ACETIC ACID IN THE HOT CORES OF SAGITARRIUS B2(N) AND W51 A. Remijan, 1 L. E. Snyder, 1 S.-Y. Liu, 2 D. Mehringer, 1,3 and Y.-J.

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1 The Astrophysical Journal, 576: , 2002 September 1 # The American Astronomical Society. All rights reserved. Printed in U.S.A. ACETIC ACID IN THE HOT CORES OF SAGITARRIUS B2(N) AND W51 A. Remijan, 1 L. E. Snyder, 1 S.-Y. Liu, 2 D. Mehringer, 1,3 and Y.-J. Kuan 4 Received 2002 February 20; accepted 2002 May 8 ABSTRACT We have detected interstellar acetic acid (CH 3 COOH) toward the hot core source W51e2. This is the first new source of interstellar CH 3 COOH since its discovery by Mehringer et al. toward the hot core source Sgr B2(N-LMH). In this paper, we report CH 3 COOH observations at two new frequencies toward Sgr B2(N- LMH) with the OVRO array and at 10 frequencies toward W51e2 with the Berkeley-Illinois-Maryland Association array. Toward Sgr B2(N-LMH) the agreement in positions, intensities, and velocities between the two lines from the previous study and the two new lines strongly indicates that all four CH 3 COOH lines are coming from a common source. Using all four detected transitions, we find an average column density of 6:1ð6Þ10 15 cm 2, a fractional abundance of ð0:8 6Þ10 10 relative to H 2 and ð3 6Þ10 2 relative to its isomer methyl formate (HCOOCH 3 ). Toward W51e2, we find a CH 3 COOH column density of 1:7ð5Þ10 16 cm 2 with a fractional abundance of 1: relative to H 2 and ð1 6Þ10 2 relative to HCOOCH 3. Furthermore, we find the distribution of CH 3 COOH toward W51e2 is coincident with HCOOCH 3, thus suggesting a similar formation mechanism. Subject headings: astrochemistry ISM: abundances ISM: individual (Sagittarius B2, W51) ISM: molecules 1. INTRODUCTION To date, interstellar acetic acid (CH 3 COOH) has only been reported in the hot core source Sgr B2(N-LMH) (Mehringer et al. 1997). However, there has been considerable interest in searching a variety of sources for interstellar acetic acid because it shares common structural elements with glycine (NH 2 CH 2 COOH), the simplest amino acid (Wootten et al. 1992; Snyder 1997; Mehringer et al. 1997). Consequently, an interstellar acetic acid source may also contain glycine. This is of great interest, for example, to astrobiologists who want to know the birthplaces of amino acids (Irion 2000). Furthermore, acetic acid is also important for astrochemical studies. Acetic acid is an isomer of both methyl formate (HCOOCH 3 ) and glycolaldehyde (CH 2 OHCHO). While methyl formate is easily detectable in many hot molecular cloud cores such as OMC-1 (Turner 1989), G (Mehringer & Snyder 1996), Sgr B2(N-LMH) (Miao et al. 1995), and W51(e1/e2 and d) (Liu, Mehringer, & Snyder 2001), glycolaldehyde has only been detected in Sgr B2(N-LMH) (Hollis, Lovas, & Jewell 2000; Hollis et al. 2001). Indeed, the structural configuration in these isomers seems to show a preference for a C O C O (e.g., methyl formate) backbone geometry rather than an O C C O (e.g., glycolaldehyde) or a C C O O (e.g., acetic acid) backbone geometry (Hollis et al. 2000). The mechanism by which a molecule is formed plays the largest role in determining its abundance. At present, there 1 Department of Astronomy, University of Illinois, Urbana, IL 61801; aremijan@astro.uiuc.edu, snyder@astro.uiuc.edu, dmehring@astro.uiuc.edu. 2 Department of Astronomy, California Institute of Technology, Pasadena, CA 91125; syl@astro.caltech.edu. 3 National Center for Supercomputing Applications, University of Illinois, Urbana, IL Department of Earth Sciences, National Taiwan Normal University; and Institute of Astronomy and Astrophysics, Academia Sinica, P.O. Box , 106 Taipei, Taiwan, Republic of China; kuan@sgrb2.geos.ntnu.edu.tw. 264 is no consensus on the formation of acetic acid in the interstellar medium. Huntress & Mitchell (1979) proposed that a gas-phase reaction of protonated ketene and water, CH 3 CO + +H 2 O! CH 3 COOH þ 2, followed by recombination with an electron could produce acetic acid. However, pure gas-phase reactions cannot account for the abundances measured for several complex species (Charnley, Tielens, & Millar 1992; Mehringer & Snyder 1996). Thus, grain surface chemistry has been invoked to account for the large abundance of complex species. Sorrell (2001) forms acetic acid, methyl formate, glycolaldehyde, and glycine by utilizing icy grain mantles photoprocessed by ultraviolet radiation. These reactions invoke a high concentration of carboxyl acid (COOH), which is closely related to formic acid (HCOOH). Ehrenfreund & Charnley (2000) suggest that protonated, mantle-formed molecules, such as methanol (CH 3 OH), offer a specific route to molecular complexity in hot cores through alkyl cation transfer to a neutral base. For example, S. B. Charnley (2000, private communication) suggests that evaporated formic acid (HCOOH) from dust grains can be the seed for more complex acids. CH 3 COOH can be formed by CH 3 OH þ 2 + HCOOH! CH 3COOH þ 2 + H 2 O, followed by recombination with an electron to give CH 3 COOH+H. Thus, it seemed that compact cores traced by enhanced abundances of HCOOH may be good regions to search for CH 3 COOH. Recently, Liu et al. (2001) mapped HCOOH emission toward five hot core sources: Sgr B2(N-LMH), W51(e1/e2 and d), and the Orion KL Compact Ridge. They found the derived HCOOH column density in the W51 hot core e2 was comparable to Sgr B2(N- LMH). Motivated by these favorable results, we searched W51 for CH 3 COOH with the Berkeley-Illinois-Maryland Association (BIMA) array. 5 As a result, CH 3 COOH was detected in W51 for the first time. In the interim, we detected 5 Operated by the University of California, Berkeley, the University of Illinois, and the University of Maryland with support from the National Science Foundation.

