Why exoplanets should have ionospheres and brown dwarfs chromospheres

Do exoplanets have an ionosphere? What does a brown dwarf need to form aurorae even without a companion? Isabel Rodriguez-Barrera and colleagues, including Christiane Helling and former member Craig Stark from the LEAP group, investigated whether it is possible to create a magnetized plasma, a medium composed of positive and negative charges but with an overall neutral electric property. The production of magnetized plasma would allow the creation of ionospheres and electromagnetic phenomena such as aurorae.

The ionosphere is the upper part of a planetary atmosphere created by the ionizing effects of stellar UV and X-ray radiation. Its importance inheres in atmospheric electricity and radio wave propagation, but also the shielding of the inner atmosphere from stellar UV radiation. The following study showed what conditions a planetary atmosphere (that is an atmosphere from a planet or a planet like object like a brown dwarf) should fulfil to produce an ionosphere or, in case of brown dwarfs, a chromosphere, such as the spectacular example of the solar chromosphere.

Radio, X-ray and Hα emission from brown dwarfs have been observed in the recent years (e.g. 2MASS J10475385+2124234 by Route & Wolszczan 2012; 2MASS J13153094-2649513AB by Burgasser et al. 2013). Similar observations are not yet available for extrasolar planets. In case of the Sun, observations of these emissions (radio, X-ray and Hα) are tracers of the solar chromosphere. These observations suggest that brown dwarfs contain ionized gas and host very strong magnetic fields, which are both needed to explain, for example, the radio emission. The aim of our study is to identify ultra-cool objects (with effective temperatures less than ~3000 K; Fig. 1.) that are most susceptible to processes leading to instabilities that trigger the emergence of strong plasma, a neutral state of matter composed of equal number of positive and negative ions.

Our theoretical work proposes a method of analysing the ionization and magnetic coupling state of objects with ultra-cool atmospheres. Our particular interest focuses on late M-dwarfs, brown dwarfs and giant gas planets.

Figure1. M-dwarfs, brown dwarfs and giant gas planets in comparison. Teide 1 is an example for a late M-dwarf, GD 165B for a cloud-forming brown dwarf of spectral type L, Gliese 229B is a cooler cloud-forming brown dwarf of spectral class T, and Jupiter is the example for a giant gas plane.

Figure 1. M-dwarfs, brown dwarfs and giant gas planets in comparison. Teide 1 is an example for a late M-dwarf, GD 165B for a cloud-forming brown dwarf of spectral type L, Gliese 229B is a cooler cloud-forming brown dwarf of spectral class T, and Jupiter is the example for a giant gas plane.

To determine the fraction of atmosphere that can be ionized, first Rodriguez-Barrera et al. consider thermal ionization only. Thermal ionization results from collisions between the gas particles according to the local gas temperature, therefore here we do not consider external ionizing affects from companions. Sources of possible irradiation are the host star in the case of planets, and a white dwarf in the case of a white dwarf – brown dwarf binary (for example WD0137-349B, Casewell et al. 2015). Such external affects can be later compared to the results of our reference study. We use the Drift-Phoenix model atmosphere grid where the local atmospheric structure is determined by the following global parameters: Teff (effective temperature), log(g) (surface gravity) and [M/H] (metallicity).

Rodriguez-Barrera et al. show that ultra-cool atmospheres with high Teff , or with high metallicity and low log(g) have large fraction of atmospheric volume where plasma processes occur, and so they are the best candidates for radio, X-ray and Hα emissions, observed from various objects as was mentioned above.

Figure 2. The volume fraction of the atmosphere that is thermally ionized, V^th_gas/V_atm, for f_e>10-7 and for M-dwarf, brown dwarf and gas giant planet atmospheres. f_e measures the extent to which a gas is ionized. (Rodriguez-Barrera et al. 2015)

Figure 2. The volume fraction of the atmosphere that is thermally ionized, V^th_gas/V_atm, for f_e>10-7 and for M-dwarf, brown dwarf and gas giant planet atmospheres. f_e measures the extent to which a gas is ionized. (Rodriguez-Barrera et al. 2015)

M-dwarfs have a considerable degree of ionization throughout the whole atmosphere, the degree of thermal ionization for a L-dwarf is low but high enough to seed other local ionization processes like Alfven ionization (see Stark et al. 2013) or electrostatic discharges, such as lightning, as seen on Fig. 2.

