 | |
| Names | |
|---|---|
| IUPAC name Hydrogenonium | |
| Identifiers | |
3D model (JSmol) | |
| ChEBI | |
| ChemSpider | |
| 249 | |
| |
| |
| Properties | |
| H+3 | |
| Molar mass | 3.024 g·mol−1 |
| Conjugate base | Dihydrogen,H2 |
| Related compounds | |
Otheranions | Hydride |
Othercations | |
Related compounds | Trihydrogen |
Except where otherwise noted, data are given for materials in theirstandard state (at 25 °C [77 °F], 100 kPa). | |
Thetrihydrogen cation orprotonated molecular hydrogen (IUPAC name:hydrogenonium ion) is acation (positiveion) withformulaH+3, consisting of threehydrogen nuclei (protons) sharing twoelectrons.
The trihydrogen cation is one of the most abundantions in the universe. It is stable in theinterstellar medium (ISM) due to the low temperature and low density of interstellar space. The role thatH+3 plays in the gas-phase chemistry of the ISM is unparalleled by any otherpolyatomic ion.
The trihydrogen cation is the simplesttriatomic molecule, because its two electrons are the onlyvalence electrons in the system. It is also the simplest example of athree-center two-electron bond system.
H+3 was first discovered byJ. J. Thomson in 1911.[1] While using an early form ofmass spectrometry to study the resultant species ofplasma discharges, he discovered a large abundance of apolyatomic ion with amass-to-charge ratio of 3. He stated that the only two possibilities wereC4+ orH+3. Since the signal grew stronger in purehydrogen gas, he correctly assigned the species asH+3.
The formation pathway was discovered by Hogness & Lunn in 1925.[2] They also used an early form of mass spectrometry to study hydrogen discharges. They found that as the pressure of hydrogen increased, the amount ofH+3 increased linearly and the amount ofH+2 decreased linearly. In addition, there was littleH+ at any pressure. These data suggested theproton exchange formation pathway discussed below.
In 1961, Martinet al. first suggested thatH+3 may be present in interstellar space given the large amount of hydrogen in interstellar space and its reaction pathway wasexothermic (~1.5 eV).[3] This led to the suggestion of Watson and Herbst & Klemperer in 1973 thatH+3 is responsible for the formation of many observed molecular ions.[4][5]
It was not until 1980 that the first spectrum ofH+3 was discovered by Takeshi Oka,[6] which was of the ν2 fundamental band (see#Spectroscopy) using a technique calledfrequency modulation detection. This started the search for extraterrestrialH+3.Emission lines were detected in the late 1980s and early 1990s in theionospheres ofJupiter,Saturn, andUranus.[7][8][9] In the textbook by Bunker and Jensen[10] Figure 1.1 reproduces part of the ν2 emission band from a region of auroral activity in the upper atmosphere of Jupiter,[11]and its Table 12.3 lists the transition wavenumbers ofthe lines in the band observed by Oka[6] with their assignments.
In 1996,H+3 was finally detected in the interstellar medium (ISM) by Geballe & Oka in two molecularinterstellar clouds in the sightlines GL2136 and W33A.[12] In 1998,H+3 was unexpectedly detected by McCallet al. in a diffuse interstellar cloud in the sightlineCygnus OB2#12.[13] In 2006 Oka announced thatH+3 was ubiquitous in interstellar medium, and that theCentral Molecular Zone contained a million times the concentration of ISM generally.[14]
The three hydrogen atoms in the molecule form anequilateral triangle, with abond length of 0.90 Å on each side. The bonding among the atoms is athree-center two-electron bond, adelocalized resonance hybrid type of structure. The strength of the bond has been calculated to be around 4.5 eV (104 kcal/mol).[15]
In theory, the cation has 10isotopologues, resulting from the replacement of one or more protons by nuclei of the other hydrogenisotopes; namely,deuterium nuclei (deuterons,2H+) ortritium nuclei (tritons,3H+). Some of them have been detected in interstellar clouds.[16] They differ in theatomic mass numberA and the number ofneutronsN:
The deuterium isotopologues have been implicated in the fractionation of deuterium in dense interstellar cloud cores.[17]
The main pathway for the production ofH+3 is by the reaction ofH+2 andH2.[18]
The concentration ofH+2 is what limits the rate of this reaction in nature - the only known natural source of it is via ionization ofH2 by acosmic ray in interstellar space:
The cosmic ray has so much energy, it is almost unaffected by the relatively small energy transferred to the hydrogen when ionizing anH2 molecule. In interstellar clouds, cosmic rays leave behind a trail ofH+2, and thereforeH+3. In laboratories,H+3 is produced by the same mechanism in plasma discharge cells, with the discharge potential providing the energy to ionize theH2.