2 ACETIC ACID IN HOT CORES OF Sgr B2(N) AND W two new CH 3 COOH transitions in Sgr B2(N-LMH) with the Caltech Owens Valley Radio Observatory (OVRO) millimeter array. 6 These detections further solidified the original CH 3 COOH identification of Mehringer et al. (1997) and gave us enough data to compare the CH 3 COOH abundances in the two astrochemically interesting hot core regions Sgr B2(N-LMH) and W51e2. 2. OBSERVATIONS The Sgr B2(N-LMH) CH 3 COOH observations were made in 1996 November and December, and in 1997 May, with the OVRO array in its E, H, and L configurations, respectively. The W51e2 CH 3 COOH observations were made in 1999 April, and in 2000 April and May, with the BIMA array in its C configuration. Table 1 summarizes the observational parameters for both arrays. The first three rows list the telescopes used, the sources observed, and abbreviated observing frequencies. The next three list the flux, phase, and passband calibration sources. The following two rows give the bandwidths and corresponding channel widths in kilometers per second. The final two contain the array configurations and corresponding synthesized beamwidths in arcseconds. Table 2 lists the spectroscopic parameters of the observed lines. Observational details are discussed below OVRO Array Observations of Sgr B2(N-LMH) Using the OVRO array, we searched for the 10; 10 9; 9 E line at 111, MHz and the 10; 10 9; 9 A line at 111, MHz in the direction of Sgr B2(N-LMH) in order to follow up the detection by Mehringer et al. (1997) of the 8; 8 7; 7 (90.2 GHz) and the 9; 9 8; 8 (100.9 GHz) A and E lines. All CH 3 COOH rest frequencies are given in Table 2; they are taken from the new study by Ilyushin et al. (2001). The Sgr B2(N-LMH) observations were taken in the direction (J2000) = 17 h 47 m 19992, (J2000) = >5. The spectral window containing the CH 3 COOH lines had a spectral resolution of 0.49 MHz. The data were calibrated using the MMA software package of Caltech. The calibrated u-v data were combined, imaged, and self-calibrated using the MIRIAD software package (Sault, Teuben, & Wright 1995). 6 Operated by the California Institute of Technology with support from the National Science Foundation BIMA Array Observations of W51e2 Using the BIMA array, we searched W51e2 for the 10 CH 3 COOH lines listed in Table 2. The phase center was coincident with the hot core source W51e2 at (J2000) = 19 h 23 m 4399, (J2000)= >7. Observations of the 21; 21 20; 20 (228.7 GHz) and the 22; 22 21; 21 (239.3 GHz) A and E lines were performed in 1999 April. The spectral window containing these CH 3 COOH lines had a bandwidth of 100 MHz and was divided into 64 channels for a spectral resolution of 1.56 MHz. Table 1 lists the specific spectral resolution of these observations. Observations of the 8; 8 7; 7 (90.2 GHz), 9; 9 8; 8 (100.9 GHz), and 10; 10 9; 9 (111.5 GHz) A and E lines were performed in 2000 April and May. The spectral window containing these CH 3 COOH lines had a bandwidth of 25 MHz and was divided into 256 channels for a spectral resolution of MHz. Table 1 lists the specific channel velocity width for each of these observations. In order to better match the spectral resolution of all the CH 3 COOH lines, the lower frequency CH 3 COOH data were smoothed to 1.0 km s 1 spectral resolution. In addition to CH 3 COOH, we also imaged a pair of HCOOCH 3 A and E lines at 90.1 and GHz (Table 2). All data were combined, imaged, and selfcalibrated using the MIRIAD software package. 3. RESULTS 3.1. Sgr B2(N-LMH) Table 3 lists the molecular species found in our frequency band toward Sgr B2(N-LMH). Figure 1 shows the detection of the CH 3 COOH lines at 3 mm. The gray scale is the GHz continuum of the Sgr B2(N) region. The contours indicate the location of the 111, and the 111, MHz CH 3 COOH emission. In addition, we have labeled the positions of the ultracompact H ii regions K1-K3 and K5 as given by Gaume & Claussen (1990). The CH 3 COOH emission peaks given by Mehringer et al. (1997) are seen as large crosses with error bars. The phase self-calibration technique that was performed on these data results in some uncertainty in absolute positions. We have attempted to minimize this uncertainty by adjusting the image headers of the present data so that the absolute positions of the Sgr B2 north and main continuum peaks agree well between this study and the previous acetic acid study of Mehringer et al. (1997). We estimate the absolute positions to be uncertain by 1 00 in both right ascension and declination. From the TABLE 1 Observational Parameters Parameter 1996 Nov Dec; 1997 May 1999 April 1999 April 2000 April 2000 April 2000 May Array... OVRO BIMA BIMA BIMA BIMA BIMA Source... Sgr B2 W51 W51 W51 W51 W51 Observed frequency (GHz) Flux calibration... Uranus Uranus Uranus Jupiter Jupiter Jupiter Phase calibration... NRAO Passband calibration... 3C 273 3C 273 3C 273 3C 454 3C 454 3C 273 Bandwidth (km s 1 ) Channel width a (km s 1 ) Array... E, H, L C C C C C Beam (arcsec) a This is the unsmoothed velocity width of each channel.