We show that the first criterion to form chromospheres, ionospheres or an aurora on an extrasolar planets or brown dwarf, namely a small but sufficient degree of ionization, can be fulfilled by thermal ionization alone without the need for additional processes. But is it possible to magnetise this ionized plasma? The second part of our study says yes! The results also give an idea of how well the different types of atmospheres can be magnetized (Fig. 3): The minimum threshold for the magnetic flux density required for electrons and ions to be magnetised is smaller than typical values of the global magnetic field strengths of a brown dwarf and a giant gas planet. This means the ionized plasma inside the atmosphere can be magnetised quite easily. A considerably lower magnetic flux density is required for magnetic coupling of the atmosphere in the rarefied upper atmosphere than in the dense inner atmosphere, meaning magnetising the plasma in the upper atmosphere is easier than in the inner parts of the atmosphere. The magnetic coupling works equally for electrons and atomic ions like Mg+ and Fe+ (Fig. 3).

Figure 3. The magnetic flux density required for electrons, Be (lower set of lines), and ions, Bi (upper set of lines), to be magnetically coupled to a background magnetic field in the object (B=10 G - giant gas planets (GP), B=103 G M-dwarfs (MD), brown dwarfs (BD); black horizontal/vertical lines). If B>Bi (or B>Be) the gas is magnetized by the external magnetic field B. (Rodriguez-Barrera et al. 2015).

Figure 3. The magnetic flux density required for electrons, Be (lower set of lines), and ions, Bi (upper set of lines), to be magnetically coupled to a background magnetic field in the object (B=10 G – giant gas planets (GP), B=103 G M-dwarfs (MD), brown dwarfs (BD); black horizontal/vertical lines). If B>Bi (or B>Be) the gas is magnetized by the external magnetic field B. (Rodriguez-Barrera et al. 2015).

To sum it up, our results suggest that it is not unreasonable to expect ultra-cool atmospheres (M-dwarfs and L & T brown dwarfs) to emit Hα or even in radio wavelength. We showed that, in particular, the rarefied upper parts of the atmospheres fulfil quite easily the plasma criteria despite having low degrees of ionization also in the case of giant gas planets. Therefore the results suggest that an ionosphere may emerge also in brown dwarf and giant gas planet atmospheres, and that the built-up of a chromosphere on brown dwarfs is likely too. Both effects will contribute to atmospheric weather features and to space weather occurrence in extrasolar, planet-like objects. An interesting result is that ultra-cool atmospheres could also drive auroral emission without the need for a companion’s wind (e.g. aurora on Earth triggered by solar wind) or an outgassing moon (e.g. aurora on Jupiter is triggered by its outgassing moon Io).

fig4

Figure 4. The dominating thermal electron donors for a subset of effective temperatures for log(g)=3,0 and solar element abundances, against the local gas pressure (Rodriguez-Barrera et al. 2015).

We further studied which of the gas species might be the best electron donors. Na+, K+ and Ca+ are the dominating electron donors in low-density atmospheres (low log(g), solar metallicity) independent of Teff. Mg+ and Fe+ dominate the thermal ionization in the inner parts of M-dwarf atmospheres. Molecules remain unimportant for thermal ionization. Chemical processes (e.g. cloud formation, cosmic ray ionization) that affect the abundances of Na, K, Mg, Ca and Fe will have a direct impact on the state of ionization in ultra-cool atmospheres.

