04</sub> (para) transition at 101.0367 GHz (*left*) and of the 5<sub>15</sub>−4<sub>14</sub> (ortho) transition (*right*) at 100.0945 GHz in L1689B. The vertical dotted lines indicate the laboratory rest frequencies listed in the CDMS database (Müller et al. 2001).

EA- substates, as these are too close in values of frequency to be spectrally resolved, so we simply assumed that the ratio of their intensities was equal to the ratio of their upper level degeneracies (which leads to equal column densities in AE- and EA-, each equal to half that of the AA- substate – see note to Table 1). For CH<sub>2</sub>CO, we find an ortho-to-para ratio of ~3.3, which is similar to the value of 3.5 found in TMC-1 (Ohishi et al. 1991), consistent with the statistical ratio of 3. Assuming a H<sub>2</sub> column density of 10<sup>23</sup> cm<sup>-2</sup> (Dapp & Basu 2009), we find abundances of a few 10<sup>-10</sup> for the four species. This is a factor of ~100 lower than in hot corinos (e.g. in NGC1333-IRAS4A – Bottinelli et al. 2004), similar to the formaldehyde abundance ratio between prestellar cores (Bacmann et al. 2003) and hot corinos (Bottinelli et al. 2004).

## 4. Discussion and conclusions

The detection of CH<sub>3</sub>OCH<sub>3</sub>, CH<sub>3</sub>OCHO, CH<sub>2</sub>CO, and CH<sub>3</sub>CHO in the gas phase of a prestellar core at temperatures of around 10 K sheds new light on COM formation in the interstellar medium. Grain-surface chemistry models such as those of Garrod & Herbst (2006) require temperatures of around 30 K to ensure that the radicals present in the ice mantle have sufficient mobility to recombine and form COMs. This scenario cannot, however, explain the presence of COMs in the gas phase of this prestellar core. Pure gas-phase scenarios assuming that COMs result from warm (~100 K) gas-phase reactions between grain mantle constituents after their evaporation from the grains (Charnley et al. 1992) are also excluded by the low temperature in prestellar cores.