3 266 REMIJAN ET AL. Vol. 576 Transition [J 0 ðk 0 1 K0 þ1 Þ JðK 1K þ1 Þ] TABLE 2 Molecular Line Parameters Frequency (MHz) CH 3 COOH a S i; j l 2 (D 2 ) E u (K) 8; 8 7; 7 E... 90, (20) b ; 8 7; 7 A... 90, (50) ; 9 8; 8 E , (20) ; 9 8; 8 A , (20) ; 10 9; 9 E , (20) ; 10 9; 9 A , (20) ; 21 20; 20 E , (50) ; 21 20; 20 A , (50) ; 22 21; 21 E , (50) 20.5 c ; 22 21; 21 A , (50) 20.5 c HCOOCH 3 d 7 2; 5 6 2; 4 E... 90, (12) ; 5 6 2; 4 A... 90, (13) ; ; 12 E , (10) ; ; 12 A , (11) a Each of the CH 3 COOH linesconsists of two a-type and two b-type degenerate transitions. The listed S i; j l 2 value is the sum of all four transitions in each group. This degeneracy is indicated by the asterisk substituted for the K_ quantum numbers (Ilyushin et al. 2001; F. J. Lovas 2001, private communication; I. Kleiner 2001, private communication). b Errors are given by Ilyushin et al The error in the frequency is one standard deviation uncertainty for the calculated frequency in MHz. c For this transition, the a and b components of the dipole moment tend to cancel each other due to the mixing of the eigenfunctions in DK 1. This causes a significant decrease in the line strength (I. Kleiner 2001, private communication). d Rest frequencies of methyl formate taken from Oesterling et al figure, we see that the emission peaks lie very close to the peaks detected in the previous investigation. The agreement in positions strongly indicates that all four CH 3 COOH lines are indeed from a common source. Least-squares Gaussian fits were made for each spectral line in order to obtain the radial velocities, peak intensities, and line widths for the detected transitions W51e1/e2 The W51 complex has been studied by numerous investigators. The centimeter emission from the W51 region was shown by Martin (1972) to consist of eight distinct components, W51a h. The strongest two of the eight components, W51e and W51d, were studied by Scott (1978), who found the two compact continuum components e1 and e2. Because both e1 and e2 are thought to be important star-forming cores, their magnetic fields were mapped interferometrically by Lai et al. (2001). At 3 mm wavelength, Liu et al. (2001) imaged four formic acid (HCOOH) transitions near the W51e2 systemic velocity of 55.0 km s 1. Ethyl cyanide (CH 3 CH 2 CN) was observed at 57.0 km s 1 and HCOOCH 3 at 55.3 km s 1. At 2 mm, Zhang, Ho, & Ohashi (1998) imaged dimethyl ether (CH 3 OCH 3 ), HCOOCH 3, and methyl cyanide (CH 3 CN) at 55 km s 1. However, they noted asymmetric profiles in their lines, especially CH 3 CN, observed near the location of the W51e2 continuum emission peak. While the central velocity coincides with the systemic velocity of W51e2, a prominent blueshifted peak showed up 2 3 km s 1 lower than the systemic velocity. At their spectral (0.5 km s 1 ) and spatial (1 00 ) resolution, they inferred this asymmetric profile to be evidence of infall. Table 4 lists the molecular species that we detected toward W51e2 at or near the systemic velocity. Figure 2a shows the CH 3 COOH spectra at 3 mm wavelength toward e2. Figure 2b shows the spectra of the 7 2; 5 6 2; 4 A and E transitions of HCOOCH 3 (Table 4). In order to determine the v LSR, leastsquares Gaussian fits of the HCOOCH 3 lines at 90,145.6 and 90,156.5 MHz, were made. Emission from the HCOOCH 3 transitions peaks in W51e2 with a v LSR of 55.1 km s 1 for 90,145.6 MHz, and 55.8 km s 1 for 90,156.5 MHz. Thus, an average systemic v LSR of the W51e2 hot core is 55.5 km s 1. The dotted line in Figure 2a is centered on this systemic velocity. The 90, MHz CH 3 COOH line is contaminated by U90.203; this unidentified line was also observed toward Sgr B2(N-LMH) by Mehringer et al. (1997). In addition, the 100, MHz CH 3 COOH line is blended with the E line of CH 3 SH, which was also detected by Mehringer et al. (1997) toward Sgr B2(N-LMH). The remaining CH 3 COOH lines at 90,246.25, 100,855.44, 111,507.27, and 111, MHz all lie near the W51e2 systemic velocity of 55.5 km s 1. Figure 2c shows the overlap between the 7 2; 5 6 2; 4 E HCOOCH 3 (thin line contours), and the 9; 9 8; 8 E transition of CH 3 COOH (thick line contours). The gray scale illustrates the continuum emission at 100,855 MHz from the W51e1/e2 region. Figure 2d shows the overlap between the 7 2; 5 6 2; 4 E HCOOCH 3 (thin line contours), and the 10; 10 9; 9 A transition of TABLE 3 Sgr B2(N-LMH) Molecular Line Identifications (MHz) Species Transition I 0 (Jy beam 1 ) Dv (km s 1 ) 111, U111, , (20)... AcA a 10; 10 9; 9 E , (5) b... H 13 COOH 5 0; 5 4 0; , U111, c 111, U111, , U111, , (20)... AcA 10; 10 9; 9 A a AcA = Acetic Acid (CH 3 COOH). b The rest frequency coincides with the 5 0; 5 4 0; 4 transition of H 13 COOH at a frequency of 100, MHz. (Pickett et al. 1998). c Blended line.

