Experimental Investigations into Astrophysically ... - ACS Publications

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Experimental Investigations into Astrophysically Relevant Ionic Reactions Wolf D. Geppert and Mats Larsson* Department of Physics, AlbaNova University Center, Stockholm University, Roslagstullsbacken 21, Stockholm SE-10691, Sweden S Supporting Information *

1. INTRODUCTION There is not a single Nobel Prize in chemistry awarded to an extraterrestrial chemical discovery. This is in sharp contrast to the Nobel Prizes in physics, where about 10 prizes have been awarded based on discoveries of extraterrestrial phenomena within the field of physics, the first being the 1936 physics prize to V. F. Hess “for his discovery of cosmic radiation” (for historical reasons it is called “cosmic radiation” or “cosmic rays” because it was initially thought to be electromagnetic radiation; now we know that the main component is protons) and no less than three prizes awarded in the 21st century (2002, 2006, and 2011). Whereas the Universe has provided the earth-bound physicists with discoveries of sufficient importance to motivate a Nobel Prize, it has failed to do likewise to the chemists. Nevertheless, the serendipitous discovery of fullerenes was arguably inspired 1 by observations by means of radio astronomy of interstellar carbon chain molecules,2−4 although the opposite view has been advanced by one of the discoverers.5 Although a few interstellar molecules (CH, CN, CH+) had been observed by absorption spectroscopy in the visible wavelength region in the diffuse interstellar medium already during the late 1930s and early 1940s (see ref 6 for a detailed history of the early work), it was the radio observations of the OH radical7 and, in particular, ammonia8 and water9 in dark interstellar clouds that represent the birth of astrochemistry. In 1971, already 20 interstellar molecules had been discovered.10 How are they formed? While searching for interstellar HCN, Buhl and Snyder11 serendipitously discovered a radio emission line at 89.190 GHz, which they were unable to assign to any known molecule; hence, they named the carrier of the line X-ogen. Klemperer12 speculated boldly that the carrier could be HCO+; protonated molecules were at that time only known in mass spectrometry, and there was no observation of a microwave, infrared, or visible spectrum from a protonated molecule. His speculation was eventually proved correct by laboratory microwave spectroscopy.13 Observation of a protonated ion in the very cold and thin interstellar medium suggested that ion−molecule reactions ought to play an important role in molecule formation, and models in that direction were developed14−17 even before the carrier of the 89.190 GHz line had been confirmed in the

CONTENTS 1. Introduction 2. Importance of Ion Reactions in Astronmical Objects 2.1. Chemistry of the Early Universe 2.2. Diffuse Clouds 2.3. Dark Clouds 2.4. Prestellar Cores and Protostars 2.5. Protoplanetary Disks 2.6. Circumstellar Envelopes 2.7. Planetary Atmospheres and Ionospheres 2.7.1. Rocky Planets 2.7.2. Gaseous Planets 2.7.3. Satellites of Gaseous Planets 2.8. Conclusion 3. Experimental Challenges To Investigate Ion Reactions 4. Ionic Reactions 4.1. Ion−Electron Reactions 4.2. Ion−Neutral Reactions 4.3. Ion−Ion Reactions 5. Conclusions Associated Content Supporting Information Author Information Corresponding Author Notes Biographies Acknowledgments References

A B B C D F H J J J L M O P Q Q W Y AA AA AA AB AB AB AB AB AB

Special Issue: 2013 Astrochemistry Received: May 10, 2013

© XXXX American Chemical Society

A

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laboratory.13 It was realized that H3+ must play an important role as the initiator of chains of reactions leading to molecule formation. However, in the early 1970s, H3+ was only known in mass spectrometry. Observation of a laboratory infrared spectrum of H3+ 18 and interstellar H3+ about 15 years later19 were key steps in establishing the importance of ion−molecule reactions in the chemistry of the interstellar medium. Presently, about 150 interstellar molecules have been observed, and the most extensive gas-phase reaction networks include 4000 reactions. The last 5 years have witnessed two remarkable developments in ion astrochemistry. First detection of an interstellar anion, C6H−, was reported in late 2006,20 and the list of anions now includes hydrogenated carbon chain anions, C4H−,21 C6H−,20 and C8H−,22,23 and three cyano anions, CN−,24 C3N−,25 and, tentatively, C5N−.26 The Herschel Space Observatory has opened the far-infrared and submillimeter wavebands (55−671 μm) to astronomical observations. This has resulted in observations of new interstellar ions, such as H2Cl+27 and HCl+,28 and observation of OH+, H2O+ and H3O+ in one and the same source, thus making it possible to strengthen our understanding of gas-phase formation of interstellar water.29 In order to understand the formation and destruction of molecules in the interstellar medium and other astrophysical sources, chemical modeling based on laboratory data on atomic and molecular processes is required. This review concerns experimental investigation of ionic reactions of astrophysical importance. More often than not this necessitates development and application of experimental methods that are not main stream in chemistry. There are several recent review articles that treat similar topics but with the focus more on the astrophysical questions, and the reader is referred to these for a comprehensive discussion of the astrophysics.6,30−32

that chemistry, including ion chemistry, played an only very limited role in the early universe. In addition, three-body processes were excluded due to the low particle density, and of course, there were no solid bodies like grains to enable surface processes, which further restricted the number of possible chemical reactions in the early universe. Thus, the range of feasible processes was fairly limited. Nevertheless, radiative association reactions

2.1. Chemistry of the Early Universe

During the first nucleosynthesis that took place about 200 s after the Big Bang33 only 3 elements were created from the protons and neutrons, which in turn formed through combination of quarks in the first second of the history of the universe: deuterium, lithium, and 2 unstable isotopes of beryllium (7Be and 8Be). At this time, the chemical elements merely existed as fully stripped ions.33 During the recombination period that started from 390 000 years after Big Bang these bare atomic nuclei associated themselves through radiative recombination with electrons, e.g.

He + H+ → HeH+ + hν

(4)

H + H+ → H 2+ + hν

(5)

H 2 + + H → H 2 + H+

(6)

+

At a later age, the HeH ions created in reaction 4 take an active part in formation of molecular hydrogen, forming H2+ ions H + HeH+ → H 2+ + He

(7)

These ions can then further react according to reaction 6 to form H2. At a red shift lower than 300, reaction 7 is the dominating destruction pathway of HeH+.38 At a red shift of z ≈ 100, radiative attachment of hydrogen atoms to H− becomes important.

(1)

H + e− → H− + hν

At this time, the universe had cooled down to around 4000 K and no photons with sufficient energy were available for the reverse process of reaction 1, namely, photoionization of the resulting neutral atoms H + hν → H+ + e−

(3)

created the first molecular ions at a red shift of z ≈ 2000.33,35 The red shift (z) is defined as the ratio between the observed wavelength and the emitted wavelength minus 1 (λobs/λemitted − 1). Scientists working in the field of early universe physics and chemistry prefer this unit, since it changes rapidly in the first fractions of seconds after Big Bang (where many important processes took place) and then alters more slowly, where things happen on a longer time scale. Ion chemistry also plays a large role in formation of H2, which is crucial for cooling of the denser clumps formed in the primeval clouds, which were formation sites of the first stars in the early universe. Without any cooling process, the temperature (and, concomitantly, the pressure) of the clumps would increase during collapse through conversion of kinetic into thermal energy until the collapse stops and star formation is halted.36 Cooling can happen through rotational excitation and subsequent radiative deactivation of the H2 molecules, allowing these clumps to continue to collapse to form the first massive (Population III) stars (atomic hydrogen is a very inefficient cooling agent below 8000 K, the temperature corresponding to the Lyman alpha line37).33 Furthermore, formation of H2 enabled the first, galaxy-mass objects to form in an expanding universe, in which Population III stars formed. At very high red shifts, the following reaction sequence is the main source of H2

2. IMPORTANCE OF ION REACTIONS IN ASTRONMICAL OBJECTS

H+ + e− → H + hν

He + He+ → He2+ + hν

(8)

At this age the photon energy in the universe becomes less than the photodetachment threshold of H−; thus, this anion can efficiently undergo associative attachment to form H2 H + H − → H 2 + e−

(2)

Nevertheless, some electrons and ions still remained after the recombination period. These could react with the present neutrals (reactions between different cations are hampered by the Coulomb repulsion between the reactants34). Given the paucity of the present elements, one might be forgiven to think

(9)

At red shifts less than 100, reaction 9 becomes the dominating formation process of molecular hydrogen. However, several destruction processes like dissociative recombination of HeH+ and mutual neutralization of H− can successfully compete with H2 formation B

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Figure 1. Predicted abundances of molecules in the early universe. Reproduced by permission of the AAS from ref 42.

HeH+ + e− → He + H

(10)

H+ + H− → 2H

(11)

Several types of clouds are found in the ISM: diffuse, dense, and intermediate translucent clouds. Diffuse clouds display a particle density of around 100 cm−3 and a temperature of 100− 200 K. Since they are transparent to UV light, both cosmic ray and UV photoionization are important in these objects. In diffuse clouds, cosmic rays and UV photons initiate the ionization of H2

although reaction 10 has been found to proceed comparatively slowly.39 Nevertheless, it played a fundamental role in the overall neutralization of the early universe.33 Rate constants of primeval ionic processes have been found to have great influence on the abundance of molecular hydrogen40 and thus the probability of star formation in the early universe. Moreover, lithium compounds take an active part in the chemistry there,41 and respective chemical models can include more than 80 reactions.37 Many of these involve ions, and their rates are not always well characterized and subject to large uncertainties. Even isotopes can play a role: The near-resonant charge-transfer process D + H+ → D+ + H

(14a)

H 2 + cosmic ray → H+ + H + e−

(14b)

+

H , which is abundant in diffuse clouds, is also formed by ionization of H atoms by cosmic rays. H2+ ions generated by reaction 14a or UV photoionization rapidly react with H2 to form H3+ H 2 + H 2+ → H3+ + H

(12)

In diffuse clouds, recombination

followed by the exothermic reaction D+ + H 2 → H+ + HD

H 2 + cosmic ray → H 2+ + e−

(13)

is a major source of HD,42 which, due to its small dipole moment, is an even better cooling agent than H2. Thus, HD could enhance the collapse of clumps, and reaction rates of deuteration processes like reaction 13 can be crucial for star formation in the primeval universe.37,43 The abundances of different molecules predicted by a standard chemical model are shown in Figure 1, and the model results clearly indicate that a huge variety of molecular processes take place in the early universe. Thus, it can be said that chemistry has a very long history.