According to the hypothesis of a grain mantle formation scenario, the COMs that we detect in the gas phase have been desorbed from the ices by non-thermal processes such as secondary UV photon (originating from cosmic rays impacting on H<sub>2</sub> molecules) heating, or direct cosmic-ray impacts. It is therefore possible that radical mobility on the grains is also provided by the energy brought to the grains by UV photons. Similar mechanisms were invoked by Öberg et al. (2010) who detected methyl formate in the protostar B1-b, possibly desorbed by UV originating from the protostar through an outflow cavity. In the case of L1689B, the gas is significantly colder than in B1-b, and the core is well-shielded from external UV, as there is no nearby protostar and no outflow cavity so that the core is not exposed to a direct UV field. In Galactic center clouds, the detection of COMs by Requena-Torres et al. (2006) in the gas phase has been interpreted as the result of the ejection of ice mantle molecules by either shocks or supersonic turbulence. This scheme does not apply to the prestellar core L1689B, in which there are no shocks and the levels of turbulence are very low, as demonstrated by
<table><caption>**Table 1.** Line fluxes<sup>a</sup> and column densities.</caption><thead><tr><th>Species</th><th>Transition</th><th><i>T</i><sub>mb</sub><br>(mK)</th><th>rms<br>(mK)</th><th>Integrated area<br>(K km s<sup>-1</sup>)</th><th>Column density<sup>c</sup><br>(cm<sup>-2</sup>)<br><i>T</i><sub>ex</sub> = 5 K</th><th>Column density<sup>c</sup><br>(cm<sup>-2</sup>)<br><i>T</i><sub>ex</sub> = 4 K</th><th>Column density<sup>c</sup><br>(cm<sup>-2</sup>)<br><i>T</i><sub>ex</sub> = 8 K</th></tr></thead><tbody><tr><td>AA-CH<sub>3</sub>OCH<sub>3</sub></td><td>4<sub>14</sub>-3<sub>03</sub></td><td>55</td><td>7</td><td>0.026 ± 0.005</td><td>(3.32 ± 0.74) 10<sup>12</sup></td><td>(5.08 ± 1.12) 10<sup>12</sup></td><td>(2.49 ± 0.55) 10<sup>12</sup></td></tr><tr><td>EE-CH<sub>3</sub>OCH<sub>3</sub></td><td>4<sub>14</sub>-3<sub>03</sub></td><td>90</td><td>7</td><td>0.037 ± 0.005</td><td>(6.81 ± 0.92) 10<sup>12</sup></td><td>(10.4 ± 1.40) 10<sup>12</sup></td><td>(5.11 ± 0.69) 10<sup>12</sup></td></tr><tr><td>EA-CH<sub>3</sub>OCH<sub>3</sub></td><td>4<sub>14</sub> - 3<sub>03</sub></td><td>36<sup>b</sup></td><td>7</td><td>0.014 ± 0.003</td><td>(1.65 ± 0.55) 10<sup>12</sup></td><td>(2.53 ± 0.84) 10<sup>12</sup></td><td>(1.24 ± 0.41) 10<sup>12</sup></td></tr><tr><td>AE-CH<sub>3</sub>OCH<sub>3</sub></td><td>4<sub>14</sub>-3<sub>03</sub></td><td>22<sup>b</sup></td><td>7</td><td>0.009 ± 0.003</td><td>(1.67 ± 0.37) 10<sup>12</sup></td><td>(2.55 ± 0.56) 10<sup>12</sup></td><td>(1.25 ± 0.28) 10<sup>12</sup></td></tr><tr><td>E-CH<sub>3</sub>CHO</td><td>5<sub>14</sub>-4<sub>13</sub></td><td>125</td><td>3</td><td>0.149 ± 0.015</td><td>(9.12 ± 0.92) 10<sup>12</sup></td><td>(2.67 ± 0.27) 10<sup>13</sup></td><td>(2.12 ± 0.21) 10<sup>12</sup></td></tr><tr><td>A-CH<sub>3</sub>CHO</td><td>5<sub>14</sub>-4<sub>13</sub></td><td>110</td><td>4</td><td>0.137 ± 0.014</td><td>(8.26 ± 0.84) 10<sup>12</sup></td><td>(1.73 ± 0.18) 10<sup>13</sup></td><td>(3.86 ± 0.39) 10<sup>12</sup></td></tr><tr><td>A-CH<sub>3</sub>OCHO</td><td>8<sub>17</sub>-7<sub>16</sub></td><td>20</td><td>3</td><td>0.025 ± 0.005</td><td>(2.01 ± 0.40) 10<sup>13</sup></td><td>(5.73 ± 1.15) 10<sup>13</sup></td><td>(5.91 ± 1.18) 10<sup>12</sup></td></tr><tr><td>A-CH<sub>3</sub>OCHO</td><td>9<sub>09</sub>-8<sub>08</sub></td><td>12</td><td>2</td><td>0.016 ± 0.005</td><td>(1.73 ± 0.54) 10<sup>13</sup></td><td>(5.50 ± 1.71) 10<sup>13</sup></td><td>(4.33 ± 1.35) 10<sup>12</sup></td></tr><tr><td>E-CH<sub>3</sub>OCHO</td><td>8<sub>17</sub>-7<sub>16</sub></td><td>19</td><td>3</td><td>0.024 ± 0.005</td><td>(1.90 ± 0.40) 10<sup>13</sup></td><td>(5.44 ± 1.15) 10<sup>13</sup></td><td>(5.59 ± 1.18) 10<sup>12</sup></td></tr><tr><td>E-CH<sub>3</sub>OCHO</td><td>9<sub>09</sub>-8<sub>08</sub></td><td>12</td><td>2</td><td>0.016 ± 0.003</td><td>(1.72 ± 0.32) 10<sup>13</sup></td><td>(5.47 ± 1.03) 10<sup>13</sup></td><td>(4.32 ± 0.81) 10<sup>12</sup></td></tr><tr><td>o-CH<sub>2</sub>CO</td><td>5<sub>15</sub>-4<sub>14</sub></td><td>140</td><td>3</td><td>0.171 ± 0.018</td><td>(1.50 ± 0.16) 10<sup>13</sup></td><td>(3.09 ± 0.33) 10<sup>13</sup></td><td>(6.75 ± 0.71) 10<sup>12</sup></td></tr><tr><td>p-CH<sub>2</sub>CO</td><td>5<sub>05</sub>-4<sub>04</sub></td><td>82</td><td>2</td><td>0.096 ± 0.010</td><td>(4.54 ± 0.47) 10<sup>12</sup></td><td>(9.64 ± 1.00) 10<sup>12</sup></td><td>(1.95 ± 0.20) 10<sup>12</sup></td></tr></tbody></table>