4 No. 1, 2002 ACETIC ACID IN HOT CORES OF Sgr B2(N) AND W Fig. 1. (a) Spectra toward Sgr B2(N-LMH) (Hanning-smoothed over three channels). The rms noise level is 80 mjy beam 1 (indicated by the vertical bar at the left of each spectrum). The spectral line labels correspond to the rest frequencies in Table 2 for a v LSR of 64 km s 1.(b) Emission contours from the 10; 10 9; 9 E transition of CH 3 COOH at GHz overlaid on the Sgr B2(N-LMH) continuum emission at GHz (gray scale). The numbers on the gray scale wedge are in units of janskys per beam. Contours indicate the location of the CH 3 COOH emission averaged between 63 and 65 km s 1. At GHz, the contour levels are 0.16, 0.2, 0.24, 0.264, and Jy beam 1. The large crosses with the error bars indicate the positions of the previous CH 3 COOH detections by Mehringer et al. (1997). The small crosses indicate the positions of the UC H ii regions K1-K3 and K5 as given by Gaume & Claussen (1990). (c) Emission contours from the 10; 10 9; 9 A transition of CH 3 COOH at GHz overlaid on the Sgr B2(N-LMH) continuum emission at GHz, as in (b). The CH 3 COOH contour levels are 0.16, 0.24, 0.28, 0.32 and 0.36 Jy beam 1. The synthesized beam of 3>7 2>3 is indicated at the bottom left corner. CH 3 COOH (thick line contours). The gray scale illustrates the continuum emission at 111,549 MHz from the W51e1/ e2 region. In both Figures 2c and 2d, there is general agreement between the locations of the e2 continuum emission, CH 3 COOH, and its isomer HCOOCH 3. Thus, we conclude all the emission is coming from a common source. Figure 3a shows the detection of CH 3 COOH at 1 mm wavelength toward W51e2. The 21; 21 20; 20 A transition at 228, MHz is blended with U The corresponding E transition is blended with the 18 5; ; 12 A transition of HCOOCH 3 at 228, MHz (Table 4). However, because of the strong line strengths of CH 3 COOH at these frequencies, least-squares Gaussian fits were possible to determine the line intensity and line width. The 22; 22 21; 21 E transition of CH 3 COOH at 239, MHz is possibly blended with the 13 12; ; 2 v 8 ¼ 1 transition of CH 3 CN at 239, MHz. The 22; 22 21; 21 A transition of CH 3 COOH at 239, MHz clearly is blended with U Note that the detection of both of these GHz CH 3 COOH lines is problematic anyway,

5 268 REMIJAN ET AL. TABLE 4 W51e2 Molecular Line Identifications (MHz) Species Transition I 0 (Jy beam 1 ) Dv (km s 1 ) 90, (12) a... MeF b 7 2; 5 6 2; 4 E , (13)... MeF 7 2; 5 6 2; 4 A , c... HCOOH 4 2; 2 3 2; d 90, U d 90, U d 90, (20)... AcA e 8; 8 7; 7 E bl f 90, U , (50)... AcA 8; 8 7; 7 A , U , U , (20)... AcA 9; 9 8; 8 E , (20)... AcA 9; 9 8; 8 A d 100,898.5 (1) g... CH 3 SH E d 111, (20)... AcA 10; 10 9; 9 E , (5) h... H 13 COOH 5 0; 5 4 0; , U , U , (20)... AcA 10; 10 9; 9 A , (10)... MeF 18 5; ; 12 E , U , (11)... MeF 18 5; ; 12 A d 228, (50)... AcA 21; 21 20; 20 E d 228, U , (50)... AcA 21; 21 20; 20 A , (50)... AcA 22; 22 21; 21 E bl f 239, (8) i... CH 3 CN 13 12; ; 2 v 8 ¼ d 239, U d 239, (50)... AcA 22; 22 21; 21 A bl f 239, U d a Rest frequencies of methyl formate taken from Oesterling et al b MeF = methyl formate (HCOOCH 3 ). c Rest frequency of formic acid taken from Liu et al d Blended line. e AcA = acetic acid (CH 3 COOH). f If the AcA transition was completely blended with an interloper, no attempt was made to determine a peak intensity (I o ) or line width (Dv) of the AcA line. g The rest frequency coincides with the E transition of CH 3 SH at a frequency of 100, MHz. (Lees & Mohammadi 1980). h The rest frequency coincides with the 5 0; 5 4 0; 4 transition of H 13 COOH at a frequency of 100, MHz. (Pickett et al. 1998). i The rest frequency coincides with the 