For more details check out the original paper on ADS:

 Rodríguez-Barrera, I.; Helling, Ch; Stark, C. R. and Rice, A. M. 2015, MNRAS 454, 3977

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DRIFT-PHOENIX Atmosphere Models – Creating new worlds

As it was promised, here comes a review on the DRIFT-PHOENIX model atmospheres, which are used by our group. (See posts: 14 Mar, 30 Apr)

DRIFT-PHOENIX is a computer code that simulates the structure of an atmosphere including the formation of clouds. The code is part of the PHOENIX-code family (Hauschildt & Baron 1999; Baron et al. 2003) which is described by one of his main authors, Peter Hauschildt, as “… a general-purpose state-of-the-art stellar and planetary atmosphere code. It can calculate atmospheres and spectra of objects all across the HR-diagram ranging from main sequence stars, giants, white dwarfs, stars with winds, TTauri stars, over novae and supernovae, to brown dwarfs and extrasolar giant planets.”

DRIFT is a separate code that was coupled with PHOENIX (Dehn 2007; Helling et al. 2008a; Witte et al. 2009). DRIFT describes the formation of mineral clouds and allows to predict cloud details, like the size of the cloud particles and their composition (Woitke & Helling 2003, 2004; Helling & Woitke 2006; Helling et al. 2008b).

I. STELLAR ATMOSPHERE PART


The modelling of a stellar atmosphere structure is a standard task in astrophysical research. It is determined by three global parameters:

  • the effective temperature, T_eff [K], (represents the total flux of a star)
  • the surface gravity, log(g) [cm/s],
 (relates the mass and radius of an objects and, in many cases, serves as a reasonable age indicator)
  • a set of element abundances (H, He, O, C, Mg, Fe, Si, …), the so-called metallicity [M/H] 
(relies on primordial element abundances combining stellar nucleosythesis with the age of the universe before an object has formed, but also involves other dynamic effects like altitude-dependent gravitational settling or electromagnetic enrichment/depletion)

Given these global parameters, the atmosphere structure is the coupled solution of

(1) the hydrostatic equilibrium which provides the local gas pressure, p_gas [dyn cm^2],
(2) the energy transfer through radiation and convection which yields the local gas temperature, T_gas [K],
(3) the chemical equilibrium which provides the number densities of those ions, atoms, and molecules that compose the atmospheric gas.

The energy transfer through radiation is described by the radiative transfer equations, and the convective energy transfer is described by the Mixing Length Theory. Both produce a flux of energy through the atmospheric gas and the total flux through the atmosphere needs to be conserved as no heat nor radiation can be created inside the atmosphere itself. Radiation is transported through absorption and re-emission by the atmospheric gas. It is, therefore, essential to know how many absorbing ions, atoms and molecules are in the atmosphere (see (3)), and with which efficiency they absorb and re-emit photons. These chemical abundances, however, change with temperature and pressure, and are therefore interwoven with (1) and (2). The solution to this can only be found by an iterative numerical procedure. The solution to the problem is found if the iterative procedure fulfills some prescribed criteria to a certain precision.

The absorption and re-emission of photons by the atmospheric gas as part of a model atmosphere requires a set of material constants (wavelength dependent absorption coefficients for each individual atom and molecule and equilibrium constants for gas-phase chemistry). Without them, the atmosphere structure can not be calculated.

A more detailed insight into the theory of stellar atmospheres can be acquired from standard text books, like e.g. Mihalas (1978).

II. CLOUD FORMATION PART

In contrast to Earth, extra-solar atmospheres may not contain any condensation seeds. The cloud formation module DRIFT includes, therefore, a description of the formation of seed particles and their subsequent growth to macroscopic particles. DRIFT models the formation of (TiO2)_N-seed particles, and the mantle growth by the 7 [s]olids TiO2[s], Al2O3[s], Fe[s], SiO2[s], MgO[s], MgSiO3[s], Mg2SiO4[s] made of 6 elements and grown by 32 surface reactions (Witte et al. 2009). The cloud particles are composed of mixed materials, and their sizes change with atmospheric height. If a cloud forms, the atmospheric gas is reduced by those elements that are turned into the cloud particle. The cloud particles start to fall through the atmosphere. During this journey, they will grow in size but also change it’s material composition. The distance of their fall determines the cloud height.

