Trihydrogen cation

Protonated molecular hydrogen, trihydrogen cation, or H₃⁺, is one of the most abundant ions in the universe. It is stable in the interstellar medium (ISM) due to the low temperature and low density of interstellar space. The role that H₃⁺ plays in the gas-phase chemistry of the ISM is unparalleled by any other molecular ion. The cation is also the simplest triatomic molecule, since its two electrons are the only valence electrons in the system. It is also an example of a three-center two-electron bond system.

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History

H₃⁺ was first discovered by J.J. Thomson^[1] in 1911. While studying the resultant species of plasma discharges, he discovered something very odd. Using an early form of mass spectrometry, he discovered a large abundance of a molecular ion with a mass-to-charge ratio of 3. He stated that the only two possibilities were C⁴⁺ or H₃⁺. Since C⁴⁺ would be very unlikely and the signal grew stronger in pure hydrogen gas, he correctly assigned the species as H₃⁺.

The formation pathway was discovered by Hogness & Lunn^[2] in 1925. They also used an early form of mass spectrometry to study hydrogen discharges. They found that as the pressure of hydrogen increased, the amount of H₃⁺ increased linearly and the amount of H₂⁺ decreased linearly. In addition, there was little H⁺ at any pressure. This data suggested the proton exchange formation pathway discussed below.

In 1961, Martin et al.^[3] first suggested that H₃⁺ may be present in interstellar space given the large amount of hydrogen in interstellar space and its reaction pathway was exothermic (~1.5 eV). This led to the suggestion of Watson^[4] and Herbst & Klemperer^[5] in 1973 that H₃⁺ is responsible for the formation of many observed molecular ions.

It was not until 1980 that the first spectrum of H₃⁺ was discovered by Takeshi Oka,^[6] which was of the ν₂ fundamental band using a technique called frequency modulation detection. This started the search for interstellar H₃⁺. Emission lines were detected in the late 1980s and early 1990s in the ionospheres of Jupiter, Saturn, and Uranus.^[7][8][9] In 1996, H₃⁺ was finally detected in the ISM by Geballe & Oka^[10] in two molecular interstellar clouds in the sightlines GL2136 and W33A. In 1998, H₃⁺ was unexpectedly detected by McCall et.al.^[11] in a diffuse interstellar cloud in the sightline Cyg OB2 No. 12.

Structure

The arrangement of the hydrogen atoms in the molecule is an equilateral triangle. The molecule has a resonance structure which represents a three-centre, two-electron bond. The strength of the bond has been calculated to be around 4.5 eV (104 kcal per mole).^[12] This molecule is a good example which illustrates the importance of electron pair delocalization which contributes to the stability of molecules.

Formation

The information for this section was from a paper by Eric Herbst.^[13] The main pathway for the production of H₃⁺ is by the reaction of H₂⁺ and H₂.

H₂⁺ + H₂ → H₃⁺ + H

The concentration of H₂⁺ is what limits this reaction. H₃⁺ can only be produced in interstellar space by a cosmic ray ionizing H₂.

H₂ + cosmic ray → H₂⁺ + e^- + cosmic ray

However, the cosmic ray has so much energy, it is relatively unaffected by the relatively small energy required to ionize an H₂ molecule. In interstellar clouds, cosmic rays leave behind a trail of H₂⁺, and therefore H₃⁺. In laboratories, H₃⁺ is produced by the same mechanism in plasma discharge cells, with the discharge potential providing the energy to ionize the H₂

Destruction

The information for this section was also from a paper by Eric Herbst.^[13] There are many destruction reactions for H₃⁺. The dominant destruction pathway in dense interstellar clouds is by proton transfer with a neutral collision partner. The most likely candidate for a destructive collision partner is the second most abundant molecule in space, CO.

H₃⁺ + CO → HCO⁺ + H₂

The significant product of this reaction is HCO⁺, an important molecule for interstellar chemistry. Its strong dipole and high abundance make it easily detectable by radioastronomy. H₃⁺ can also react with atomic oxygen to form OH⁺ and H₂.

H₃⁺ + O → OH⁺ + H₂

OH⁺ then usually reacts with more H₂ to create further hydrogenated molecules.

OH⁺ + H₂ → OH₂⁺ + H

OH₂⁺ + H₂ → OH₃⁺ + H

At this point, the reaction between OH₃⁺ and H₂ is no longer exothermic in interstellar clouds. The most common destruction pathway for OH₃+ is dissociative recombination, yielding four possible sets of products: H₂O + H, OH + H₂, OH + 2H, and O + H₂ + H. While water is a possible product of this reaction, it is not a very efficient product. Different experiments have suggested that water is created anywhere from 5% - 33% of the time. Water formation on grains is still considered the primary source of water in the interstellar medium.

The most common destruction pathway of H₃⁺ in diffuse interstellar clouds is dissociative recombination. This reaction has multiple products. The major product is dissociation into three hydrogen atoms, which occurs roughly 75% of the time. The minor product is H₂ and H, which occurs roughly 25% of the time.