13 12; ; 2 v 8 ¼ 1 transition of CH 3 CN at a frequency of 239, MHz. (Pickett et al. 1998). because their intensities are lowered by the line strength decrease noted in Table 2. Once again, to determine the v LSR, least-squares Gaussian fits were made using the 18 5; ; 12 A and E lines of HCOOCH 3. Emission from the transitions peaks in the W51e2 region at 53.5 and 53.4 km s 1, respectively. However, with our coarse spectral resolution (2 km s 1 ), the v LSR of both the HCOOCH 3 and CH 3 COOH transitions at 1 mm wavelength is well within 1 of the v LSR Gaussian fit at 3 mm wavelength. In Figure 3b, the gray scale is the 228,691 MHz continuum emission toward W51e2 and W51e1. Thick contours show the emission from the 21; 21 20; 20 A transition of CH 3 COOH and thin contours trace 18 5; ; 12 E transition of HCOOCH 3. Here the emission from the two isomers is offset by 2 00, which may not be significant with respect to our 4>3 3>2 beam (Table 1). In all, we find the distribution of CH 3 COOH toward W51e2 is coincident with HCOOCH 3, thus suggesting a similar formation mechanism. 4. DISCUSSION 4.1. Column Densities In several hot core sources, the measured abundances of many large molecular species were higher than what was predicted from gas-phase chemistry (Millar et al. 1988; Charnley et al. 1992; Kuan & Snyder 1994; Mehringer & Snyder 1996). The abundances found may be accounted for in two ways. First, the molecules may be formed by grain surface chemistry then evaporated into the gas phase. Second, precursor species with enhanced abundances may be evaporated off the mantles, and then warm gas-phase reactions may form the larger molecules. Therefore, determining column densities is very important because it puts constraints on the formation mechanism of large molecular species. Observationally, the total column density can be determined from a rotational temperature diagram (a plot of the logarithm of normalized column density vs. rota-

6 Fig. 2. (a) Spectra toward W51e2 (Hanning-smoothed over three channels). The rms noise level is 30 mjy beam 1 (indicated by the vertical bar at the left of each spectrum). The spectral line labels correspond to the rest frequencies in Table 2 for a v LSR of 55.5 km s 1. The dotted line is centered on the systemic velocity of W51e2 (55.5 km s 1 ). The ordinate of each spectrum ranges from 0 to 0.2 Jy beam 1 except for the spectrum at GHz. There, the ordinate ranges from 0 to 0.4 Jy beam 1.(b) The 7 2; 5 6 2; 4 A and E emission spectra of methyl formate (HCOOCH 3 ) toward W51e2 (Hanning-smoothed over three channels). The rms noise level is 0.2 Jy beam 1 (indicated by the vertical bar at the left of the spectrum). Once again, the dotted line is centered on the systemic velocity of W51e2 (55.5 km s 1 ). (c) Emission from the 9; 9 8; 8 E transition of CH 3 COOH at GHz (thick contours) and the 7 2; 5 6 2; 4 E transition of HCOOCH 3 at GHz (thin contours) averaged between 53 and 57 km s 1 overlaid on the W51e1/e2 continuum emission at GHz (gray scale). The numbers on the gray-scale wedge are in units of janskys per beam. The CH 3 COOH contour levels are 0.04, 0.054, 0.074, 0.094, and Jy beam 1. The HCOOCH 3 contour levels are 0.4, 0.6, 0.8, 1.0, 1.2, 1.4, 1.6, 1.8, and 2.0 Jy beam 1. Both sets of contours peak around W51e2. The synthesized beam of 8>8 6>4 is indicated at the bottom left corner. (d) Emission from the 10; 10 9; 9 A transition of CH 3 COOH at GHz overlaid on the W51e1/e2 continuum emission at GHz. The CH 3 COOH contour levels are 0.08, 0.08, 0.10, 0.12, and 0.14 Jy beam 1. The HCOOCH 3 contour levels are the same as in Fig. 2c. The synthesized beam of 8>1 5>7 is indicated at the bottom left corner.

