Detection of interstellar hydrogen peroxide
P. Bergman, B. Parise, R. Liseau, B. Larsson, H. Olofsson, K. M. Menten, R. Güsten
aa r X i v : . [ a s t r o - ph . GA ] J un Astronomy&Astrophysicsmanuscript no. hooh c (cid:13)
ESO 2018October 31, 2018 L etter to the E ditor Detection of interstellar hydrogen peroxide ⋆ P. Bergman , B. Parise , R. Liseau , B. Larsson , H. Olofsson , K. M. Menten , and R. G ¨usten Onsala Space Observatory, Chalmers University of Technology, SE-439 92 Onsala, Sweden e-mail: [email protected] Max Planck Institut f¨ur Radioastronomie, Auf dem H¨ugel 69, 53121 Bonn, Germany Department of Earth and Space Sciences, Chalmers University of Technology, SE-439 92 Onsala, Sweden Department of Astronomy, Stockholm University, AlbaNova, SE-10691 Stockholm, SwedenReceived ?; accepted ?
ABSTRACT
Context.
The molecular species hydrogen peroxide, HOOH, is likely to be a key ingredient in the oxygen and water chemistry in theinterstellar medium.
Aims.
Our aim with this investigation is to determine how abundant HOOH is in the cloud core ρ Oph A.
Methods.
By observing several transitions of HOOH in the (sub)millimeter regime we seek to identify the molecule and also todetermine the excitation conditions through a multilevel excitation analysis.
Results.
We have detected three spectral lines toward the SM1 position of ρ Oph A at velocity-corrected frequencies that coincidevery closely with those measured from laboratory spectroscopy of HOOH. A fourth line was detected at the 4 σ level. We also foundthrough mapping observations that the HOOH emission extends (about 0.05 pc) over the densest part of the ρ Oph A cloud core. Wederive an abundance of HOOH relative to that of H in the SM1 core of about 1 × − . Conclusions.
To our knowledge, this is the first reported detection of HOOH in the interstellar medium.
Key words. astrochemistry – ISM: abundances – ISM: individual objects: ρ Oph A – ISM: molecules
1. Introduction
Hydrogen peroxide, HOOH, is believed to play an importantrole in the Earth’s atmospheric ozone and water chemistry. Itis a key constituent in the gas- and liquid-phase radical chem-istry and has an oxidizing potential in the liquid phase. Gas-phase HOOH has been seen in the Martian atmosphere byground-based observations (Clancy et al. 2004; Encrenaz et al.2004). However, recent Mars observations with the
Herschel
Observatory (Hartogh et al. 2010) failed to detect HOOH atlevels below those previously seen. The non-detection was at-tributed to seasonal variations.Interestingly, HOOH is among the simplest molecules thatshow internal rotation. The internal rotation, or torsion, mani-fests itself as a rotation of the two O-H bonds about the O-Obond. This hindered internal rotation can be described with atorsion potential in which the two minima (the most stable con-figurations) do not coincide with the cis or trans alignment ofthe two O-H bonds . The twofold barrier gives rise to a quartetof sublevels for each torsional state. These sublevels are denoted τ = , , , C † h pointgroup, see Hougen (1984) for a discussion), with only c -typetransitions and with a dipole moment of 1.6 D (Cohen & Pickett ⋆ Based on observations with the Atacama Pathfinder EXperiment(APEX) telescope. APEX is a collaboration between the Max-Planck-Institut f¨ur Radioastronomie, the European Southern Observatory, andthe Onsala Space Observatory. When the two O-H bonds point in the same direction, this is referredto as the cis position, while the 180 degree opposite case is called the trans position. For HOOH, the trans potential barrier height is 557 K,while the cis barrier height is almost 4000 K (Pelz et al. 1993). . × − with respect to H from their Orion spectral scandata. Boudin et al. (1998) reported an upper limit of 5.2 % ofsolid HOOH relative to H O ice toward NGC 7538 IRS9.As in the Earth’s atmosphere, HOOH is expected to beclosely connected to the water and molecular oxygen chemistryalso for those physical conditions prevailing in molecular clouds.In current pure gas-phase models, HOOH is formed by reactionof H with HO or via the reaction involving two OH radicals.However, these reactions proceed very slowly . Alternatively,Tielens & Hagen (1982) suggested that HOOH could be formedon grain surfaces by the successive additions of H atoms tomolecular oxygen. If this is the case, HOOH could be closelyrelated to the amount of O on grains.The 119 GHz line of O was detected toward the cloud ρ Oph A with an abundance of 5 × − relative to H byLarsson et al. (2007) using the Odin satellite. The ρ Oph Amolecular cloud, at a distance of about 120 pc, has been thesubject of several studies. Continuum observations (Andr´e et al.1993; Motte et al. 1998) and C O observations (Liseau et al.2010) revealed several cores. Very recently, Bergman andcoworkers found a very high degree of deuteration toward theSM1 core in ρ Oph A from observations of deuterated H CO ∼ eric P. Bergman et al.: Detection of interstellar hydrogen peroxide
Table 1.