**Notes.** (a) The fluxes were derived by fitting Gaussians to the lines using the dedicated CLASS routine. The uncertainty was calculated as $\sqrt{\sigma_{stat} + \sigma_{cal}}$, where $\sigma_{stat}$ is the statistical uncertainty in the Gaussian fit (given by the CLASS routine) and $\sigma_{cal}$ is the calibration uncertainty taken as 10%. (b) Both EA- and AE- CH<sub>3</sub>OCH<sub>3</sub> lines are so close in frequency that they are indistinguishable in our spectra. In the table, we assume that the ratio of their intensities was equal to the ratio of their upper level degeneracy because the transitions have nearly equal frequencies, Einstein coefficients, and equal level energies. This hypothesis is supported by it also applying to the EE- and AA- transitions (the ratio of their respective intensities is consistent with the ratio of their respective upper level degeneracies). (c) The column densities in Cols. (6)–(8) are those of the corresponding substates listed in Col. (1), as derived from the line fluxes given in Col. (5).

the narrow linewidths. If COMs form on grains, the most likely mechanisms to bring radical mobility and lead to molecular desorption are thus secondary UVs, cosmic rays, or the release of chemical energy due to exothermic reactions on the grains (e.g. chemical explosions, Garrod et al. 2007).

The simultaneous detection of CH<sub>3</sub>OCH<sub>3</sub> and CH<sub>3</sub>OCHO in very cold gas is particularly interesting. According to Öberg et al. (2011), COMs could form sequentially and CH<sub>3</sub>OCHO and CH<sub>3</sub>CHO are formed by UV irradiation of cold CO-rich ices and CH<sub>3</sub>OCH<sub>3</sub> forms in warmer (CH<sub>3</sub>OH-rich) ice after CO has evaporated (see also Öberg et al. 2009). Our observations do not support this scenario, and show that either CH<sub>3</sub>OCH<sub>3</sub> forms in the same ice as CH<sub>3</sub>OCHO (since prestellar core have rather uniform temperature and do not undergo temperature changes similar to those of protostars) or that the formation mechanisms of both types of molecules are different with, e.g., CH<sub>3</sub>OCH<sub>3</sub> forming in the gas phase, as suggested by Peeters et al. (2006), though this idea is disputed by Requena-Torres et al. (2006).

The main conclusion of this paper is the detection of methyl formate and dimethyl ether for the first time in a very cold (~10 K) dark core, here the prestellar core L1689B. Acetaldehyde and ketene, which were previously detected in TMC-1, are also detected. The detection of these species is also further evidence that they do not specifically trace hot core or hot corino regions, but that they are also present in the cold and dense ISM. Since prestellar cores are thought to be the precursors of low-mass protostars, the presence of COMs at the prestellar stage means that we have to re-investigate COM formation scenario in protostars, especially in the case of dimethyl ether, which is not expected to form in similar ices as methyl formate. Following our results, it is likely that at least part of the COM contents in young protostars originates from the prestellar phase.

*Acknowledgements.* This work has been supported by l'Agence Nationale pour la Recherche (ANR), France (contract ANR-08-BLAN-022).

## References

Bacmann, A., André, P., Puget, J.-L., et al. 2000, A&A, 361, 555

Bacmann, A., Lefloch, B., Ceccarelli, C., et al. 2003, ApJ, 585, L55

Bennett, C. J., & Kaiser, R. I. 2007, ApJ, 661, 899

Blake, G. A., Sutton, E. C., Masson, C. R., & Phillips, T. G. 1987, ApJ, 315, 621