Figure 1. Nucleation (seed formation), dust growth (and evaporation), gravitational settling (rain-out) and element replenishment are processes involved into the formation of a cloud. The inner part of an atmosphere is typically warmer than the outer part in a brown dwarf, and no cloud particles can form. Atmospheres of brown dwarfs and giant gas planets are convective (”boiling”) which provides the mechanisms for element replenishment. Original figure in Woitke & Helling (2004).

Figure 1. Nucleation (seed formation), dust growth (and evaporation), gravitational settling (rain-out) and element replenishment are processes involved into the formation of a cloud. The inner part of an atmosphere is typically warmer than the outer part in a brown dwarf, and no cloud particles can form. Atmospheres of brown dwarfs and giant gas planets are convective (”boiling”) which provides the mechanisms for element replenishment. (Woitke & Helling 2004)

The processes describing the cloud formation are (Woitke & Helling 2003, 2004; Helling & Woitke 2006; Helling et al. 2008a) (Fig. 1)

(4)  nucleation, growth & evaporation,
(5)  gravitational settling of cloud particles,
(6)  element conservation for those element involved in cloud formation, 
incl. convective element replenishment.

Process (4) requires the knowledge of the local gas temperature and of the number densities of those gas species involved into the cloud particle formation. Process (5) requires the knowledge of the local gas density (gas pressure), and process (6) requires the local gas velocity. In return, process (6) feeds back into the local gas-phase chemistry (3) with reduced element abundances, and (4)−(6) provide details for the opacity calculations as input for (2). Also the cloud formation part requires a set of material constants (e.g. material densities, molecular cluster data for nucleation) which are needed to calculate the cloud structure.

The text book by Gail & Sedlmayr (2014) provides a detailed derivation of the basic theory of dust formation in astrophysical objects. Helling & Fomins (2013) summarise additional developments for gravitational dominated atmospheres.

III. MODEL OUTPUT

A model atmosphere simulation provides detailed information for all the complexes (1)−(6). To keep the amount of numerical data sensible, only a subset of information is included into any output files. Table 1 summarises the most common output quantities. The easiest way to compare model results with observations of an astrophysical objects, like a brown dwarf, is compare the spectral energy distribution (Fig. 3), also called the spectrum, with the observed spectrum. The spectrum is the wavelength-dependent flux, F_λ, plotted versus the wavelength, λ.

Screen Shot 2014-06-30 at 17.08.56

Figure 2. Left: Drift-Phoenix (T_gas(z), p_gas(z))-structures for solar metallicity, log(g)=3.0 and effective temperatures, T_eff, as indicated. Right: Typical composition of clouds made of mixes of minerals that change throughout the cloud. The change in grain size is indicated (big grains / small grains).

Figure 2. Left: DRIFT-PHOENIX (T_gas(z), p_gas(z))-structures for solar metallicity, log(g)=3.0 and effective temperatures, T_eff, as indicated. Right: Typical composition of clouds made of mixes of minerals that change throughout the cloud. The change in grain size is indicated (big grains / small grains).

DF_Spec_metal

DRIFT-PHOENIX synthetic spectra for varying initial element abundances mimicking objects from the early universe for [M/H]=-6.0. The global parameters are given for each of the model results shown. (Witte et al. 2009).

4. DRIFT-PHOENIX MODEL ATMOSPHERE GRID

A model grid is a whole set of model atmospheres calculated for a whole set of global parameters (T_eff , log(g), [M/H]). The most up-to-date DRIFT-PHOENIX model atmosphere grid was presented in Witte et al. (2011). It contains an updated equilibrium chemistry (ACES). It spans over the following global parameter:

  • T_eff = 1000K . . . 3000K (in steps of 100K),
  • log(g)= 3.0 . . . 5.5 (in steps of 0.5),
  • [M/H]= −0.6 . . . + 0.3 (in steps of 0.3).