H3+

(15)

are mainly destroyed by dissociative

H3+ + e− → H 2 + H

(16a)

H3+ + e− → 3H

(16b)

which makes the H3+ to H2 abundance ratio a gauge for the cosmic-ray ionization rate ζ.45 The predicted value for ζ directly depends on the rate constant for dissociative recombination of H3+, which was very controversial for decades and is still debated.6,46 Although diffuse clouds are not a favorable environment for ion−neutral reactions due to the low density prevailing in these objects, H3+ can, due to the low proton affinity of H2, act as proton donor to common neutral molecules like CO

2.2. Diffuse Clouds

Since the first (Population III) stars were very massive and their nuclear processes very fast, they burned out very quickly (on the time scale of several million years). The intense UV light emitted by them led to a reionization of the universe. Finally, they exploded as gigantic supernova ejecting a great number of heavy elements, dust, and even molecules like CO and SiO.44 These ejecta can then form molecular clouds.

H3+ + CO → HCO+ + H

(17)

forming ions that can undergo subsequent reactions. Also, as can be seen below (reactions 19−25), oxygen and carbon ions created through charge transfer H+ + O → O+ + H C

(18)

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HCl+ + H 2 → H 2Cl+ + H

can start chains of reactions leading to important interstellar species. Interestingly, the velocity of reaction 18 is dependent on the spin−orbit state of the oxygen atom; while it is fast for the 3P1 level, it proceeds slowly for the ground 3P2 state.47 In diffuse clouds ion reactions lead to the important molecules CO, OH, and H2O O+ + H 2 → HO+ + H

(19)

OH+ + H 2 → H 2O+ + H

(20)

H 2O+ + H 2 → H3O+ + H

(21)

H3O+ + e− → H 2O + H

(22a)

→ OH + 2H

(22b)

→ OH + H 2

(22c)

→ O + H + H2

(22d)

C+ + OH → CO+ + H

(23)

CO+ + H 2 → HCO+ + H

(24)

HCO+ + e− → CO + H

Since H2Cl does not react with H2 and H, it is mainly destroyed by dissociative recombination

+

+

(28a)

→ Cl + 2H

(28b)

→ Cl + H 2

(28c)

C+ + H 2 → CH+ + H

Interestingly, intermediates OH , H2O , and H3O have recently been observed in diffuse interstellar gas using the Herschel Space telescope29,48 (see also Figure 2). Also, it has

(29)

is endoergic by about 0.4 eV and only proceeds with a rate constant of k = 1.0 × 10−10 exp(−4640/T) cm3 s−1.52 Thus, it is unfeasible in dark and diffuse clouds. It has been suggested that CH+ is formed in higher velocity regions present in shock zones53,54 or caused by Alfvén waves52 and turbulent dissipation.55 However, formation by shocks is questioned by the fact that similar radial velocities have been observed for CH+ and the corresponding neutral CH, which is at odds with such a scenario,56 so Alfvén waves and turbulent dissipation remain as possible sources. To throw some light on CH+ formation in diffuse clouds, Indriolo et al.57 attempted to gauge the temperature and molecular hydrogen fraction of a diffuse cloud by determining the CH3+/CH+ ratio along the diffuse cloud sight line toward Cyg OB2 No. 12. However, these efforts were thwarted by nondetection of CH3+ in this object, so it could only be concluded that CH+ (as the low rate constant at low temperature implies) is not formed under equilibrium conditions prevalent in normal diffuse clouds. The exact formation pathway of CH+, however, still remains a mystery.

Figure 2. Spectra of H2O+ 111−000 (green) and OH+ N = 1−0 (blue) transitions obtained toward W49N. Credit: Neufeld et al. Astron. Astrophys. 2010, 521, L10 (ref 301), reproduced with permission © ESO.

2.3. Dark Clouds

been proposed that OH could act as a chemical thermometer in diffuse clouds. Once this molecule has been generated, it forms CO through reactions 23−25. It is evident that the branching ratios between the different pathways of the dissociative recombination (reaction 22) have a pivotal influence on the abundance of OH and H2O. Experimental investigations using an ion storage ring yielded that dissociative recombination of H3O+ leads to water with a probability of 25% (pathway 22a) and to OH with one of 60% (pathways 22b and 22c).49 Model calculations yielded that whereas CO is the most abundant molecule in present dark clouds, in which higher metallicities (abundances of elements heavier than helium50) are prevalent, OH should dominate in the ones present during the reionization epoch, whose metallicity was much lower.51 A build-up similar to H2O has been invoked for, e.g., the interstellar molecule HCl Cl+ + H 2 → HCl+ + H

H 2Cl+ + e− → HCl + H

This mechanism is particularly enhanced by the fact that chlorine is mainly present in its cationic form in the interstellar medium, since the ionization potential (13.0 eV) lies below that of atomic hydrogen (13.6 eV). The feasibility of formation of HCl via reactions 26−28 has been corroborated through detection of the intermediates HCl+ (even with higher abundances than predicted by chemical models)28 and H2Cl+ by Herschel toward Galactic star-forming regions.27 However, there are also problems associated with formation pathways of interstellar molecules through ionic processes. If one wants to invoke a build-up of hydrocarbons (e.g., methane) from carbon ions via a pathway similar to water formation (reactions 19−22), one faces the problem that the process analogue to reaction 19

(25) +

(27)

+

Diffuse clouds can condense and further evolve into dark clouds, of which an example is shown in Figure 3. During this process the chemical composition of these objects changes substantially. Atomic hydrogen is turned into molecular hydrogen by reactions on grain surfaces, and atomic carbon and CO become the main carbon reservoirs (as opposed to C+ in diffuse clouds).58,59 Thus, the term ‘molecular clouds’ is often used for them. They are very productive chemical factories; a large percentage of the hitherto identified interstellar molecules has been detected in dark clouds.60 Temperatures of gas-phase molecules and ions in dark clouds lie mostly in the range 7−20 K.61 Complex turbulent motions give dark clouds a very complicated velocity structure,62 and the magnetic field of the cloud influences the motion of charged particles such as ions and electrons. Dark clouds are not homogeneous. First, their density increases toward their center, and second, they contain clumps of denser material surrounded

(26) D

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employed to this end.68 However, studying both ratios simultaneously can be difficult, since they are not independent of each other.69 More recently, the ratio between C6H− and C6H has also been used for assessment of the electron abundance (fractionation).70 The electron density can also be affected greatly by formation of anions. In particular, polyaromatic hydrocarbons (PAHs) can play an important role in this respect due to their high electron affinity and propensity to undergo radiative attachment. They therefore can effectively reduce the number of electrons in dark clouds.33,70 Also, interstellar grains can efficiently scavenge electrons. Due to the abundance of molecular hydrogen in dark clouds, ionization of H2 with subsequent reaction of the H2+ ion with H2 (reactions 14a, 14b, and 15) is the main ion formation pathway. As mentioned, H3+ can easily protonate other molecules like CO (reaction 17), N2, and H2O

Figure 3. Picture of the dark cloud B68 in the visible (left) and infrared (right). As can be seen, the scattering of the light is less in the infrared, so that the dark cloud becomes translucent at infrared wavelengths. Reprinted with permission from the European Southern Observatory.

+

(30)

(31)

+

in which kDR(H2D ) is the rate constant of the dissociative recombination of H2D+ and kx the rate constant for protonation of different molecules X (e.g., N2, CO, H2O) with H3+. kf and kb are the forward and backward reactions of the deuterium exchange reaction +

+

H3 + HD → H 2D + H 2

(34)

(35)

are involved as first steps of molecule formation in dark clouds.71 However, the latter process has a small barrier of approximately 7 meV72 (such small thresholds are not uncommon for H-abstraction reactions); thus, its rate constant is quite low at dark cloud temperatures. o-H2 (J = 1) is, however, sufficiently high in energy to power this reaction, so that its rate is determined by the o/p-H2 ratio. Nevertheless, relative velocities between ions and neutrals of up to 5 km s−1 have been observed in magnetohydrodynamic shocks, which would enable ion−neutral reactions even if they possess small energy barriers.73,74 As HCl+ and OH+, NH+ can further react with H2 under subsequent hydrogen abstraction

in which n(DCO ) is the number density of DCO and so on. The factor f is defined as −1 n(X ) ⎤ ⎥ ∑ kx n(H 2) ⎥⎦ X

H3+ + N2 → N2H+ + H 2

N+ + H 2 → NH+ + H

+

⎡ k (H D+)ne f = k f ⎢k b + DR 2 + ⎢⎣ n(H 2)

(33)

which leads to protonated species that are important for observational astronomy. For example, N2H+ is used as a tracer for N2 in radioastronomy, since N2 lacks a dipole moment and thus cannot be observed through radiotelescopes. Ion reactions are of paramount importance also for another reason. At the dark cloud formation stage, H2 is probably the only important neutral molecule in the gas phase. Unfortunately, no atoms (apart from fluorine) can undergo exoergic, barrierless reactions with H2 to form a diatomic molecule HX. Thus, ion reactions such as eqs 19, 26, or

by more diffuse gas. Production of ions in different parts of the clouds proceeds along different pathways depending on the density (opacity) of the region. In the dense inner parts of dark clouds, cosmic rays are the main source of ionization, since they can uniquely penetrate into these environments. In the outer regions with lower opacity (density) UV photoionization by the interstellar radiation field is the dominating ionization process.63 Nevertheless, the inhomogeneous, clumpy structure of dark clouds might make it possible for some UV radiation to penetrate into their inner parts via less dense regions.64 The overall cosmic-ray flux in our galaxy is difficult to determine since low-energy cosmic-ray particles (100 K) and densities (106−108 cm−3) occur in these objects. Hot cores consist of prestellar material not consumed by the star96 and contain short-lived complex species.87 Enhanced abundances of polyatomic molecules like methanol, formaldehyde, ammonia, dimethyl ether, and methylformate are found in hot cores,91,97−99 which suggests a very evolved chemistry. This

Figure 4. Different phases of star formation. Reprinted with permission from ref 91. Copyright 1998 Annual Reviews.