7 270 REMIJAN ET AL. Vol. 576 Fig. 3. (a) Spectra toward W51e2 (Hanning-smoothed over three channels). The rms noise level is 0.2 Jy beam 1 in the upper spectrum and 0.3 Jy beam 1 in the lower spectrum (indicated by the vertical bars at the left). The spectral line labels correspond to the rest frequencies in Table 2 for a v LSR of 55.5 km s 1. The dotted line is centered on the systemic velocity of W51e2 (55.5 km s 1 ). The 21; 21 20; 20 A transition of CH 3 COOH is blended with U The corresponding E transition is blended with the 18 5; ; 12 A transition of HCOOCH 3 at 228, MHz. The 22; 22 21; 21 E transition of CH 3 COOH at 239, MHz is possibly blended with the 13 12; ; 2 v 8 ¼ 1 transition of CH 3 CN at 239, MHz. The 22; 22 21; 21 A transition of CH 3 COOH at 239, MHz clearly is blended with U (b) Emission from the 21; 21 20; 20 A transition of CH 3 COOH at GHz (thick contours) and the 18 5; ; 12 E transition of HCOOCH 3 at GHz (thin contours) overlaid on the W51e1/e2 continuum emission at GHz (gray scale). The numbers on the gray scale wedge are in units of janskys per beam. The CH 3 COOH emission is averaged between 53 and 60 km s 1 and the contour levels are 0.6, 0.8, 1.0, 1.2, 1.4, 1.6, and 1.8 Jy beam 1. The HCOOCH 3 contour levels are 2.0, 3.0, 4.0, 5.0, 6.0, 7.0, 8.0, 9.0, and 10.0 Jy beam 1. The crosses indicate the positions of the W51e1/e2 and e8 hot cores as given by Zhang et al. (1998). Both CH 3 COOH and HCOOCH 3 contours peak around W51e2, and the peaks are offset by 2 00 or less. The synthesized beam is 4>3 3>2 as shown at the bottom left. tional energy). With the assumption of LTE and optically thin emission, we have N T ¼ 2: R DIdv a b Q r e Eu=Tr ð1þ 3 S ij l 2 from Miao et al. (1995), and N u ¼ g u e E u=t r ð2þ N T Q rot from R the Boltzmann equation. In equations (1) and (2), DIdv is the integrated line intensity (Jy beam 1 km s 1 ), N T is the total column density (cm 2 ), N u is the upper level column density (cm 2 ), a and b are the FWHM beam sizes (arcsec), Q rot is the rotational partition function, E u is the upper state energy level (K), T r is the rotational temperature (K), is the rest frequency (GHz), S ij l 2 is the product of the transition line strength and the square of the electric dipole moment (D 2 ), and g u is the upper-level rotational degeneracy factor. Because CH 3 COOH is an asymmetric rotor, its rotational partition function is well approximated by Q r ¼ CTr 3=2, where C ¼ 14:1 K 3=2 (Mehringer et al. 1997). By combining the above two equations with a single rotational temperature and a single total column density, one finds that lnðn u =g u Þ is a linear function of E u with the slope of the line equal to 1=T r. Thus, a plot of lnðn u =g u Þ versus E u will yield a rotational temperature and a total column density CH 3 COOH Column Density in Sgr B2(N-LMH) In order to analyze our Sgr B2(N-LMH) data, we convolved it with a beam size of 11>2 4>3, which is the beam size of the previous observations of Mehringer et al. (1997). This allows a more meaningful comparison of column densities. Because of the large ordinate scatter and close E u spacings of the observed transitions, we found that a plot of lnðn u =g u Þ versus E u does not produce a reliable temperature. Instead we used the temperature found by Pei, Liu, & Snyder (2000) of 170(13) 7 K from their observations of methanol (CH 3 OH). From that temperature, we determined 7 The uncertainty (1 ) is shown in parenthesis, with the number of digits referring to the last digit(s) given.

8 No. 1, 2002 ACETIC ACID IN HOT CORES OF Sgr B2(N) AND W a total CH 3 COOH column density from all four transitions of N CH3 COOH ¼ 6:1ð6Þ10 15 cm 2, which is comparable to the value of 7: cm 2 found by Mehringer et al. (1997); however, they used T rot ¼ 200 K. For an HCOOCH 3 column density range of N HCOOCH3 ¼ ð1 2Þ10 17 cm 2 (Kuan & Snyder 1994; Liu et al. 2001), the relative CH 3 COOH/HCOOCH 3 abundance ratio is ð3 6Þ10 2. This is comparable to the previous measurements made by Mehringer et al. (1997) of ð4 7Þ10 2. For a H 2 column density range of N H2 ¼ð1 8Þ10 25 cm 2 (Lis et al. 1993; Kuan, Mehringer, & Snyder 1996), the CH 3 COOH fractional abundance range is X CH3 COOH ¼ð0:8 6Þ Mehringer et al. (1997) found ð0:9 7Þ CH 3 COOH Column Density in W51e2 For the CH 3 COOH emission from W51e2, the plot of lnðn u =g u Þ versus E u showed large scatter at 3 mm and has only two points at the same upper level energy at 1 mm wavelength. Consequently, we used two approaches. First, we calculated lnðn u =g u Þ using the true beam sizes (Table 1). This gave a