This set of global parameter (T_eff , log(g), [M/H]) allows combinations typical for:

  • M-dwarfs: Teff > 2700K, solar metallicity,
  • brown dwarfs: Teff < 2700K, log(g)≥4.5 for old brown dwarfs and log(g)<4.5 , for young brown dwarfs, solar metallicity and lower,
  • giant gas planets: Teff < 2700K, log(g)≤3.0, solar metallicity

The grid contains 527 model atmospheres in total ([M/H]=0.0: 148, [M/H]=-0.3: 130, [M/H]=+0.3: 121, [M/H]=-0.6: 128).

5. WHICH PAPERS TO CITE IF USING DRIFT-PHOENIX

Hauschildt & Baron (1999); Baron et al. (2003); Woitke & Helling (2003, 2004); Helling & Woitke (2006); Helling et al. (2008a, 2008b); Witte et al. (2009, 2011)

 

The PHOENIX code:

 http://www.hs.uni-hamburg.de/EN/For/ThA/phoenix/index.html

DRIFT-PHOENIX models are stored here:

http://svo2.cab.inta-csic.es/theory/newov/index.php

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Cosmic Rays enhance Prebiotic Chemistry on Sunless Worlds

In the International Journal of Astrobiology, Camille Bilger, Christiane Helling and Paul Rimmer (2014) presented a proof-of-concept on the potential effect of cosmic rays in the upper atmospheres of exoplanets (atmospheric pressures between 0.000001 bar and 0.000000001 bar, where 1 bar is the pressure of Earth’s atmosphere at sea level). We model the cosmic ray transport through the atmosphere of a planet with elemental composition and surface gravity similar to (but not the same as) Jupiter, if Jupiter were much farther away from the sun. This model atmosphere was produced using the Drift-Phoenix code (an upcoming post will be available soon on the Drift-Phoenix code).

Cosmic rays are charged particles (electrons, protons, bare nuclei) hurled through our galaxy at relativistic speeds by supernovae. When they strike the upper atmosphere of a planet, found to change its chemistry.

Cosmic Ray air shower

Figure 1. Cosmic Ray air shower

The combination of ultraviolet light from a star and cosmic ray ionization involves a delicate interplay between physics and chemistry, and is a hard problem to solve. It is simpler to consider cosmic ray chemistry on planets without daylight. These planets are often gas giants far from their host star, or rogue planets without a host star at all. These gas giants, like Jupiter, have an atmosphere made up not of mostly oxygen and nitrogen, but of mostly hydrogen, with a significant amount of nacient atomic hydrogen.

This is therefore a reducing atmosphere, and provides an ideal environment for making molecules believed to be important for the origin of life. Cosmic rays would ionize the molecular hydrogen, and this would make the atmosphere even more reducing.

Artistic impression of cosmic rays entering Earth's atmosphere. (Credit: Asimmetrie/Infn via CERN).

Figure 2. Artistic impression of cosmic rays entering Earth’s atmosphere. (Credit: Asimmetrie/Infn via CERN).

A reducing atmosphere is the standard initial chemical environment used in “origin-of-life” experiments, such as the Urey-Miller experiment. In the Urey-Miller experiment, an ionizing source in the form of an electrical discharge is initiated in a molecular gas, and so long as the atmosphere is reducing, prebiotic molecules are formed, including the twenty common amino acids found in living systems. If a reducing atmosphere, such as one dominated by oxygen and nitrogen, is used, the experiment produces no organic compounds.

The ion-neutral reactions made possible by cosmic ray ionization allow more of the hydrogen to be liberated from its molecular form, and increases the rate of reducing reactions. These reactions are found to be responsible for much of the prebiotic chemistry. Specifically, cosmic rays help to make biologically important molecules such as ammonia and acetylene. How much do cosmic rays help? They increase the abundance of some of these species by 10x or 100x in some cases.

Figure 3. Volume fraction of various species as a function of the gas pressure, p [bar] for the model atmosphere of a free-floating giant gas planet. The results assuming chemical quenching of C2H2 and C2H4 at height at ∼10−3 (dashed), and the results with cosmic ray ionization (solid) are all presented in this plot. A thick black horizontal line indicates the pressure above which termolecular (??) reactions may dominate.