Observed HOOH lines
Frequency Transition E u A ul (MHz) J ′ K ′ a , K ′ c − J ′′ K ′′ a , K ′′ c τ ′ − τ ′′ (K) (s − )219166.86 3 , − , − . × − , − , − . × − , − , − . × − , − , − . × − , − , − . × − , − , − . × − (Bergman et al. 2011). These observations suggest that grainsurface reactions were at work to produce the very high deu-terium levels observed in the gas-phase material.In this Letter we continue our study of the ρ Oph A cloud byreporting on observations of several HOOH transitions. We havedetected four HOOH lines and two of these lines were mappedover the central part of this cloud. In Sect. (2) we describe ourobservations and present our results. There, we also describe inmore detail the energy level structure and symmetry of HOOH,which is important for discussing the detected lines. Then, inSect. (3), we discuss the implications of our results.
2. Observations and results
We have used the APEX 12 m telescope located at about5100 m altitude in the Chilean Andes (G¨usten et al. 2006)to observe HOOH. For the lower frequency lines we usedthe Swedish heterodyne facility instruments APEX-1 andAPEX-2 (Vassilev et al. 2008). The 7-pixel longer wavelength(450 µ m) module of the MPIfR-built CHAMP + receiver array(Kasemann et al. 2006) was used for the observations of a high-frequency line. The 7 pixels, spaced by 18 ′′ , are arranged in ahexagon around a central pixel. The observations took place onseveral occasions during 2010; April 2-11, July 7, August 4-8,and September 10-13.The targeted HOOH lines are listed in Table 1. The fre-quency, transition quantum number designation, energy of upperlevel, and Einstein A -coe ffi cient are listed. All values have beencompiled from the JPL catalogue (Pickett et al. 1998). The fre-quency uncertainty for the listed lines is 0.1 MHz or better andthis corresponds to 0 .
14 km s − at 219 GHz. We also present inFig. 1 the energy diagrams for all levels below 100 K. Owingto the symmetry of HOOH, four di ff erent radiatively decoupledladders occur: A ↔ , A ↔ , B ↔ , and B ↔ . The subscript in-dicates the pair of torsional quantum numbers τ involved. The c -type electric dipole transitions must also obey τ = ↔ τ = ↔ A -specieshave a nuclear spin weight of 1 and the B -species have a spinweight of 3 (Hougen 1984). This is, of course, due to the nuclearspin directions of the two H atoms (in the same way as the orthoand para symmetries occur for H O or H CO).The ground-state symmetry species is that of A ↔ with the B ↔ state only about 2.5 K higher energy. This means that theenergy di ff erence between the species with di ff erent nuclear spinweights is much smaller than the corresponding di ff erence forH O or H CO. The torsional τ = , τ = , ff erence stems from the tunnel-ing through the trans barrier and is comparable to the gas kinetictemperatures of 20 −
30 K found in ρ Oph A (Loren et al. 1990;Liseau et al. 2003; Bergman et al. 2011).
Fig. 1.