Ortho/Para-H₃⁺

  The most abundant molecule in dense interstellar clouds is H₂. When a H₃⁺ molecule collides with H₂, stoichometrically there is no net yield. However, a proton transfer still can take place, which can potentially change the total nuclear spin of the two molecules depending on the nuclear spins of the protons. Two different spin configurations for H₃⁺ are possible, called ortho and para. Ortho-H₃⁺ has all three proton spins parallel, yielding a total nuclear spin of 3/2. Para-H₃⁺ has two proton spins parallel while the other is anti-parallel, yielding a total nuclear spin of 1/2. Similarly, H₂ also has ortho and para states, with ortho-H₂ having a total nuclear spin 1 and para-H₂ having a total nuclear spin of 0. When an ortho-H₃⁺ and a para-H₂ collide, the transferred proton changes the total spins of the molecules, yielding instead a para-H₃⁺ and an ortho-H₂.^[13]

Spectroscopy

The spectroscopy of H₃⁺ is challenging. Due to its lack of a permanent dipole moment, pure rotational spectroscopy of H₃⁺ is impossible. Ultraviolet light is too energetic and would dissociate the molecule. Rovibronic spectroscopy provides the ability to observe H₃⁺. Rovibronic spectroscopy is possible with H₃⁺ because one of the vibrational modes of H₃⁺, the ν₂ asymmetric bend mode, has a weak dipole moment. Since Oka's initial spectrum,^[14] over 900 absorption lines have been detected in the infrared region. H₃⁺ emission lines have also been found by observing the atmospheres of the Jovian planets. H₃⁺ emission lines are found by observing molecular hydrogen and finding a line that cannot be attributed to molecular hydrogen.

Astronomical detection

H₃⁺ has been detected in two types of celestial environments: Jovian planets and interstellar clouds. In Jovian planets, it has been detected in the planet’s ionospheres, the region where the Sun’s high energy radiation ionizes the particles in the atmosphere. Since there is a high level of H₂ in these atmospheres, this radiation can produce a significant amount of H₃⁺. Also, with a broadband source like the Sun, there is plenty of radiation to pump the H₃⁺ to higher energy states from which it can relax by stimulated and spontaneous emission.

Planetary atmospheres

The detection of the first H₃⁺ emission lines was reported in 1989 by Drossart et.al.,^[15] found in the ionosphere of Jupiter. Drossart found a total of 23 H₃⁺ lines with a column density of 1.39 * 10⁹ cm^-2. Using these lines, they were able to assign a temperature of the H₃⁺ of ~1,100 K, which is comparable to temperatures determined from emission lines of other species like H₂. In 1993, H₃⁺ was found in Saturn by Geballe et.al.^[16] and in Uranus by Trafton et.al.^[17]

Molecular interstellar clouds

H₃⁺ was not detected in the interstellar medium until 1996, when Geballe & Oka^[18] reported the detection of H₃⁺ in two molecular cloud sightlines, GL2136 and W33A. Both sources had temperatures of H₃⁺ of about 35 K and column densities of about 10¹⁴ cm^-2. Since then, H₃⁺ has been detected in numerous other molecular cloud sightlines, such as AFGL 2136,^[19] Mon R2 IRS 3,^[20] GCS 3-2,^[21] GC IRS 3,^[21] and LkHα 101.^[22]

Diffuse interstellar clouds

Unexpectedly, three H₃⁺ lines were detected in 1998 by McCall et.al.^[23] in the diffuse cloud sightline of Cyg OB2 No. 12. Before 1998, the density of H₂ was thought to be too low to produce a detectable amount of H₃⁺. McCall detected a temperature of ~27 K and a column density of ~ 10¹⁴ cm^-2, the same column density as Geballe & Oka. Since then, H₃⁺ has been detected in many other diffuse cloud sightlines, such as GCS 3-2,^[21] GC IRS 3,^[21] and ζ Persei.^[24]

Steady-state model predictions

To approximate the pathlength of H₃⁺ in these clouds, Oka^[25] used the steady-state model to determine the predicted number densities in diffuse and dense clouds. As explained above, both diffuse and dense clouds have the same formation mechanism for H₃⁺, but different dominating destruction mechanisms. In dense clouds, proton transfer with CO is the dominating destruction mechanism. This corresponds to a predicted number density of 10^-4 cm^-3 in dense clouds.

n(H₃⁺) = (ζ / k_CO)[n(H₂) / n(CO)] ≈ 10^-4 cm^-3

n(H₃⁺) = (ζ / k_e)[n(H₂) / n(C⁺)] ≈ 10^-6 cm^-3

In diffuse clouds, the dominating destruction mechanism is dissociative recombination. This corresponds to a predicted number density of 10^-6 cm^-3 in diffuse clouds. Therefore, since column densities for diffuse and dense clouds are roughly the same order of magnitude, diffuse clouds must have a pathlength 100 times greater than that for dense clouds. Therefore, by using H₃⁺ as a probe of these clouds, their relative sizes can be determined.

References