X-ray fields, and thus, these species can serve as excellent tracers for X-ray-dominated regions. Such tracers are important in distinguishing between X-ray- and UV-induced chemistry near YSOs, since X-rays are not directly observable if the young star is still embedded.92

Figure 5. Schematic of a star-forming region. Reprinted with permission from ref 87. Copyright 2005 Cambridge University Press. G

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molecular abundance also implies a reduction of primary, cosmic-ray-produced ions such as H3+, He+, and C+, since they can readily react with molecules evaporated from grains.93 Whereas reactions of these species with He+ are highly exoergic and expected to lead to molecular fragmentation, reactions with C+, which is less reactive, often imply charge transfer.100 As in dark clouds, ion−neutral reactions leading to protonated molecules followed by dissociative recombination were previously held responsible for forming the complex organic compounds like alcohols, esters, and ethers encountered in hot cores. However, recent experiments have demonstrated that protonated alcohols and ethers form the unprotonated substances with only small branching ratios upon dissociative recombination.101,102 In the case of methyl formate, it was even predicted by ab initio calculations that the protonated precursor molecules cannot be formed at hot core temperatures.103 Model calculations suggest that methylformate is formed by surface processes, whereas dimethyl ether and formic acid can be formed by ion reactions in the gas phase.104 This shows that a universal catch-all mechanism for forming molecules in star-forming regions does not exist, and generation of each species has to be explored individually.

produced by them also has to be taken into account. Also, the decay of radioactive nuclides such as 26Al and 60Fe initiates ionization.108 On the whole, ionization can diminish molecular abundances in protoplanetary disks by up to a factor of 500 compared with the level in dark clouds.105 The relative dominance of the different ionization processes is dependent on the luminosity of the central star. Models indicate that in a protoplanetary disk surrounding a star with a luminosity of 1028 erg s−1 ionization is mainly dominated by cosmic rays throughout the disk. With increasing luminosity, the difference between the midplane (dominated by cosmic rays) and surface ionization rates (governed by UV and X-ray flux) rises dramatically,107 as the UV field from the central star can exceed the interstellar radiation field by 4 orders of magnitude. Closest to the star, X-ray ionization dominates and ion fractions decrease with the distance from the central star, since the X-ray intensity drops faster with the radius than the density of the disk.109 In the inner disk, observed ion abundances consequently can amount to 109 cm−3, while the lower limit of the global ionization fraction in the disk was deduced to be 2 × 10−10.110 Figure 6 shows protoplanetary disks in the Orion Nebulae.

2.5. Protoplanetary Disks

Young stars can be surrounded by protoplanetary disks that can subsequently evolve into planetary systems. Protoplanetary disks are short-lived features, with lifetimes of 300 000−1 000 000 years, consisting mostly of gas and dust. They are rotating, and their emission extends up to 10 and 1000 au from the star. In these disks, dust grains coagulate into planetesimals, which further assemble into planets, comets, and asteroids. Ionization of atoms and molecules present in the disk by cosmic radiation and light from the star initiates a multitude of chemical processes, making protoplanetary disks very interesting objects.105 Formation of ions also significantly affects the accretion of the disk, so the ion fraction has a major influence on the evolution of such objects into stars and planetary systems.106 Ions also play a particularly pivotal role in protoplanetary disks by coupling the gas with the magnetic field. The rich chemistry in protoplanetary disks subsequently influences the composition of planets and their atmospheres as well as asteroids and comets.107 Radially, protoplanetary disks can be divided into a warm inner zone (the site of planet formation) and a cold outer zone.105 Perpendicular to the rotational axis, protoplanetary disks can be separated into three zones at radii greater than ∼100 au: the dense midplane (where molecules are frozen out on grains), the surface (where chemistry is photon dominated), and the intermediate (warm molecular) zone between these two layers. Planets are finally formed in the disk midplane, where, at large radii, most molecules are frozen out on grains and planetesimals. With increasing axial distance from the midplane, the density decreases. UV radiation from the star and the interstellar radiation field can easily penetrate the surface zone and photodissociate and photoionize most molecules. Gas-phase molecular abundances thus tend to peak in the intermediate zone, where molecules are desorbed from the grains but not yet photolyzed. In the midplane ion production is triggered mainly by cosmic rays. With increasing axial distance from the plane, X-rays and UV photons become the main ionization sources. Due to the high density of protoplanetary disks, cosmic rays usually are absorbed in the disk and ionization from secondary particles

Figure 6. Protoplanetary disks in the Orion Nebulae take with the Jet Propulsion Laboratory’s Wide Field Planetary Camera 2 on the Hubble Space Telescope (NASA Credit PRC95-45b).

Models of protoplanetary disks predict that in the layer nearest the disk surfac, C+ is the dominant ion. In the intermediate plane, metal ions, HCO+, and HCNH+ are prevalent in the warmer region (at radii below 3 au). HCO+, which is formed by protonation of CO by reaction 17, dominates in outer regions (R > 30 au) at a later stage of the evolution of protoplanetary disks. At intermediate radii, He+ initiates an elaborate chemistry by forming C+ and S+ ions through charge transfer, leading to formation of carbon chains. H

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H3+ and its isotopomers trace the midplane of the disk (see Figure 7), apart from very close to the star where NO+

neutrals, e.g., carbon chain molecules (HCn), which are thought to be formed by reaction of CnH2+ ions with electrons113 CnH 2+ + e− → CnH + H

(48)

In the case where dissociative recombination is the only degradation route for ions in protoplanetary disks and cosmic radiation their only formation route, ion fractionation, xe (nion/ nH), can be expressed as xe(eq) =

ς kDR nH

(49)

if the disk is assumed to be in chemical equilibrium. In this expression, ζ is the cosmic-ray ionization rate, kDR the average dissociative recombination rate constant, and nH the hydrogen number density. This expression, however, applies only if chemical processes occurring on dust grains are excluded. However, such grains and PAHs play an essential role in protoplanetary disk chemistry. Some PAHs can even survive the strong radiation field of the young star during planet formation and end up on planet surfaces.114 Charged grains can serve as recombination sites, and the probability of ion recombination products sticking to grains also substantially affects the cation abundance predictions of the models.109 Also, ions of PAHs greatly influence the chemistry of protoplanetary disks. Ionization of PAHs by UV, cosmic rays, but also through charge transfer with C+ ions leads to formation of fast electrons that cause gas heating of these objects. PAH cations can undergo dissociative recombination reactions, and the rates of such processes have been determined experimentally at room temperature.115 Due to their ability to undergo vibrational relaxation of highly excited states, neutral PAHs are also thought to efficiently attach electrons; however, experimental data on such processes are still scarce. This is somewhat unfortunate, since sensitivity analyses indicate that ionization and electron attachment rates significantly affect the outcome of model calculations of protoplanetary disks, whereas neutral processes involving PAHs have less impact. Temperature also exerts a strong influence, since the rate of dissociative

Figure 7. Schematics of a protoplanetary disk. Reproduced by permission of the AAS from ref 387.

dominates due to its chemical inertness.109 At larger radii, N2H+ and HCO+ also become important, although the former ion is destroyed quickly by the stellar radiation near the disk surface. Overall, HCO+ is the most abundant ion in protoplanetary disks.110 Grains can be important charge carriers in the midplane; they can be both positively and negatively charged, and even multiple charges might be present on grains.111 As in star-forming regions, deuteration by ions, H2D+ and CH2D+, augments the deuterium fractionation of molecules and ions in protoplanetary disks,104 and DCO+ has been detected in these objects.112 In the inner zone, the density is high enough that three-body reactions can occur. In other regions, chemistry is restricted to two-body processes. Ion−neutral reactions dominate the chemistry of protoplanetary disks, and dissociative recombination is the primary destruction process. Dissociative recombination also serves as the main formation pathway for some

Figure 8. Chemical structure of a circumstellar envelope. Credit: Decin et al. Astron. Astrophys. 2010, 516, A69 (ref 388), reproduced with permission © ESO. I

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MgNC5H+ + e− → C4H + MgCN

recombination usually diminishes by an exponential factor of between 0.5 and 1, whereas electron attachment to neutrals is less temperature dependent. Thus, higher temperatures lead to somewhat greater abundances of cations, but this dependence is strongly influenced by the size of the above-mentioned exponential factor. Again, exact knowledge of the behavior of rate constants of ion reactions is crucial for the understanding of ion chemistry in astronomical objects.

This scheme is corroborated by the observation that MgNC and C4H have nearly identical distributions in IRC+10216.132 Many species found in the molecular envelope of IRC+10216 are carbon-chain radicals, such as C5H, C6H, H2C3, H2C4, and C8H,133 which have formed in the external layers of the CSE. In the innermost part of the envelope, dominated by thermodynamic equilibrium,134 many stable molecules are found, for example, CO, HCN, C2H2, SiO, CS, and SiS. At large radii, photodissociation by the interstellar radiation field destroys all molecules apart from molecular hydrogen.125 In oxygen-rich circumstellar envelopes like VY Canis Majoris HCO+, SO2, and even NaCl have been detected, which also points to an evolved chemistry in these objects. Generally, oxygen-rich circumstellar envelopes are chemically less complex than IRC+10216, which is also illustrated by the fact that they do not show as many unidentified lines in radioastronomy observations. In the inner part near the star, molecules are shielded from interstellar UV photons by the circumstellar envelope. In this region the temperature is high and the chemistry rich and dominated by neutral−neutral reactions leading to closed-shell species. Molecular species are easily formed under these conditions and can interact with dust grains,135,136 which are pushed outward together with gas under the radiation pressure.137 Also, the grains moving out exert a drag force on the gas of the envelope. In the outer regions of the circumstellar envelope, the interstellar UV radiation field causes photodissociation and photoionization, allowing a different type of chemistry. At a later stage, circumstellar envelopes evolve into protoplanetary and further into planetary nebulae with a hot, exposed central star. In these objects, synthesis of molecular ions becomes important and their chemistry is governed by radical- and ion-induced reactions.138 Numerous molecules like HCO+, HCN, and CCH are found in planetary nebulae such as the Helix, and even CH+ has been detected in these objects.139 They are then slowly expelled into the interstellar medium and may be linked to the variety of compounds found recently in diffuse clouds, thus closing the circle of chemistry during the life cycle of a star.121 In fact, the material ejected into the ISM by circumstellar envelopes may account for nearly 80% (by mass) of the ISM.140 The “recycling” of material from circumstellar envelopes and protoplanetary nebulae might also partly explain the large presence of complex species like carbon-chain molecules in dark clouds.121 As we pointed out in previous sections, reactions involving ions play a crucial role during the whole life cycle of stars: however, it would be by far exceeding the scope of this review to discuss all participating processes. We thus refer the reader to the excellent review on the ion chemistry in the interstellar medium by Snow and Bierbaum.31