nonphysical negative temperature, which suggests that the CH 3 COOH emission source in W51e2 is highly concentrated with respect to our beam size. Indeed, Zhang et al. (1998) found a source size of 2>4 from methyl cyanide (CH 3 CN) in the e2 core. When 2>4 is adopted for the CH 3 COOH source size, we derived a rotational temperature of 201(40) K and a column density of 1:7ð5Þ10 16 cm 2 (Fig. 4). Using C 2 H 5 CN, Liu et al. (2001) estimated an average rotational temperature of K for their five hot core regions, including W51e2. Thus, our temperature result is in general agreement. Because we have only four HCOOCH 3 lines with two separate energy levels (Table 2), a rotational temperature diagram is not meaningful. Clearly, the W51e1/e2 region would be a meaningful target for future extensive interferometric observations of HCOOCH 3 spectra in order to establish a rotational temperature diagram. In the meantime, we have adopted a rotational temperature of 201 K for HCOOCH 3, which gives a column density of ð3 15Þ10 17 cm 2 for the HCOOCH 3 Fig. 4. Rotational temperature diagram for the observed CH 3 COOH transitions toward W51e2. Ordinate:lnðN u =g u Þ. Abscissa: E u (K). The leastsquares fit is shown as a solid line. The resulting rotational temperature is 201(40) K for an assumed source size of 2>4. lines in Table 2. 8 Thus, our best estimate for the relative CH 3 COOH/HCOOCH 3 abundance ratio ranges from to This abundance ratio in W51e2 is in general agreement with the abundance ratio found by Mehringer et al. (1997) toward Sgr B2(N-LMH) [ð4 7Þ10 2 ]. Lastly, for an H 2 column density of cm 2 (Jaffe, Becklin, & Hillebrand 1984), the best estimate for the CH 3 COOH fractional abundance is X CH3 COOH ¼ 1: This fractional abundance in W51e2 is higher than the estimated fractional abundance range found by Mehringer et al. (1997) toward Sgr B2(N-LMH) [ð0:9 7Þ10 10 ] Formation of Acetic Acid As mentioned in x 1, there are formation mechanisms for interstellar acetic acid both in the gas phase and on grain surfaces. The radiative association mechanism proposed by Huntress & Mitchell (1979) is CH 3 CO þ + H 2 O! CH 3 COOH þ 2, followed by dissociative recombination with an electron to give CH 3 COOH+H. They also outlined a formation mechanism for the acetic acid isomer methyl formate. Wlodarczak & Demaison (1988) used this approach to predict a theoretical fractional abundance of for acetic acid and for methyl formate in Sgr B2. This gives a CH 3 COOH/HCOOCH 3 abundance ratio of Wootten et al. (1992) used the same radiative association approach with a slightly different set of assumptions to predict a CH 3 COOH/HCOOCH 3 ratio of 0.1 at a gas temperature of 30 K, but the ratio drops to 0.01 at 100 K. We note three problems associated with these radiative association predictions of the acetic acid and methyl formate abundances. The first problem is that the abundance ratios determined from our observations are larger than the high-temperature radiative association predictions (T rot > 100 K). Typically, gas-phase reaction models have a very steep inverse temperature dependence in their reaction rates. A temperature of 200 K should, for example, reduce the ratio predicted by Wootten et al. (1992) to less than 0.01, which would put the prediction for W51e2 even further from our observed values for the CH 3 COOH/HCOOCH 3 abundance ratio. The second problem is that many of the radiative association rate coefficients are not well determined. The third problem is that many of the key atomic and molecular fractional abundances (e.g., C þ,ch þ 3,CHþ 5 ) and the fractional electron abundances for hot molecular cores such as Sgr B2(N-LMH) and W51e2 are not well known. Correct fractional abundances are essential to accurate predictions. Ehrenfreund & Charnley (2000) proposed a warm gas-phase route in which reactions of protonated methanol (CH 3 OH þ 2 ) and HCOOH evaporated from grain surfaces can result in the formation of CH 3 COOH in the gas phase. However, they do not provide an abundance prediction based on their model. Recently, Ikeda et al. (2001) compared the observed fractional abundances of several complex species, including HCOOCH 3, C 2 H 5 OH, and C 2 H 5 CN, to calculated fractional abundances found from gas-phase formation models. Their observed fractional abundances (10 9 to 10 8 ) in some cases were several orders of magnitude greater than what was predicted from their gas-phase formation models. While our fractional abundan- 8 In the HCOOCH 3 calculations, we assumed a source size of 2>4, and Q r ¼ 12:454T 3=2 r in eq. (1) (Liu et al. 2001).