There are many destruction reactions forH+3. 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.
The significant product of this reaction isHCO+, an important molecule for interstellar chemistry. Its strongdipole and high abundance make it easily detectable byradio astronomy.H+3 can also react with atomicoxygen to formOH+ andH2.
OH+ then usually reacts with moreH2 to create furtherhydrogenated molecules.
At this point, the reaction betweenOH+3 andH2 is no longer exothermic in interstellar clouds. The most common destruction pathway forOH+3 isdissociative recombination, yielding four possible sets of products:H2O + H, OH + H2, OH + 2H, and O + H2 + H. Whilewater 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 ongrains is still considered the primary source of water in the interstellar medium.
The most common destruction pathway ofH+3 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 isH2 and H, which occurs roughly 25% of the time.[18]
The protons of[1H3]+ can be in two differentspin configurations, calledortho andpara.Ortho-H+3 has all three proton spins parallel, yielding a totalnuclear spin of 3/2.Para-H+3 has two proton spins parallel while the other is anti-parallel, yielding a total nuclear spin of 1/2.
The most abundant molecule in dense interstellar clouds is1H2 which also hasortho andpara states, with total nuclear spins 1 and 0, respectively. When aH+3 molecule collides with aH2 molecule, a proton transfer can take place. The transfer still yields aH+3 molecule and aH2 molecule, but can potentially change the total nuclear spin of the two molecules depending on the nuclear spins of the protons. When anortho-H+3 and apara-H2 collide, the result may be apara-H+3 and anortho-H2.[18]
Thespectroscopy ofH+3 is challenging. The pure rotational spectrum is exceedingly weak.[19] Ultraviolet light is too energetic and would dissociate the molecule.Rovibronic (infrared) spectroscopy provides the ability to observeH+3. Rovibronic spectroscopy is possible withH+3 because one of thevibrational modes ofH+3, the ν2 asymmetric bend mode (seeexample of ν2) has a weak transition dipole moment. Since Oka's initial spectrum,[6] over 900absorption lines have been detected in the infrared region.H+3 emission lines have also been found by observing the atmospheres of the Jovian planets.H+3 emission lines are found by observing molecular hydrogen and finding a line that cannot be attributed to molecular hydrogen.
H+3 has been detected in two types of theuniverse environments:jovian planets andinterstellar clouds. In jovian planets, it has been detected in the planets'ionospheres, the region where theSun'shigh energy radiation ionizes theparticles in the planets'atmospheres. Since there is a high level ofH2 in these atmospheres, this radiation can produce a significant amount ofH+3. Also, with abroadband source like the Sun, there is plenty of radiation to pump theH+3 tohigher energy states from which it can relax byspontaneous emission.
The detection of the firstH+3emission lines was reported in 1989 by Drossartet al.,[7] found in the ionosphere of Jupiter. Drossart found a total of 23H+3 lines with acolumn density of 1.39×109/cm2. Using these lines, they were able to assign a temperature to theH+3 of around 1,100 K (830 °C), which is comparable to temperatures determined from emission lines of other species likeH2. In 1993,H+3 was found inSaturn by Geballeet al.[8] and inUranus by Traftonet al.[9]
H+3 was not detected in theinterstellar medium until 1996, when Geballe & Oka reported the detection ofH+3 in twomolecular cloud sightlines, GL 2136 andW33A.[12] Both sources had temperatures ofH+3 of about 35 K (−238 °C) and column densities of about 1014/cm2. Since then,H+3 has been detected in numerous other molecular cloud sightlines, such asAFGL 2136,[20]Mon R2 IRS 3,[20]GCS 3–2,[21]GC IRS 3,[21] andLkHα 101.[22]
Unexpectedly, threeH+3 lines were detected in 1998 by McCallet al. in thediffuse interstellar cloud sightline ofCyg OB2 No. 12.[13] Before 1998, the density ofH2 was thought to be too low to produce a detectable amount ofH+3. McCall detected a temperature of ~27 K (−246 °C) and a column density of ~1014/cm2, the same column density asGeballe & Oka. Since then,H+3 has been detected in many other diffuse cloud sightlines, such as GCS 3–2,[21] GC IRS 3,[21] andζ Persei.[23]
To approximate the path length ofH+3 in these clouds, Oka[24] 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 forH+3, 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.
In diffuse clouds, the dominating destruction mechanism is dissociative recombination. This corresponds to a predicted number density of 10−6/cm3 in diffuse clouds. Therefore, since column densities for diffuse and dense clouds are roughly the same order of magnitude, diffuse clouds must have a path length 100 times greater than that for dense clouds. Therefore, by usingH+3 as a probe of these clouds, their relative sizes can be determined.