2.6. Circumstellar Envelopes

Circumstellar envelopes are formed from the outflow of carbon- and oxygen-rich stars in their late stage of evolution, like asymptotic giant branch stars and red supergiants. They are generated by thermal pulses in the star interior and radiation pressure on dust116,117 and are roughly spherical in shape and not gravitationally bound to the central star. Circumstellar envelopes are quite large, extending to 104−105 stellar radii, see Figure 8. Near the stellar photosphere, the envelope material (all gaseous) is warm (T = 1500 K) and dense (n = 1010 cm−3). With increasing distance from the star temperature decreases with a factor of 1/r (although the factor is larger close to the star118) and the density with one of 1/r2.119,120 Near the outer edges of the envelope, the temperature has sunk to 25 K and the density to 105 cm−3. Circumstellar envelopes harbor a complex chemical environment involving hot, thermodynamically controlled synthesis, freeze out of molecules on grains, shock-initiated reactions, and photochemistry governed by radical mechanisms. More than 60 different chemical compounds have been identified in the well-studied circumstellar envelope of the carbon-rich star IRC+10216.121 In this object also the first interstellar anions have been detected,20 and anion chemistry could play an important role in circumstellar envelopes.122 This is illustrated by the fact that 5 anions (C4H−, C6H−, C8H−, CN−, and C3N−) have been observed in IRC +10216.123,124 These ions can be produced by radiative attachment, but also anion−neutral reactions might play a significant role125,126 C2 H 2 + H− → C2 H + H 2

(50)

C7− + N → CN− + C6

(51) −

In circumstellar envelopes, H is predicted to form mainly by the cosmic-ray-induced ion pair formation H 2 + cosmic ray → H− + H+

(52)

127

albeit the process is fairly slow. Anions are mostly destroyed by mutual neutralization (with C+), associative detachment (with H atoms), and photodetachment.125 Many associative detachment and photodetachment processes have been studied experimentally,128,129 but data on mutual neutralization reactions is still rather sparse, albeit such processes are thought to be important in the interstellar medium.130 This is unfortunate since models overestimate the abundance of HC4− in IRC+10216,125 and consequently, efficient destruction mechanisms for this ion are to be found. In carbon-rich envelopes the chemistry is dominated by species containing long carbon chains and metal compounds (mostly metal halides and cyanides), which makes it quite different from the one encountered in dark clouds. Reactions involving ions are held responsible for formation of metal cyanides,131 e.g. Mg + + HC5N → MgNC5H+ + hν

(54)

2.7. Planetary Atmospheres and Ionospheres

2.7.1. Rocky Planets. The ion chemistry around the 4 rocky planets of the solar system (Mercury, Venus, Earth, and Mars) differs greatly because of the dissimilarity of the composition of their atmospheres. Since a discussion of the ion chemistry of the terrestrial atmosphere would by far exceed the limits of this review, we confine us to the atmospheres of the three other planets. Mercury has only a thin atmosphere (surface pressure 10−15 bar) consisting mainly of atomic oxygen, hydrogen, helium, and traces of metals (calcium, magnesium, sodium, and potassium).

(53) J

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Due to its sparseness it has often been called a “surface-bound exosphere”.141 Mercury also has a relatively strong magnetic field which deflects the solar wind particles during very strong solar wind conditions. Nevertheless, ion chemistry is not unimportant there. In phases of strong solar wind, ion sputtering caused by its particles could be a main source of exospheric material.141 This material can be photoionized by solar UV photons, and the emerging ions can be trapped in Mercury’s magnetotail and transported to the night side of the planet.142 The transported ions can then impinge on the night side of the planet,143 where they are neutralized or chemically bonded144 and lead to a net transport of atmospheric material from the day side to the night side. Furthermore, photoionization is also one of the most important sinks of neutrals in the Mercurian atmosphere. Also, negative ions can be on surfaces of rocky planets. In the case of planets (and satellites) without an atmosphere anion formation through ejection from the surfaces of planets can be important also. Three mechanisms are held responsible: charge inversion by proton scattering, secondary negative-ion emission, and micrometorite impact. The latter process is most important for Mercury, which possesses a magnetic field deflecting the solar wind and thus hinders ionization by surface impact of solar wind particles. Photodetachment is the most important anion degradation route because of the low density in Mercury’s surroundings caused by the lack of a dense atmosphere.145 The atmosphere of Venus is very dense with a surface pressure of ∼90 bar and a surface temperature of 740 K. It is dominated by carbon dioxide (95%) and nitrogen (3.5%). Trace substances include sulfur dioxide (200 ppm), argon (70 ppm), water (30 ppm), carbon monoxide (17 ppm), helium (12 ppm), and neon (7 ppm).146 Also, tiny amounts of hydrogen chloride and hydrogen fluoride are present. Venus has no a magnetic field, and the UV radiation hitting its atmosphere creates a distinctive ionosphere. Interestingly, interaction between the Venusian atmosphere and solar wind plays only a minor role.147 However, the Venus Express and mission detected ionospheric layers created by X-rays at an altitude of 125 km above Venus.148 In higher layers of the atmosphere of Venus, ionization of the dominant atmospheric component carbon dioxides through EUV radiation leads mainly to O2+ and O+. Photoionization can even result in formation of the dications. In the diurnal Venusian atmosphere, the main processes of dication generation are photoionization and electron-impact ionization. The exact mechanisms of formation of these ions have not yet been fully elucidated, but double valence shell ionization and Auger ionization have been held responsible. At an altitude below 60 km, where solar UV and X-ray photons cannot penetrate and on the night side of the planet, cosmic rays become the main source of ions in the Venusian atmosphere.149 Cosmic-ray ionization mainly leads to O2+, CO2+, and CO+.150 Since molecular oxygen (which can be formed through dissociative recombination of CO2+) reacts with CO2+ to produce O2+ on Venus, an important fraction of the CO2+ ions produced is transformed into O2+ in the Venusian atmosphere.151 Other possible ion sources are lightning152 and impacts by micrometeorites.153 Since the (retrograde) rotation of Venus is very slow (a Venusian day equals 58 Earth days), transportation of ions from the day side to the night side is a crucial factor in the total ion balance.154

Dissociative recombination is one of the most important decay pathway of ions in the atmospheres of Venus. Moreover, dissociative recombination of CO2+ leads to the excited CO molecules that decay under the emission of the fourth positive (A 1P → X 1S) and Cameron (a 3P → X 1S) bands.155 The green nightglow on Venus could be due to excited oxygen atoms produced by dissociative recombination of O2+.156 Dissociative recombination of oxygen-containing ions is also an important source of suprathermal oxygen atoms.157 Also, charge-transfer reactions between ions, like O2+ + NO → NO+ + O2

(55)

can be important. Reaction 55 is the main source of the stable NO+ ion on the night side of Venus.153 Charge-transfer reactions can also lead to fast (“hot”) H atoms (H*) O2+ + H 2 → OH+ + H*

(56)

which can escape the gravitational field of the planet. Such hot atoms are thought to account for a considerable part of the total hydrogen escape from the planet’s atmosphere.158 On heavy planets like Earth and Venus, however, part of these hot atoms will fall down again under the influence of gravity, forming “hot atom” coronae.159 Regarding ion chemistry, the Venusian atmosphere has some peculiarities due to its density and absence of a magnetic field: Loss of ions via pick up from the solar wind plays a dominant role in hydrogen and oxygen escape and, consequently, loss of water from the Venusian atmosphere.157,160 The high density of the atmosphere fosters clustering of ions, mainly with water and carbon dioxide.161 Also, Mars has a CO2-dominated atmosphere (mole fraction 95%), which is, however, much thinner than Venus’s with a surface pressure of only 60 mbar. Main minor components are nitrogen (2.7%), argon (1.6%), oxygen (0.13%), carbon monoxide (0.07%), and water vapor (0.03%). Photoionization and cosmic-ray ionization are the most important primary ion production mechanisms in the atmosphere of Mars. Their contribution differs with altitude. Below 80 km, cosmic rays are the main source of ionization. Above this level, solar EUV photons and X-rays are mainly responsible for producing CO2+, N2+, O+, CO+, and O2+ ions in the atmosphere of Mars. CO2+ can also be efficiently produced through charge transfer between neutral carbon dioxide and N2+, Ar+, and CO+. Ionization by solar EUV photons peaks at an altitude of 125− 145 km (F1 peak), X-ray ionization at 100−112 km (E peak), and cosmic-ray ionization at 25−35 km.162,163 As on Venus, the dominance of CO2 in the atmosphere also manifests itself in the products of EUV radiation ionization, namely, O2+ in the lower and O+ in the upper atmosphere. Also, the existence of dications like CO22+ has been predicted by models.164,165 In the Martian ionosphere, dissociative recombination is one of the most important decay processes of ions. In the case of CO2+, it has been regarded to produce the majority of carbon atoms in the atmosphere of Mars, since dissociative recombination of this ion was determined to lead to atomic carbon with a branching fraction of 9%.166 However, these findings have been disproved in a later experiment, which yielded that O and CO are the only products of this process.167 Dissociative recombination can also lead to isotope fractionation in planetary atmospheres. Such fractionation of nitrogen has been observed in the atmospheres of Mars and Titan, where 15N/14N ratios are substantially larger than in Earth’s atmosphere.168,169 Due to the law of conservation of K

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momentum, the 14N product receives a slightly larger share of the kinetic energy release in the dissociative recombination of 14 15 + N N 14