Figure 3. Volume fraction of various species as a function of the gas pressure, p [bar] for the model atmosphere of a free-floating giant gas planet. The results assuming chemical quenching of C2H2 and C2H4 at height at ∼10−3 (dashed), and the results with cosmic ray ionization (solid) are all presented in this plot. A thick black horizontal line indicates the pressure above which termolecular reactions may dominate. (Rimmer et al. 2014, Fig. 4.)

The third picture in this article is from our paper, showing how much various molecules are enhanced or reduced because of cosmic rays. Some of the molecules that are enhanced, like ammonia, are key ingredients in the formation of the amino acid Glycine. These ingredients often must overcome a reactive barrier in order to form the amino acid, and overcoming this barrier may be made possible by electrostatic activation of ammonia (see our previous post here).

An image of the HR8799 planetary system from the December 2010 press release.

Figure 4. An image of the HR8799 planetary system from the December 2010 press release. (Credit: NRC-HIA, C. Marois, and Keck Observatory)

Our work may be relevant for directly imaged exoplanets orbiting HR 8799. Since these planets are ten times farther from their host star than Jupiter is from the sun, they are forever shrouded in the dark of night. It will be good to see how a star would change the chemistry, but this is a difficult problem. In order to find out what a star does in the upper atmosphere, it will be necessary to consider how the starlight passes through the atmosphere, and how that will affect thetemperature. The temperature will change the chemistry, and the chemistry will in turn affect how the starlight passes through theatmosphere. Much work still needs to be done to solve the problem for planets closer to their sun.

For more details check out the original paper on arXiv:

P. B. Rimmer, Ch. Helling and C. Bilger 2014, IntJAstrobio, 13, 173

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Electrostatic activation of prebiotic chemistry in substellar atmospheres

How do prebiotic molecules, necessary to the origin of life, form? What energies are needed for the formation of e.g. formaldehyde, ammonia or glycine? Do dust grains of exoplanetary atmospheres have a key role in these processes? These and similar questions were investigated by Dr. Craig Stark and his collaborators.

Dust is present all over the Universe, growing in a variety of diverse environments, for example in the atmospheres of gas giant exoplanets, where mineral dust clouds form, as earlier works of our group have demonstrated. In this atmospheric plasma the dust grains become charged and attract positive ions accelerated from the plasma. “The energy of the ions upon reaching the grain surface may be sufficient to overcome the activation energy of particular chemical reactions that would be unattainable via ion and neutral bombardment from classical, thermal excitation. As a result, prebiotic molecules or their precursors could be synthesized on the surface of dust grains that form clouds in exoplanetary atmospheres.”

Miller-Urey_experiment-en.svg

Figure 1. Miller-Urey experiment

The famous experiment of Stanley Miller and Harold Urey was one of the first to show how important electricity is in the synthesis of prebiotic molecules. (Fig. 1.) They showed that in a planetary atmosphere composed of H2, CH4, NH3 and H2O it is possible to form prebiotic amino acids and other biological molecules important for life if electrical discharge is present. However nowadays scientists believe that the atmospheric composition used in the above mentioned experiment does not correspond to the one existed on the young Earth, although a recent study showed that during volcano outbursts, where reduced gases and lightning processes are also present, it is possible to produce prebiotic molecules. As our paper says: “For atmospheres more representative of primitive Earth, no significant organic molecules are produced using electrical (sparking) discharges. However, the presence of hydroxy-acids in the famous Murchison meteorite indicate that the so-called Strecker amino-acid synthesizing mechanism (triggered by a Miller-Urey-type electrical discharge) may be responsible for the extraterrestrial synthesis of amino acids.”

Why may dust grains have a significant role in the occurrence of particular chemical reactions? Because in a plasma containing dust particles the absorption of different kind of species is electrostatically driven, which means that energies can easily exceed the activation energy required for the formation of prebiotic molecules.