9 272 REMIJAN ET AL. Vol. 576 ces (10 10 to 10 9 ) are slightly lower than that of Ikeda et al. (2001), our value could be considered a lower limit, given the compact gas distribution and beam sizes used in our observations (Liu & Snyder 1999). Thus, our derived acetic acid fractional abundance could be significantly greater than those predicted values given by gas-phase reaction models. Many formation models of large species on grain mantles possess a minimal set of reaction mechanisms: (1) activationless C, O, N, and H atom additions to atoms and radicals, (2) H atom addition reactions, possessing activation energy barriers, to multiply bonded molecules, and (3) C atom additions to stable organic radicals, leading to carbon chain growth (Allen & Robinson 1977; Charnley et al. 1992; Hiraoka et al. 1998). In a recent model by Sorrell (2001), a combination of surface hydrogenation, photolysis, and radical reactions is considered. Sorrell assumes grain mantles are composed of a mixture of CO and hydrogenated species such as H 2 O, CH 4, and NH 3. Then a high concentration of free OH, CH 3, and NH 2 radicals is created in the mantle mainly by photolysis reactions. Once these radicals are created, they remain frozen in position until the grain heats up. As the grain heats up, the radicals become mobile and undergo chemical reactions with themselves and other molecules to produce complex organic substances. The advantage of radical-radical reactions is that they occur with almost no activation energy. Accordingly, Sorrell (2001) suggested that it is possible to form all three isomers, CH 3 COOH, HCOOCH 3, and CH 2 OHCHO, in this scenario. If all these isomers are indeed synthesized on grain mantles, there is still a need to specify an efficient grain desorption mechanism that can deposit large molecules into the gas phase, where they are observed. It is generally thought the molecules formed via surface reactions are evaporated by thermal heating and are thereafter deposited into the gas phase. The estimated CH 3 COOH rotation temperature of over 100 K appear to be consistent with this picture, as most mantle species are thought to be returned to the gas phase when the dust temperature exceeds 100 K (Tielens & Hagen 1982; Hasegawa, Herbst, & Lueng 1992). On the other hand, Sorrell (2001) also suggested a more violent desorption mechanism. In the model, significant amounts of chemical energy released from radical reactions will lead to the explosion of grain mantles, which then releases surface species into the gas phase. We note that the location of acetic acid emission, nearly coincident with the emission from other large molecular species thought to be formed by grain-surface chemistry in both W51 and Sgr B2(N) (Mehringer et al. 1997; Liu et al. 2001), strongly suggests that grain-surface chemistry is important in the formation of acetic acid. However, the proximity of emission does not necessarily imply a common formation mechanism. For example, CO is a good tracer of molecular hydrogen (H 2 ) and yet each has a very different formation mechanism (see, e.g., Hollenbach & Salpeter 1971 and Frerking, Langer, & Wilson 1982 for a review of H 2 formation and its distribution relative to CO). Nevertheless, the fact that acetic acid is coincident with other mantleformed species may still trace its formation pathway. For example, if acetic acid is formed via the warm gas-phase reaction paths suggested by Ehrenfreund & Charnley (2000), the acetic acid emission will still be coincident with the emission from mantle-formed species because the reactions take place after the precursor molecules are evaporated off the grain surface. Finally, Sorrell s model requires a substantial UV flux ( photons cm 2 s 1 ) to photoprocess the grain mantles for subsequent interstellar chemistry to occur. However, the UV flux in dense interstellar clouds is primarily due to cosmic rays and is on the order of 10 3 photons cm 2 s 1 (Ehrenfreund et al. 2001). Also, we note that our observations show the relative abundance ratio of CH 3 COOH/HCOOCH 3 is ð1 6Þ10 2 in W51e2 and ð3 6Þ10 2 in Sgr B2(N-LMH). However, Sorrell s model does not account for the preference in the formation pathway leading to HCOOCH 3 ; a future grain surface chemistry model will be needed to account for this. 5. SUMMARY We have observed two more emission lines in Sgr B2(N- LMH) from interstellar CH 3 COOH near GHz using the OVRO millimeter array. The positions, intensities, and velocities of these lines agree well with the two lines previously observed by Mehringer et al. (1997) at lower frequencies with the OVRO and BIMA arrays. The agreement between all four CH 3 COOH lines strongly indicate they are indeed coming from a common source. Using a beam size of 11>2 4>3, we derive an average column density of N CH3 COOH ¼ 6:1ð6Þ10 15 cm 2 from all four observed transitions, a fractional abundance of X CH3 COOH ¼ ð0:8 6Þ10 10,andaCH 3 COOH/HCOOCH 3 abundance ratio of ð3 6Þ10 2. Our measurements agree very well with the previous measurements made by Mehringer et al. (1997). We used the BIMA array to make the first detection of interstellar CH 3 COOH toward W51. This is the first new CH 3 COOH source detected outside of Sgr B2(N-LMH). The emission peak of CH 3 COOH and its isomer HCOOCH 3 are nearly cospatial suggesting a similar formation mechanism. The formation mechanisms that have been proposed to form interstellar CH 3 COOH include both gasphase reactions and radical-radical reactions on grain surfaces. Using a source size of 2>4, we derive a rotational temperature of 201(40) K and an average column density of N CH3 COOH ¼ 1:7ð5Þ10 16 cm 2,aCH 3 COOH fractional abundance of X CH3 COOH ¼ 1:7 10 9, and a relative CH 3 COOH/HCOOCH 3 abundance ratio of ð1 6Þ10 2. At present, the HCOOCH 3 column densities in both Sgr B2(N-LMH) and W51e2 are not as well determined as the CH 3 COOH column densities. Clearly, future interferometric measurements of HCOOCH 3 will be important for lowering the current uncertainties in the abundance ratios of CH 3 COOH to HCOOCH 3. Even within the uncertainties, the abundance ratios of CH 3 COOH to HCOOCH 3 in both Sgr B2(N-LMH) and W51e2 do not support published ion-molecule formation chemistry models. Furthermore, the proximity of CH 3 COOH emission to the emission from other large molecular species in both regions supports a formation mechanism using either grain-surface reactions or warm gas-phase reactions that occur after mantle-formed species are evaporated off the grain surface. We thank F. J. Lovas, I. Kleiner, and J. T. Hougen for their many useful calculations on the acetic acid transitions and line strength interpretations. We thank S. Charnley for providing us with a possible formation pathway for acetic

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