N15 N+ + e− → 14 N + 15 N

In the case of Saturn, the picture is slightly different: Reaction of H+ with water from the Saturnian rings can act as a destruction mechanism in the Kronian atmosphere

(57)

This kinetic energy difference can cause preferential escape of the lighter 14N isotope and, consequently, the observed enrichment of 15N in the atmosphere of Mars.170 Dissociative recombination can also account for the net loss of atmospheric components. In the Martian atmosphere, the escape of fast oxygen atoms produced through dissociative recombination of O2+ is the most important nonthermal oxygen atom loss process171−173 and also significantly influences water loss from the planet’s atmosphere.174 This process shows several pathways leading to different electronic states of the product oxygen atoms O2+ + e− → O(3P) + O(3P) + 6.99eV

(58a)

O2+ + e− → O(1D) + O(3P) + 5.02eV

(58b)

O2+ + e− → O(1D) + O(1D) + 3.06eV

(58c)

O2+ + e− → O(1S) + O(3P) + 2.80eV

(58d)

O2+ + e− → O(1S) + O(1D) + 0.83eV

(58e)

H+ + H 2O → H + H 2O+

(59)

H 2O+ + H 2 → H + H3O+

(60)

The resulting H3O+ ions are efficiently destroyed through dissociative recombination, which causes a significant decrease of the H+ abundance in the Saturnian ionosphere. This has consequences for its whole structure:181 Whereas H+ is the dominant ion in the higher layers of the Kronian ionosphere, H3+ dominates below the electron abundance peak. In addition, there is also a strong diurnal variation. Jupiter and Saturn have strong magnetic fields and thus possess large magnetospheres, which even extend as far as their satellites. Due to the high mass of Jupiter and Saturn, meteorites can play a role in ionization of their atmospheric molecules. The consequently very high impact velocity of these bodies hitting the atmosphere of these planets enables quick ionization processes through collisions of ablated atoms with atmospheric atoms and molecules. The presence of several ions, such as Mg+, O+, C+, Si+, Na+, and S+, in gaseous planet ionospheres is caused by meteoritic influx (which amounts to 20 000 ton a day in the case of Jupiter). Models predict the existence of a metal ion layer originating from meteorite impact at an altitude of 350−450 km in the Jovian atmosphere.182 On Titan, meteoritic ablation can contribute ionization at an altitude of 700 km, creating an ionospheric layer containing the metal ions Si+, Fe+, and Mg+.183 The influence of impactors on the ionosphere chemistry of gaseous planets can even be more dramatic: Upon collision with the fragments of the comet Shoemaker-Levy in 1994, the ionosphere became contaminated by molecules such as S2, H2S, CS2, and NH3 from lower layers of the atmosphere. S2 can be photolyzed by solar UV irradiation and the generated S atoms photoionized

Since the escape energy for O atoms from the Martian atmosphere at an altitude of 200 km is approximately 2.0 eV, only oxygen atoms emerging from the first channels 58a and 58b have enough kinetic energy to escape the gravitational field.172 Thus, the escape rate is highly dependent on the branching ratios of the different pathways of the dissociative recombination of O2+. These have been investigated by Petrignani et al.175 in a storage ring experiment, which yielded that at relative reactant energies below 280 meV (i.e., at conditions relevant to the Martian atmosphere) the main channel is 58b (branching fraction 33−54%), followed by channel 58a (branching fraction 22−33%) and 58c (branching fraction 17−29%). The rest is made up by channel 58e. This illustrates that dissociative recombination can contribute to oxygen loss in the Martian atmosphere. 2.7.2. Gaseous Planets. Gaseous planets as well as Titan (the only satellite with a dense atmosphere in the solar system) show extensive ion chemistry. This is illustrated by the fact that the important ion H3+ was discovered in the Jovian atmosphere before it was observed in the interstellar medium.176 In the thermospheres of gaseous giant planets, solar EUV radiation and charged particles impinging on molecules are the most important ion sources. Whereas solar EUV radiation dominates at low to middle latitudes, particle irradiation is most important in circumpolar regions, where the magnetic field lines cross the atmosphere. At lower altitudes, where solar UV rays and charged particles do not reach, ionization by cosmic rays may become important.177 Since H2 is by far the dominant molecule in the exosphere of all gaseous planets, ion production mainly starts through ionization of H2 and the produced H2+ ions quickly react to form H3+ according to reaction 15. This ion is, however, rapidly consumed through dissociative recombination due to the high electron densities in the ionospheres of gaseous planets.178,179 However, ionization in gaseous planets also leads to H+, which only reacts slowly with electrons by inefficient degradation mechanisms like radiative recombination and thus becomes the dominant ion in the Jovian atmosphere.180

S2 + hν → 2S

(61)

S + hν → S+

(62)

Models predict that S and S+ had become the most dominant species in the Jovian ionosphere for a short time. This should also manifest itself in the appearance of unusual sulfurcontaining ions, like H3CS+ and SO+.184 Ion influx is not restricted to meteorites and (in rare cases) comets. In Jupiter’s magnetosphere, fast sodium atoms created by the dissociative recombination of sodium-containing ions that probably originate from Io’s exosphere185 form the extensive sodium nebula around Jupiter.186 The same reaction is also thought to be responsible for the sodium observed in cometary tails187 and near Io.188 Another possible ionization source in atmospheres of gaseous planets is lightning, since evidence for this phenomenon in the atmospheres of the outer gaseous planets has been found.189 Generally, H3+ is a highly versatile diagnostic tool in the ionospheres of gaseous planets. It was used as a tracer of energy input into the ionosphere190 and to study the energy balance of the upper thermosphere.191 Studies of its abundance and distribution are vital not only to understanding the ion chemistry of gaseous planet ionospheres but also to assessing the energetics and dynamics of these environments.192 Emission spectra of H3+ also provide detailed information L

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about the temperature regime in the atmosphere. 180 Furthermore, the ion serves as an ionospheric coolant.193 The versatile ionosphere chemistry in gaseous planets manifests itself in a multitude of reactions involving ions. H3+ ions rapidly react with ion acceptors present in the Jovian atmosphere such as CH4, C2H2, and C2H4 under proton transfer. This mechanism of production of hydrocarbons dominates at higher altitudes (∼400 km). Further down, photoionization of CH4 and CH3 leading to CH+, CH2+, and CH4+ takes over. The produced ions react efficiently with molecular hydrogen to form the stable species CH3+ and CH5+. In the absence of larger hydrocarbons the former ion undergoes radiative association with H2 to form the latter. In high layers of the Jovian atmosphere CH5+ is destroyed by dissociative recombination, leading mostly to CH3 radicals.79,193−195 At altitudes lower than 380 km, where the hydrocarbon density is higher, CH3+ and CH5+ ions can react with acetylene and methane to produce hydrocarbon ions with two or three carbon atoms CH3+ + CH4 → H 2 + C2H5+

Figure 9. Sources of ionization in Titan’s atmosphere. Medium ionization rate by magnetospheric electrons is shown together with the one under a strong precipitation event. Reprinted with permission from ref 389. Copyright 2009 Elsevier.

its atmosphere, whereby high-energy protons (>30 keV) can even reach down to an altitude of 650 km.205,206 Because different species penetrate to different layers of the ionosphere, the individual particles create different layers of ionization in the ionosphere.207 Whereas photoionization is thought to dominate on the day side, the picture is less clear on the night side of the moon, and transportation of ions from the day side to the night side may also play a role.208,209 Below an altitude of 1400 km, however, chemistry rules the composition of the ionosphere since chemical lifetimes are inferior to transport time constants there.210 At very low altitudes (95%) by dissociative recombination, whereas CO+ + RGM → C + O + RGM+ accounts for less than 5%. In the experiment 13CO was used in order to avoid contamination of N2+ in the ion beam. Peaks resulting from a hit by a single 13C atom and a single O atom are resolved and well separated from the C + O peak. Reprinted with permission from ref 284 (http://link.aps.org/abstract/ PRA/v57/p4462). Copyright 1998 by the American Physical Society. (b) Pulse-height spectrum recorded in a single-pass merged-beam experiment with a CO+ ion beam at 400 keV,285 i.e., almost a factor of 10 lower beam energy than in a. This leads not only to a lower resolution (peaks from the C and O atoms are not resolved) but in combination with a higher background pressure to a much larger contribution (85−90%) to the C + O peak from the CO+ + RGM → C + O + RGM+ charge-transfer process. Reprinted with permission from Reference 285. © IOP Publishing. Reproduced by permission of IOP Publishing. All rights reserved.

molecule networks developed to describe the chemistry in interstellar clouds14−17 the main focus was on the final products. Thus, it was important to know not only the rate coefficients for ion−electron reactions but also the product branching ratios. Although some important work has been done in determining product branching ratios in plasma-afterglow experiments,273−275 the lion’s share has been determined with another technique even more remote from classical chemical experiments: ion storage cooler rings. Just as the experimental break through in the study of ion−electron reactions drew heavily from the wartime R&D in electronics, the ion storage rings for atomic, molecular, and chemical physics are off-springs of the development in high-energy physics. Before we take a closer look at the storage rings, the dissociative recombination process deserves closer scrutiny. In order for a free electron to be captured by an ion, it must surrender its kinetic energy. In the absence of a third body, a naked ion such as H+ can only capture a free electron by emission of a photon−radiative recombination. The end result is formation of a neutral H atom, but the rate coefficient is very small, on the order of 10−12 cm3s−1. Atomic ions with bound electrons can also recombine by means of dielectronic recombination, in which the kinetic energy is removed by excitation of a bound electron and the capture stabilized by emission of a photon. Photon emission is a slow process, and the captured electron will mostly be ejected by autoionization, resulting in a small recombination rate. In contrast, a molecular ion recombines far more effective. The reason is that it has a

degree of freedom not available to an atomits nuclear motion. The process can be regarded as taking place in two steps AB+ + e− → AB*

(86a)

The electron surrenders its kinetic energy by exciting a bound electron, leading to formation of a doubly excited molecule AB*. The potential curve for AB** is in most cases repulsive, and immediately following the capture, A and B start to separate on a femtosecond time scale. Dissociation traps the electron, and the molecule falls apart

AB* → A + B

(86b)