Bottinelli, S., Ceccarelli, C., Lefloch, B., et al. 2004, ApJ, 615, 354

Cazaux, S., Tielens, A. G. G. M., Ceccarelli, C., et al. 2003, ApJ, 593, L51

Charnley, S. B., Tielens, A. G. G. M., & Millar, T. J. 1992, ApJ, 399, L71

Cummins, S. E., Linke, R. A., & Thaddeus, P. 1986, ApJS, 60, 819

Dapp, W. B., & Basu, S. 2009, MNRAS, 395, 1092

Faure, A., Szalewicz, K., & Wiesenfeld, L. 2011, J. Chem. Phys., 135, 024301

Friberg, P., Hjalmarson, A., Madden, S. C., & Irvine, W. M. 1988, A&A, 195, 281

Garrod, R. T., & Herbst, E. 2006, A&A, 457, 927

Garrod, R. T., Wakelam, V., & Herbst, E. 2007, A&A, 467, 1103

Gerakines, P. A., Schutte, W. A., & Ehrenfreund, P. 1996, A&A, 312, 289

Horn, A., Møllendal, H., Sekiguchi, O., et al. 2004, ApJ, 611, 605

Irvine, W. M., Friberg, P., Kaifu, N., et al. 1989, ApJ, 342, 871

Kirk, J. M., Ward-Thompson, D., & André, P. 2007, MNRAS, 375, 843

Matthews, H. E., Friberg, P., & Irvine, W. M. 1985, ApJ, 290, 609

Müller, H. S. P., Thorwirth, S., Roth, D. A., & Winnewisser, G. 2001, A&A, 370, L49

Öberg, K. I., Garrod, R. T., van Dishoeck, E. F., & Linnartz, H. 2009, A&A, 504, 891

Öberg, K. I., Bottinelli, S., Jørgensen, J. K., & van Dishoeck, E. F. 2010, ApJ, 716, 825

Öberg, K. I., van der Marel, N., Kristensen, L. E., & van Dishoeck, E. F. 2011, ApJ, 740, 14

Ohishi, M., Kawaguchi, K., Kaifu, N., et al. 1991, in Atoms, Ions and Molecules: New Results in Spectral Line Astrophysics, ed. A. D. Haschick, & P. T. P. Ho, ASP Conf. Ser., 16, 387

Peeters, Z., Rodgers, S. D., Charnley, S. B., et al. 2006, A&A, 445, 197

Pickett, H. M., Poynter, R. L., Cohen, E. A., et al. 1998, JQSRT, 60, 883

Redman, M. P., Keto, E., Rawlings, J. M. C., & Williams, D. A. 2004, MNRAS, 352, 1365

Requena-Torres, M. A., Martín-Pintado, J., Rodríguez-Franco, A., et al. 2006, A&A, 455, 971

Requena-Torres, M. A., Marcelino, N., Jiménez-Serra, I., et al. 2007, ApJ, 655, L37

Taquet, V., Ceccarelli, C., & Kahane, C. 2012, A&A, 538, A42

van Dishoeck, E. F., Blake, G. A., Jansen, D. J., & Groesbeck, T. D. 1995, ApJ, 447, 760
## Appendix A: Spectroscopic line parameters

The frequencies, upper energy levels, upper level degeneracies, and Einstein coefficients of the observed transitions were taken from the JPL spectroscopic database (Pickett et al. 1998) and are summed up here in Table A.1.