HOOH energy level diagrams. The energy is given in Kon the vertical axis, and at the bottom the quantum numbers K a and τ are shown. To the left of each level the rotational quantumnumbers J K a , K c are listed. The upper diagram shows the levelswith a nuclear spin weight of 1 ( A ↔ and A ↔ ), while thoselevels with a spin weight of 3 ( B ↔ and B ↔ ) are shown inthe lower diagram. The allowed c -type transitions within eachof the four ladders are indicated as arrows. The red arrows rep-resent detected lines, while the red-dashed arrows indicate non-detections.In Fig. 2 we show the HOOH spectra toward the core of ρ Oph A. The upper five spectra are toward the SM1 core, α (J2000) = h m . s and δ (J2000) = − ◦ ′ ′′ . The670 GHz CHAMP + spectrum is an average of all pixels andis centered 30 ′′ north of the SM1 position (usually denotedSM1N). In Fig. 3 spectra of the 219166 MHz line for the ob- . Bergman et al.: Detection of interstellar hydrogen peroxide 3 Fig. 2.
HOOH spectra toward ρ Oph A. The transition is in-dicated in each spectrum. The T ∗ A intensity scale is in K andthe velocity ( v LSR ) scale is in km s − . The velocity resolutionis 0 .
25 km s − . Table 2.
Observed HOOH line velocities, widths, and intensities
Freq. Beam size v LSR
FWHM R T mb dv a (MHz) (arcsec) ( km s − ) ( km s − ) (K km s − )219167 28 3.8 0.84 0.167(0.018)251915 25 3.7 0.75 0.165(0.018)268961 23 3.7 (1.2) 0.040(0.011)318223 20 3.8 0.78 0.106(0.013)318712 20 < . < . b < . ca integr. from 2.5 to 4 . − , errors are 1 σ , upper limits are 3 σ b average value from three pixels closest to the SM1 position c average value from all pixels, central pixel on SM1N served map positions are shown. In the southern part the line iscentered at 3 . − , while further north, at o ff set (0 ′′ , + ′′ ),the line is at 3 . − . This NS velocity gradient is almostidentical to the one seen for H CO and its deuterated variants(Bergman et al. 2011). The narrow peaks to the NW are adjacentto where the sulphur species peak as noted by the same authors.In addition to the map in Fig. 3, we also mapped the 251 GHzline, albeit with a poorer S / N ratio and will not discuss it furtherhere.The beam sizes, velocity-integrated main-beam brightnesstemperatures ( R T mb dv ) , fitted LSR velocities ( v LSR ), and linewidths (FWHM) are tabulated in Table 2. The T mb -scale wasestablished assuming main-beam e ffi ciencies of 0.75, 0.73, and0.4, for APEX-1, 2, and CHAMP + , respectively. The upper lim-its are 3 σ and for the CHAMP + line two intensities are listed inTable 2; one for the three pixels closest to the SM1 position andthe other by averaging data from all pixels. The tabulated errorsand upper limits depend only on the channel noise. Fig. 3.
Map spectra of the 219166 MHz HOOH line toward ρ Oph A. The map o ff set (0 ′′ , ′′ ) corresponds to the SM1 posi-tion. The T ∗ A intensity scale, in K, and the velocity ( v LSR ) scale,in km s − , are indicated in the upper right spectrum. The velocityresolution is 0 .
25 km s − .Using the integrated line intensities in Table 2 (corrected bya beam-filling factor corresponding to a source size of 24 ′′ , seeBergman et al. (2011)), we performed a rotation diagram anal-ysis (Goldsmith & Langer 1999). From this we can determinethe rotation temperature, T rot , as well as the HOOH column den-sity, N (HOOH). The resulting rotation diagram is displayed inFig. 4. The fit is based on the detected τ = ↔ T rot = ± ± × cm − where the errors depend on the uncer-tainty of the integrated intensities and a calibration uncertaintyof 10 − CO at the SM1 posi-tion. Obviously, the non-detection of the τ = ↔ T rot and N (HOOH). Hence, we conclude that the τ = , ffi cient is large (Table 1) andfor a typical value of a collision coe ffi cient ( ∼ − cm s − ), thecritical density of the 670 GHz transition is about two orders ofmagnitude higher than the H density determined from H COand CH OH observations in the same source (Bergman et al.2011). From the energy diagrams (Fig. 1) it is clear that thereis a lack of radiative de-excitation routes out of the K a = τ = , τ = , τ = , K a = ffi cient. However,a full statistical equilibrium analysis is needed for understand-ing the details of the excitation. This also requires some basicknowledge of the collision coe ffi cients. For now we assume thatthe τ = , τ = , P. Bergman et al.: Detection of interstellar hydrogen peroxide
Fig. 4.