The process is effective because dissociation on a repulsive potential curve competes victoriously with autoionization, and it is complex because even if the incoming electron is in the millielectronvolts energy range, electron capture brings the R

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required in order to mix particles of opposite charge at very low relative energies. The ability to control beams of positively charged atoms and molecules was a prerequisite when Thomson invented mass spectrometry.277 Thomson’s student Ashton improved the technique278 to the point where he could make important chemical discoveries (isotopes of nonradioactive elements) worthy of a Nobel Prize in chemistry (in 1922). The ability to control two beams of differently charged particles that could interact at very low energies (merged beams) took much longer time to develop.279 In the early 1990s accelerator technology had advanced even further thanks to developments in particle and nuclear physics, and the merged-beam technique could then be combined with ultrahigh vacuum and beam-cooling techniques, which allowed new opportunities in atomic and molecular physics. Figure 10 shows schematically an ion storage ring and how an experiment is conducted. The ion storage ring technique has the following advantages. (1) Ultrahigh vacuum (107 cm3 s−1 instead of the 10−11 cm3 s−1 inferred from flowing afterglow experiments.307 Considering the controversy, the astrophysical importance of H3+, and the fact that it was

eminently suited for a storage ring experiment, dissociative recombination of H3+ was chosen for the first experiment at CRYRING in Stockholm.310,311 The results agreed with Amano’s309 and were strongly in variance with the afterglow experiments.307 The CRYRING result suggested that relaxation of vibrationally hot H3+ ions to their lowest vibrational level was achieved by storage in the ring for about 10 s, but a definitive proof did not come until 2002,312 when experiments at the storage ring TSR unambiguously revealed that infrared radiation depletes vibrationally excited levels in H3+ on the anticipated time scale. However, the TSR also revealed a problem; despite long storage it was impossible to remove highly excited rotational levels. The question naturally arose: what if the rate coefficient varied strongly with the rotational quantum number? The experimentalists responded to the challenge, and at the same time a theoretical team put forward a plausible mechanism that for the first time could explain a high recombination rate.313,314 It was realized313 and implemented into a theoretical framework314 that the incoming electron would distort the equilateral triangle H3+ such that electron capture would become much more efficient than previously thought. Simultaneously, a Stockholm−Berkeley collaboration designed, constructed, and characterized a discharge pinhole supersonic jet source which could deliver rotationally cold H3+ ions and used it as injector to CRYRING in a cross-section measurement of dissociative recombination.45,315 Figure 14 is a beautiful example of a comparison of a laboratory absorption spectrum of H3+ and an absorption spectrum of a diffuse cloud toward the star ζ Persei.

Figure 14. Spectra of two H3+ transitions arising from the two lowest rotational levels, which are the only levels populated in diffuse clouds. R(1,1)u originates from the lowest para level (J=1, K = 1), whereas R(1,0) comes from the lowest ortho level (J = 1, K = 0). Top trace is a laboratory absorption spectrum recorded with a cavity ring-down spectrum of a supersonic discharge expansion source, while the lower trace is a spectrum of a diffuse cloud toward ζ Persei. Reprinted with permission from ref 45. Copyright 2003, Rights Managed by Nature Publishing Group.

The results from CRYRING45,315 and the theoretical calculations314 were very encouraging in that they gave almost the same rate coefficient at an electron temperature of 300 K, α = 0.7 × 10−7 cm3 s−1, and soon after, experiments at TSR confirmed the CRYRING results.316 Although this was an unprecedented agreement in the history of H3+, it did not solve all problems. However, it strongly suggested45 that the cosmicray ionization rate ζ2 in a diffuse cloud such as the one toward ζ Persei is considerably larger than the one in a dense cloud, V

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which was assumed to be about 3 × 10−17 s−1. It had been believed that the rate was similar in diffuse and dark clouds, but the CRYRING rate coefficient45,315 applied to reaction 100 suggested otherwise. Many more diffuse cloud sight lines have now been analyzed with respect to H3+ absorption and the cosmic-ray ionization rate found to be in the range (1.7 ± 1.3) × 10−16 s−1 < ζ2 < (10.6 ± 8.2) × 10−16 s−1.317 Application of the cosmic chemistry of H3+ to derive the cosmic-ray ionization rate is a beautiful example of the interplay between astrochemistry and high-energy astroparticle physics,318 which has probably not been appreciated to the full extent either by chemists or by physicists. Interstellar ionization is caused by those particles in the cosmic rays that have an energy of less than about 100 MeV. This is the energy range for which the solar wind and interplanetary magnetic fields prevent the particles (protons) from reaching the Earth, thus making terrestrial measurements of the low-energy tail of the cosmic rays impossible. The origin of the cosmic rays is a central problem in contemporary astroparticle physics. Laboratory astrophysics, astronomical observations, and astrochemistry combine in assisting to provide clues to this problem. There are a few remaining problems that need to be solved for H3+. Experiments at TSR have shown that the rotational temperature in both the CRYRING and the TSR experiments were higher than previously believed,46,319 probably due to heating effects when the ions are extracted and accelerated from the ion source. Despite this uncertainty, the agreement between the absolute cross-section measurements in the two storage rings is remarkable,45,46,315 as shown in Figure 15. The new Cryogenic Storage Ring (CSR)320 at the Max Planck Institute für Kernphysik (Heidelberg) will provide optimal conditions for finally solving the problem.

H3+ gives a complete coverage of all the fascinating aspects of this molecular ion, so central to astrochemistry. Space limitations have also prevented us from giving a more exhaustive description of the many experimental results on product branching ratios coming from, in particular, CRYRING in Stockholm. The reader is referred to the very recent review on ion chemistry in space6 and the research monograph on dissociative recombination275 for complete references. Of the ion storage rings that produced so many results for astrophysics and other branches of physics during the 1990s and 2000s, only the TSR is still in operation but only until the new CSR has been commissioned and is ready for experiments. In Lanzhou in China, the Molecular Ion Research Facility (MIRFL)322 at the storage ring CSRe is under construction. It has a much higher magnetic rigidity (bending power) than the previous generation storage rings and will make possible studies of electron recombination with heavier ions (100−150 amu) than was accessible in the smaller rings. The higher energy will also increase the resolution in the decay channels. Together with the CSR it is going to be the main accelerator and storage ring facility for studies of dissociative recombination during the next decade. 4.2. Ion−Neutral Reactions

Ion−molecule reactions in weakly ionized plasma are part of the chemistry of high-energy reactions since the initial step involves imparting some tens of electronvolts to molecules and thus causing ionization. The ionizing agent differs depending on the situation and the environment. In the Earth’s ionosphere ionization is caused by the solar radiation, whereas it is primarily the cosmic rays that ionize the dominant interstellar molecule H2. It was the desire to understand the chemistry of the terrestrial ionosphere that initiated a research program at the then National Bureau of Standards in Boulder, CO, in 1962.323 Plasma-afterglow techniques had been around for more than a decade,270,271 but something more chemically versatile was needed. This led to development of the flowing afterglow technique.324,325 The technique was then developed in several laboratories in directions which suited the particular application. The selected ion flow tube (SIFT) (see Figure 16) is noteworthy326 since it was used to study ion−molecule reactions of interstellar significance.327 The comprehensive review by Graul and Squires328 is a good starting point for learning about afterglow techniques in gas-phase ion chemistry. The CRESU (Cinétique de Reaction en Ecoulment Supersonic Uniformé) technique329 uses a supersonic expansion to reach the low temperatures prevailing in the interstellar medium. It has been used to study ion−molecule reactions330 but far more to study neutral−neutral reactions.331 The third technique is based on ion trapping. The reviews, both in 1968, on ion−molecule reactions332 and ion cyclotron resonance (ICR) trapping techniques333 do not discuss astrochemistry, simply because it was not yet on the scientific agenda. This changed, of course, in the early 1970s, as discussed in the Introduction, and experimentalists responded to the challenge of obtaining understanding of detailed molecular processes in laboratory experiments. An ion-trap study334 of dissociative recombination was a prelude to the first measurement of an astrophysically important ion−molecule reaction rate coefficient at interstellar temperature; the radiative association reaction

Figure 15. Dissociative recombination rate coefficient in the electron− ion merging region (the relative velocity times the effective cross section) as a function of the center-of-mass collision energy. Black squares are data from the TSR,46 whereas gray dots are data from CRYRING.45,315 Data from the two storage rings are absolute and just put on top of each other. Small deviation at very low center-of-mass energy is due to the slightly colder electron beam in the TSR. Reprinted with permission from ref 46 (http://link.aps.org/abstract/ PRA/v82/e042715). Copyright 2010 by the American Physical Society.

Space limitations have prevented us from giving a complete picture of the fascinating and complex H3+ ion. We decided to focus on the interplay between laboratory astrophysics and astrochemistry to address a central problem in astroparticle physics. The proceedings321 from the recent Royal Society Discussion Meeting on Chemistry, Astronomy, and Physics of W

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set in Figure 18 is obtained. The reader is referred to the tutorial by Wester342 and the review by Gerlich and Smith343 for more details.

Figure 16. Schematic diagram of a typical SIFT apparatus illustrating its major features. Quadrupole mass filter is used to select the ion. Carrier gas flow brings the ions into contact with the various reactant gases. Product ions resulting from the ion−neutral reactions are mass selected and detected by a channeltron. Flow tube is typically 100 cm long and 8 cm in diameter. Reprinted with permission from ref 390. Copyright 1988 Elsevier (Academic Press).

CH3+ + H 2 → CH5+ + hν

Figure 18. Measured decay of H2+ and formation of H3+ as a result of reaction 15 and as a function of storage time. Linear fit to the data set for H2+ (straight line) gives the reaction rate coefficient. Reprinted with permission from Reference 342. © IOP Publishing. Reproduced by permission of IOP Publishing. All rights reserved.