<table><caption>**Table A.1.** Parameters of detected lines.</caption><thead><tr><th>Species</th><th>Transition</th><th>Frequency (MHz)</th><th>E<sub>up</sub> (K)</th><th>g<sub>up</sub></th><th>A<sub>ul</sub> (s<sup>-1</sup>)</th></tr></thead><tbody><tr><td>AA-CH<sub>3</sub>OCH<sub>3</sub></td><td>4<sub>14</sub>-3<sub>03</sub></td><td>99 326.072</td><td>10.2</td><td>90</td><td>5.533 &times; 10<sup>-6</sup></td></tr><tr><td>EE-CH<sub>3</sub>OCH<sub>3</sub></td><td>4<sub>14</sub>-3<sub>03</sub></td><td>99 325.217</td><td>10.2</td><td>144</td><td>5.533 &times; 10<sup>-6</sup></td></tr><tr><td>EA-CH<sub>3</sub>OCH<sub>3</sub></td><td>4<sub>14</sub>-3<sub>03</sub></td><td>99 324.362</td><td>10.2</td><td>36</td><td>5.533 &times; 10<sup>-6</sup></td></tr><tr><td>AE-CH<sub>3</sub>OCH<sub>3</sub></td><td>4<sub>14</sub>-3<sub>03</sub></td><td>99 324.364</td><td>10.2</td><td>54</td><td>5.533 &times; 10<sup>-6</sup></td></tr><tr><td>E-CH<sub>3</sub>CHO</td><td>5<sub>14</sub>-4<sub>13</sub></td><td>98 863.328</td><td>16.6</td><td>22</td><td>3.104 &times; 10<sup>-5</sup></td></tr><tr><td>A-CH<sub>3</sub>CHO</td><td>5<sub>14</sub>-4<sub>13</sub></td><td>98 900.948</td><td>16.5</td><td>22</td><td>3.108 &times; 10<sup>-5</sup></td></tr><tr><td>A-CH<sub>3</sub>OCHO</td><td>8<sub>17</sub>-7<sub>16</sub></td><td>100 490.682</td><td>22.8</td><td>34</td><td>1.435 &times; 10<sup>-5</sup></td></tr><tr><td>E-CH<sub>3</sub>OCHO</td><td>8<sub>17</sub>-7<sub>16</sub></td><td>100 482.241</td><td>22.8</td><td>34</td><td>1.434 &times; 10<sup>-5</sup></td></tr><tr><td>A-CH<sub>3</sub>OCHO</td><td>9<sub>09</sub>-8<sub>08</sub></td><td>100 683.368</td><td>24.9</td><td>38</td><td>1.472 &times; 10<sup>-5</sup></td></tr><tr><td>E-CH<sub>3</sub>OCHO</td><td>9<sub>09</sub>-8<sub>08</sub></td><td>100 681.545</td><td>24.9</td><td>38</td><td>1.472 &times; 10<sup>-5</sup></td></tr><tr><td>A-CH<sub>3</sub>OCHO</td><td>9<sub>19</sub>-8<sub>18</sub></td><td>100 080.542</td><td>25.0</td><td>38</td><td>1.444 &times; 10<sup>-5</sup></td></tr><tr><td>E-CH<sub>3</sub>OCHO</td><td>9<sub>19</sub>-8<sub>18</sub></td><td>100 078.608</td><td>25.0</td><td>38</td><td>1.444 &times; 10<sup>-5</sup></td></tr><tr><td>o-CH<sub>2</sub>CO</td><td>5<sub>15</sub>-4<sub>14</sub></td><td>100 094.514</td><td>27.5</td><td>33</td><td>1.031 &times; 10<sup>-5</sup></td></tr><tr><td>p-CH<sub>2</sub>CO</td><td>5<sub>05</sub>-4<sub>04</sub></td><td>101 036.630</td><td>14.6</td><td>11</td><td>1.104 &times; 10<sup>-5</sup></td></tr></tbody></table>

## Appendix B: Column density determination

The column densities were determined within the LTE approximation, using

$$N = \frac{8\pi\nu^3 Q(T_{\text{ex}})}{c^3 g_{\text{up}}A_{\text{ul}}} \frac{e^{E_{\text{up}}/kT_{\text{ex}}}}{e^{h\nu/kT_{\text{ex}}} - 1} \frac{1}{J_{\nu}(T_{\text{ex}}) - J_{\nu}(T_{\text{bg}})} \int T_{\text{mb}}dv, \quad (\text{B.1})$$

where $\nu$ is the frequency of the transition, $c$ the velocity of light, $h$ and $k$ the Planck and Boltzman constants, respectively, $g_{\text{up}}$ the upper level degeneracy, $A_{\text{ul}}$ the spontaneous emission Einstein coefficient, $T_{\text{ex}}$ the excitation temperature, $Q(T_{\text{ex}})$ the partition function at $T_{\text{ex}}$, $E_{\text{up}}$ the upper level energy, and $J_{\nu}(T)$ is the radiation temperature that is defined as

$$J_{\nu}(T) = \frac{\hbar\nu}{k} \frac{1}{e^{h\nu/kT} - 1}. \quad (\text{B.2})$$

Equation (B.1) assumes that the lines are optically thin. We note that since the temperatures are very low, $J_{\nu}(T_{\text{bg}})$ is not negligible with respect to $J_{\nu}(T_{\text{ex}})$.