HOOH rotation diagram. The detected lines (only τ = ↔ σ upper limits (not included in the fit) areshown as open squares and downward arrows and originate fromtransitions between τ = , τ = , ± × cm − .
3. Discussion
Given the good agreement of the velocities of our detectedHOOH lines (together with the ≈ . − accuracy of the lab-oratory frequencies) with those from other species we are veryconfident that the lines belong to HOOH. From our mapping ofthe 219 GHz line it is also evident that the NS velocity gradientseen for HOOH reflects that of other species. Moreover, the de-rived rotation temperature of 22 ± ρ Oph A cloud core. For this fairly low excitation temper-ature one would not expect many lines from other species to bepresent and, using the JPL and Cologne databases (Pickett et al.1998; M¨uller et al. 2005), we found no lines from other speciesthat could possibly interfere with the identification. Of course,the narrow lines (with FWHM typically < − ) seen to-ward the ρ Oph A cloud core also make line confusion muchless likely.From the H CO and CH OH analysis of the SM1 coreBergman et al. (2011) determined an H column density of 3 × cm − . Assuming that the HOOH level populations mainlyreside in the τ = , × − . This is well below the limit of 4 . × − foundtoward Orion KL by Blake et al. (1987).According to current gas-phase schemes (e.g., the OSUchemical reaction database), formation of HOOH in the gasphase is not e ffi cient. Only two very slow reactions are proposedfor its formation, reaction of H with HO , or reaction of twoOH radicals.On grains, HOOH is formed through successive hydrogenadditions to O . This was first proposed by Tielens & Hagen(1982), based on theoretical arguments. Recent laboratory exper-iments have been made to investigate this route, up to formationof water molecules:O + H + H → HOOH (1)HOOH + H → H O + OH . (2)Miyauchi et al. (2008) have investigated the reaction of Hatoms with solid O at 10 K. Subsequently, Ioppolo et al. (2008) have investigated the same reaction in the temperature range 12-28 K. Both studies showed that the conversion of O stops atsome point before exhaustion of O because of shielding of O in the deepest layers. The experiment of Ioppolo et al. (2008)shows that the shielding decreases with increasing temperature,pointing to the fact that the O ice may become more porouswhen close to its sublimation temperature (30 K).Because this shielding may not be very relevant in space,Oba et al. (2009) investigated the formation of HOOH and H Owhen codeposing O and H in the temperature range 10-40 K.The H O / HOOH ratio in the formed ices is observed to dependstrongly on the temperature, and on the O / H flux. The mea-sured H O / HOOH ratio is lower than 5 in all experiments ( T =
10, 20 K and O / H-flux between 3 . × − and 1 . × − ), al-though higher values may be obtained in the case of lower O / H-flux. The present detection of HOOH may open the possibility toquantify the importance of reactions (1) and (2) in the formationof water.Ioppolo et al. (2008) have modeled the formation of water intypical dense clouds (their Figure 4). The laboratory results leadto a revision of the energy barriers involved in the models, andtheir new model (only accounting for the three main routes ofwater formation on the grains) predicts a fractional abundancefor HOOH of a few 10 − with respect to H nuclei. This is morethan three orders of magnitude lower than our detection.Further understanding will require detailed chemical model-ing of grain chemistry. This will be the scope of a forthcomingpaper. The observation of water in the ρ Oph A region with the
Herschel
Observatory as well as the confirmation of the O de-tection would also be very valuable in setting constraints on themodels. Acknowledgements.
We acknowledge the excellent observational support fromthe APEX sta ff . We are grateful to A. Gusdorf for doing some of the CHAMP + observations. BP is funded by the Deutsche Forschungsgemeinschaft (DFG) un-der the Emmy Noether project number PA1692 / References
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