In the cold part (≤20 K) of the dark molecules clouds, in the regions where stars are formed, H3+ is converted into its deuterated isotopologues by ion−molecule reactions

(101)

was measured at 13 K.335 The detailed description336 of the Penning ion-trap technique used to measure reaction 101 also gives an interesting view of the difference between the physicist’s trapping technique (hyperbolically shaped Penning trap) and that of chemists (cubic ICR cell). The trapping time in the Penning trap is much longer, something which served the physicists well in precision experiments which demanded long, nonintrusive observation time (≫hours) and rendering Nobel Prizes in physics in 1989337 and 2012.338 However, the Penning trap never became the laboratory work horse for the study of ion−molecule reactions. Gerlich developed a radiofrequency (rf) ion-trap apparatus which combined an ion source in which the ions could be trapped and thermalized, a quadrupole ion guide, and an rf ion trap that could be cooled down to 10 K.251,339,340 In particular, the 22-pole trap (see Figure 17) with its almost square-well-like potential is very well suited for collision experiments.341 Consider the exothermic H3+-producing ion molecule in reaction 98. If H2+ ions are injected and stored in the trap and a controlled amount of H2 is leaked into the trap, H3+ will be produced while H2+ is depleted. By emptying the trap after various storage times and mass analyzing the content, the data

H3+ + HD → H 2D+ + H 2 + 232K +

+

H 2D + HD → D2 H + H 2 + 187K +

+

D2 H + HD → D3 + H 2 + 234K

(102) (103) (104)

where the negatives of the reaction enthalpies for simplicity are given in units of temperature (K). The backward reactions are endothermic and hence not energetically possible at 20 K. In contrast to the symmetric, equilateral triangle H3+ (and D3+), which lacks a rotational spectrum and can only be observed in the infrared, H2D+ and D2H+ can be studied by rotational spectroscopy in the far-infrared and submillimeter wavelength regions. This makes them excellent probes into the star-forming regions, where infrared spectroscopy cannot reach.344−346 The new Atacama Large Millimeter/submillimeter Array (ALMA), the largest ground-based telescope in the world, was inaugurated on March 13, 2013. It will allow the recording of high-resolution maps of prestellar H2D+. A remaining question which must be addressed by laboratory work is the rate constant for reaction 102. The 22-pole trap was used in two independent experiments,83,347 which differ by about a factor of

Figure 17. Cross-sectional view of the 22-pole ion trap described in refs 339−341: (left) view parallel to the symmetry axis; (middle) perpendicular view. Photo to the right shows a photo of the 22-pole trap. Reprinted with permission from Reference 342. © IOP Publishing. Reproduced by permission of IOP Publishing. All rights reserved. X

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CnN− + H → CnNH + e−

4, the older result being systematically lower. This is an unsolved problem. The laboratory studies of ion−electron and ion−molecule reactions have in some cases led to the realization that not all molecules are formed in the gas phase. One such example is interstellar methanol. A possible formation route in the gas phase would be radiative association of CH3+ and water followed by dissociative recombination CH3+ + H 2O → CH3OH 2+ + hν

(105)

CH3OH 2+ + e− → CH3OH + H

(106)

The reaction efficiency (the ratios of experimental rate constants to calculated collision rate constants) was found to fall in the range 0.23−0.35, essentially independent of n.354 The anions CnN2− have a slightly lower reactivity, but the absence of a dependence on n persists also for carbanions containing two N atoms.354 A third N atom leads to no reactivity.354 Thus, although atomic hydrogen is very reactive, the reaction efficiency with CnNx− anions is much less than unity. They are therefore likely to be highly stable and can survive in the interstellar medium. 4.3. Ion−Ion Reactions

A plasma afterglow experiment at room temperature and at 225 K348 combined with theoretical extrapolation to 20 K and some additional estimates were used to extract a binary rate coefficient of 2.5 × 10−9 cm3 s−1.123 It was also assumed that reaction 106 could dominate the dissociative recombination of protonated methanol, which would suggest that reactions 105 and 106 would be an effective route of forming interstellar methanol. This conclusion, however, did not survive a measurement by means of the 22-pole trap of the rate coefficient for reaction 105 at 20 K to be 2 × 10−12 cm3 s−1 and a storage ring measurement of the branching ratio of eq 106 to be 0.03.349 Hydrogenation of CO on grain surfaces is a more probable formation route and consistent with astronomical observations of rotational transitions in 12CH3OH and 13 CH3OH.350 This illustrates the importance of combining the best available laboratory experiments with astronomical observations in order to solve a concrete problem. Grain surface chemistry is not necessarily a universal recipe for producing interstellar molecules. The abundance of methyl formate, HCOOCH3, cannot be explained by surface chemistry, and gas-phase formation routes must be found.351 Methyl formate’s isomer glycolaldehyd (HCOCH2OH), which contains both a functional group (aldehyde) of sugar and the hydroxyl group, has also been observed toward the galactic center352 and in a solar-type protostar.353 A flowing afterglowselected ion flow tube (FA-SIFT) apparatus was used351 to measure ion−molecule reactions involving protonated methanol and formic acid

Due to the detection of anions in interstellar clouds, circumstellar envelopes, and Titan’s atmosphere, mutual neutralization of anions and cations has moved into the focus of interest of astrochemists. The process was first identified by Thomson and Rutherford in 1896, who also measured the first rate constants.265 The boost of development of mass spectrometers and microwave discharge devices after World War II enabled further measurements.355,356 Mutual neutralization reactions have two important properties that distinguish them from dissociative recombination. First, the detachment energy of the anion has to be counted into the total reaction enthalpy, which makes the process less exoergic and constrains the number of possible reaction pathways. As opposed to dissociative recombination, which often leads to three or more neutral fragments, mutual neutralization can thus be expected to be a “softer process” that is more likely to conserve the structure of the reactants. Part of the reaction enthalpy can also be used for the internal rovibronic excitation of neutral product molecules. Production of excited states automatically reduces the amount of energy available for fragmentation of the products. The lower energy yield of mutual neutralization might make it a feasible final step of the synthesis of more complex molecules in the interstellar medium from protonated species (e.g., methanol from CH3OH2+) in the interstellar medium, since dissociative recombination has often been found not to lead to the “desired” product but to result in fragmentation.275 Second, as opposed to electron recombination partners in mutual neutralization, kinetic energy can be released without one of the neutrals formed from the reactant cation having to dissociate or radiate off energy. Also, this property can help the conservation of the structure of complex ions during mutual neutralization. Mechanistically, mutual neutralization can proceed as a binary or ternary (buffer gas assisted) process. Also, an electron-catalyzed mechanism for mutual neutralization reactions has been invoked.357 Due to the low densities in the interstellar medium, only binary processes are relevant for this environment. Regarding astrochemistry, mutual neutralization has a very long historyit is even thought to play a role in the early universe. Reaction of hydride anions with hydrogen cations

CH3OH 2+ + HCOOH → HC(OH)OCH3+ + H 2O (107)

CH3OH 2+ + HCOOH → CH3OH 2+ − HCOOH

(108)

CH3OH 2+ + HCOOH → HCOOH 2+ + CH3OH

(109)

(110)

It was found that the overall rate coefficient was 3.19 × 10−10 cm3 s−1 and that the branching ratio for the adduct ion in eq 108 was 0.95, whereas the branching ratio for protonated methyl formate was found to be 0.05 and that for proton transfer negligible.351 However, it is conceivable that the branching ratio for the adduct ion is smaller in the interstellar medium, where the pressure is lower than in the flow tube. As always, the closer to interstellar conditions one can perform an experiment the better. Observation of interstellar anions20−26 has stimulated laboratory studies of their chemistry. Three of the observed anions are nitrogen-containing carbanions.24−26 The FA-SIFT technique was used to study reactions of CnN− (n = 1−6) anions with H atoms,354 which proceeds mainly through associative detachment

H+ + H− → 2H

(111)

was very probably a crucial process there since it removed H− ions formed through radiative attachment of hydrogen atoms. Hydride anions played a pivotal role in the primeval universe, since they form molecular hydrogen in the later phase of the early universe358 through Y

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of reaction 111 at low energies using a merged-beam device has been carried out.372 Again, the rate constant measured was in contradiction with the result by Moseley et al.369 but agreed well with previous merged beam findings of Szucs et al.370 and Peart and Hayton.371 The results are also in line with the most recent theoretical calculations.364 The sensitivities of the outcome of model calculations of protogalactic gas collapse (i.e., the temperature and final densities predicted by them) to the mutual neutralization rate constants used as inputs into these model have been evaluated by Glover et al.358 Whereas changes in the rate constant of reaction 111 only had a minor effect on the results if cold initial conditions (12 K) were assumed, they exerted a quite dramatic influence if collapse was assumed to start from a hot (104 K) highly ionized gas under hot initial conditions. In this scenario, a high rate constant of reaction 111 leads to less efficient H2 formation. The sensitivity of the model is even greater when the presence of external UV radiation is assumed. This example shows that it is vital to obtain the exact rate coefficient for mutual neutralization reactions at a variety of different temperatures. The double storage ring DESIREE located at the Physics Department of Stockholm University, which has been commissioned this year will be used to study mutual neutralization reactions between astrophysically important anions and cations under a large range of collision energies, covering the low-temperature regime present in dark clouds as well as the possibly hotter temperatures present in the early Universe.256,373,374 The importance of mutual neutralization is not restricted to small cations and anions in the early universe. For example, in the F layer of Earth’s atmosphere mutual neutralization of oxygen ions with opposite charge is a major source of excited oxygen atoms which emit characteristic UV radiation.375 Inclusion of mutual neutralization reactions of polyaromatic hydrocarbons (PAHs) into chemical models of astronomical environments was found to have a profound influence on the abundance of many species. Generally, degradation of molecular cations by mutual neutralization is slower than that by dissociative recombination. Atomic ions, on the other hand, decay faster by mutual neutralization with PAH anions than by electron−ion reactions like dielectronic or radiative recombination. Mutual neutralization can consequently lead to a reduction of the total ionization fraction predicted by chemical models of astronomic environments if reactions involving PAHs are included.376 Furthermore, as pointed out by Lepp and Dalgarno already in 1988, PAH− anions become the major carriers of the negative charge in the ISM if the abundances of PAHs relative to H2 exceed a certain value (10−8 in their model).131 This could allow mutual neutralization to compete with dissociative recombination as a main neutralization process of positive ions, even if rate constants of mutual neutralizations tend to be 1 or 2 orders of magnitude lower. It has further been predicted that mutual neutralization of atomic ions can greatly affect the charge balance and even the abundance of neutral carbon, sulfur, silicon, and metals in photon-dominated regions through reactions of the type377