“Consider a dusty plasma in the atmosphere of a substellar object such as a giant gas exoplanet. The dust particles will be negatively charged and as a result a plasma sheath (an electron depleted region) forms around the particle. As a consequence, the ionic flux at the grain surface increases as the plasma ions are attracted to and are accelerated towards the grain surface. Upon reaching the surface the ions have fallen through an electrostatic potential and have been energized. In comparison to the neutral case, the ionic flux is enhanced and the ionic energy amplified, increasing the probability that chemical reactions will occur and that reactions with higher activation energies are accessible. In this way, charged particle surfaces help catalyze chemical reactions otherwise inaccessible at such low-temperatures present in planetary atmospheres.”

In this paper we investigated the energization of ions as they are accelerated to the surface of a charged dust grain. We were mainly interested in the electrostatic activation of particular chemical reactions in the atmospheres of exoplanets.

Simulations were made using an example substellar atmosphere, which was created by the Drif-Phoenix atmosphere and cloud formation code. The atmosphere was defined by the following parameters: Surface gravity (log(g)) = 3.0; effective temperature (Teff) = 1500 K and solar metalicity ([M/H]=0.0).

Figure 2. was taken from the paper. As it says, the figure “shows the mean grain size  as a function of atmospheric pressure pgas. In the nucleation-dominated upper atmosphere (pgas ≈ 10-11 bar) seed particles form with a mean grain size <a> ≈ 10-7  cm. The dust particles gravitationally settle and grow as they fall, increasing in size. In the lower atmosphere (pgas ≈ 1 bar) the mean particle size is <a> ≈ 10-5 cm.”

Figure 2.  Mean grain particle size  as a function of gas pressure.

Figure 2. Mean grain particle size as a function of gas pressure (Stark et al. 2013)

As the electron temperature increases the local electrostatic field will rise as well which will result in that the ions accelerated by this field gain more energy. The more energy the ions gain the more likely to form molecules with higher activation energies.

We used the formation of glycine (NH2CH2COOH) to give an example “of the electrostatic activation of prebiotic chemistry on the surface of a charged dust grain.” The following chemical sequence shows the individual steps, which lead to the formation of glycine, with the formation energies (an indicator of the activation energy of the chemical reaction)

Screen Shot 2014-01-28 at 14.24.12

Figure 3. Synthesis of glycine

Figure 3. Synthesis of glycine

As our results show, “high in the atmosphere where pgas ≈ 10-15 bar, the ion temperature is ≈ 600 K and the resulting thermal energy of an ion is ≈ 0.08 eV, which is lower than the energies required to form the reaction products above. However when the electron temperature Te = 1 eV the ions are accelerated to Etot ≈ 7.8 eV which surpasses the required formation energies, increasing the likelihood that these reactions will occur. At lower atmospheric pressures where it is hotter, the thermal energy can reach ≈ 0.45 eV (≈ 5200 K) and there will exist a population of higher energy ions that, once neutralized on the surface, may be energetic enough to activate the required chemical reactions to form glycine.”

We showed in our paper that in atmospheric plasma where dust grains become negatively charged, ions can be accelerated to energies high enough to produce chemical reactions which would be inaccessible via classical thermal ion and neutral fluxes. As a consequence the creation of prebiotic molecules on grain surfaces may increase significantly.

The importance of this paper is that it “establishes the feasibility of the electrostatic activation of prebiotic chemistry. This idea can be developed to explicitly model the surface chemical kinetics, describing the incoming accelerated ions interacting with the grain surface. In this way, the effect of the plasma ionic species, the composition of the grain surface and the effect of the grain charge on the resulting surface chemical reactions can be quantified.”

For more details check out the original paper on arXiv:

C. R. Stark, Ch. Helling, D. A. Diver, and P. B. Rimmer 2013, IntJAstrobio 13, 165

Also look at the press release of the paper:

http://www.st-andrews.ac.uk/news/archive/2014/title,233471,en.php