(112)

Thus, the rate of reaction 111 determined the amount of H− ions that was converted into molecular hydrogen in the early universe at later stages (at red shifts lower than z = 100), when photodetachment was less important and collisional process dominated the H− destruction.359,360 Substantial variations in the rate coefficients of reactions 111 and 112 could then result in significant differences in the production rate of H2.360 The presence of molecular hydrogen is fundamental for the cooling of primeval clouds in the early universe to enable collapse into the first Population III stars, since it can provide cooling at temperatures lower than 10 000 K, where Lyman-a cooling by hydrogen atoms it not possible anymore. Another possibly interesting mutual neutralization process in the early universe is the reaction H 2 + + H− → H 2 + H

(113)

This process has been thoroughly investigated by Bates and Lewis361 using a Landau−Zener approach, in which the reaction is assumed to proceed along covalent potential curves crossed by an attractive ion-pair potential. Hickman362 formulated a general complex potential model for mutual neutralization reactions. Two different mechanisms of mutual neutralization were assumed in this approach: (a) simple electron and (b) association forming a neutral molecule from the two reaction partners. As a result, the following general formula for the rate constant of mutual neutralization was derived for T < 1000 K k = 2.28 × 10−5(T /300)EA −0.4μ−0.5 cm 3s−1

(114)

where EA denotes the electron affinity (in eV) and μ the reduced mass in atomic units. A further theoretical study on mutual neutralization processes using a quantum close-coupling approach has been published,363 and recently, Stenrup et al.364 presented a further theoretical treatment of reaction 111. Also, the rate constant of the mutual neutralization reaction of Li+ and H−, Li+ + H− → Li + H

(115)

which is of possible significance in the early universe as well as in thermalization of stellar atmospheres,365 has been investigated theoretically at temperatures up to 10 000 K using a Landau−Zener approach by Croft et al.366 However, it has to be stated that not all possible mutual neutralization processes play a fundamental role in the early universe. For example, model calculations predicted that reaction 112 lacks importance for this environment.367 Compared with other ion reactions there has not been much experimental data on mutual neutralizations. Many studies also only have covered high collision energies. There have, however, been flowing-afterglow Langmuir probe (FALP) studies of a number of mutual neutralization reactions a room temperature,368 but the first measurement of the mutual neutralization of H+ with H− at low collisional energies relevant to the early universe (the energy range covered was 0.15−300 eV) has been performed by Moseley et al.369 in superimposed beam experiments. In this study, a rate constant of (4.0 ± 1.8) × 10−7 cm3 s−1 was obtained, which exceeds the value determined in the theoretical investigation by Bates and Lewis362 by a factor of 3. Subsequent experimental studies,370,371 however, resulted in lower cross sections and a rate constant more in agreement with Bates and Lewis.361 Very recently, a further measurement

PAH− + C+ → PAH + C

(116)

Mutual neutralization consequently also favors transformation of C+ to CO via C in these environments. Model calculations of photon-dominated regions also yield that, as mentioned above, mutual neutralization dominates over radiative recombination Z

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as a neutralization process of atomic metal ions if PAHs are present. It should be kept in mind that model predictions are highly dependent on the rate constants of mutual neutralizations serving as input to these models. These can be calculated using the assumption that cation and anion have to approach each other at a certain critical distance to allow the reaction to take place. In the case of large (PAH) anions the anion radius can be used as the critical distance and the following formula can be derived for the rate constant k = πa 2 v0 [1 + (e 2 /akBT )]

More recently, the VENDAMS method (variable electron and neutral density attachment mass spectrometry), a further development of the flowing-afterglow Langmuir-probe technique, has been employed to determine rate coefficients for ion− ion mutual neutralization along with neutral product branching fractions.381,382 It allows one to study mutual neutralization of anions formed by electron attachment to fragments of molecules in a noble gas plasma. Also, electron attachment rate constants can be determined.383,384 At the moment, however, the method is limited at present to noble gas positive ions, which are not the main reactants in the interstellar medium, and it has to be hoped that the range of reactants to be studied can be extended. Furthermore, the employed ion concentrations lie several orders of magnitudes above interstellar conditions. Nevertheless, an interesting trend of rate constant of mutual neutralization has been found in the VENDAMS experiment: reactions involving polyatomic ions tend to proceed about at least an order of magnitude faster than those involving only mono- and diatomic ones.384,385 These findings might allow the conclusion that mutual neutralization is a process that can compete with other processes if fairly large ions are involved. The electrostatic storage ring DESIREE256,373,374 will be able to handle very heavy ions and thus be a very feasible device for investigating the properties of mutual neutralizations of cluster ions.

(117)

in which a is the anion radius, e the elementary charge, kB Boltzmann’s constant, and ⟨v0⟩ the average relative kinetic energy of the reactants. This energy can be expressed as v0 = (8kBT /πμ)1/2

(118)

in which μ denotes the reduced mass of the reactants. In case of very large anions at elevated temperatures, eq 117 can be simplified to k = πa 2 v0

(119)

This implies that in the case of large anions (e.g., PAHs) the rate constant simply scales with the square of the ion radius at a certain temperature, which represents a simple ballistic regime. If the ion, on the other hand, is small and the temperature low the expression in the brackets of eq 118 is dominated by the second summand, which results in a linear dependence of the rate constant on the ion radius.376 For the mutual neutralization of a PAH anion with a radius of 4 Å and a cation with a mass of 30 amu, application of eq 118, this leads to a mutual neutralization rate constant of 2.0 × 10−7 cm3 s−1. Another group of large ions for which mutual neutralization plays a substantial role are cluster ions. Large water cluster ions like NH4+(H2O)n and HSO4−(H2O)n have been detected in the ionosphere of Earth but are also predicted to be present in the lower Martian atmosphere. In this environment, cluster anions are thought to constitute the majority of the negatively charged particles. They are also thought to play an important role in the D layer of the terrestrial ionosphere. Cluster ions are mainly destroyed by mutual neutralization in planetary ionospheres. Large neutral clusters formed by mutual neutralization can function as nucleation sites for formation of aerosols. Due to this fact,378 their degradation and formation processes of cluster ions are of great interest for atmospheric scientists. Despite earlier and current research efforts there is still a great need for experimental data on rate constants and product branching ratios of mutual neutralization reactions of clusters, although such reactions were already studied a long time ago. For example, Smith et al.379 presented a flowing afterglow Langmuir probe study on mutual neutralization of the H3O+(H2O)3 cluster ions with different negative ions (NO2−, NO3−, and their water clusters) as early as 1976. In this investigation, reaction rate constants around 5.6 × 10−8 cm3 s−1 were measured, which are close to those found for monomer ions. Further studies concerned reaction of hydronium−water H 3 O + (H 2 O) 3 and hydronium−acetonitrile H 3 O + H 2 O(CH3CN)3 clusters with monomer and cluster ions like NO3−, HSO3−, NO3−(H2O)n, and HSO3−(H2O)n that are expected to be common in the ionosphere. Also, in these measurements fairly similar rate constants were obtained.380

5. CONCLUSIONS In the earlier days of astrochemistry reactions involving ions were regarded as forming the backbone of the synthesis of interstellar molecules due to the absence of kinetic barriers in many of these processes. Later experimental investigations of ion−neutral reactions using ion traps as well as studies of ion− electron reactions employing storage rings and afterglow devices often yielded that the suggested reactions were either too slow or did not lead to the desired products. Thus, surface reactions were increasingly held responsible for formation of many observed interstellar molecules.386 However, we hope to have convinced the reader that ionic processes play a decisive role in both the interstellar medium and the planetary atmospheres. This was corroborated by the fact that many ions have been detected with the Herschel space telescope during its lifetime. We can further anticipate that the ALMA interferometric array will lead to further ion detections in the interstellar medium. With its unprecedented resolution, ALMA will also yield a cornucopia of information about the distribution of different ions in star-forming regions, circumstellar envelopes, and protoplanetary disks. Experimental and theoretical investigations will be necessary to rationalize these findings. With new methods like improved ion traps, cooled storage rings, and plasma afterglow apparatuses, such studies will be possible, and we can be looking forward to a golden age of ion chemistry in the interstellar medium and planetary atmospheres.

ASSOCIATED CONTENT S Supporting Information *

Complete reference (all authors listed) for refs 27, 29, 48, 134, 239, 297, and 298. This material is available free of charge via the Internet at http://pubs.acs.org. AA

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AUTHOR INFORMATION

ACKNOWLEDGMENTS This review was supported by the Swedish Research Council. One of the authors (M.L.) is thankful to T. Oka and for support from the Enrico Fermi Institute at The University of Chicago for making it possible to get the writing started in a stimulating environment. Support for W.G. from the COST Action CM0805 “The Chemical Cosmos: Understanding Chemistry in Astronomical Environments” is gratefully acknowledged.

Corresponding Author

*E-mail: [email protected]. Notes

The authors declare no competing financial interest. Biographies

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Wolf Geppert studied chemistry at the University of Vienna (Austria) and the University of York (UK), where he obtained his Ph.D. degre in 2000. He then obtained a postdoctoral position at the University of Bordeaux France with a focus on Astrochemistry and has been active in the field since then. After a year at the University of Helsinki he was selected for a Marie Curie Individual Fellowship at Stockholm University. In 2007 he got a Senior Researcher position at the same institution and was awarded Associate Professorship in 2009. His main research interests concern studies of ion reactions relevant for dark clouds, circumstellar envelopes, and atmospheres of planets and their satellites. He also performs radioastronomical observations and model calculations of astrochemical reaction networks.

Mats Larsson obtained his Ph.D. degree in Physics at Stockholm University in 1982. After a period as research associate at the Manne Siegbahn Institute, he became Associate Professor at the Royal Institute of Technology in 1991 and Full Professor at Stockholm University in 1996. His main interest is laboratory studies of ion reactions important to astrochemistry. He has been a JILA Visiting Fellow in Boulder and Miller Visiting Professor at Berkeley. He chaired the Space Research Advisory Committee of the Swedish National Space Board for 10 years, and he is member of the Royal Swedish Academy of Sciences and its Physics Class. AB

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