L9m

flg. 12.20. Contour plot of the size distribution for the haze particles from Lavvas et al. 2007b. Size in ^m, number density in cm~3. Dashed curve corresponds to the optically effective radius of the particles. Numbers on contours correspond to the log10 number density. Box is Voyager I estimate.

physics. The interaction of solar radiation with the atmosphere (photochemistry) and haze determines the chemical species that are formed and the temperature structure. The chemical rates that depend on the temperature, control the pathways that lead to the formation of the haze precursors. The chemical growth of the latter, proceeds until a size beyond which the microphysical growth of the particles is more efficient. From this point on, the coagulation rate between the particles (which depends on the atmospheric temperature) defines their size and hence the optical properties of the haze layers. Finally, the interaction of the haze layers with the radiation field controls the solar radiation flux and the temperature structure. Under this self-consistent picture it is possible to validate the possible pathways through comparison with the observed temperature structure, the geometric albedo, the chemical structure and other observed parameters of the atmosphere. Such a model was developed by Lavvas et al. (2007a,b). The following paragraph presents some of the results of this simulation which are compared with the Cassini/Huygens observations.

Based on the chemical precursors of the suggested pathways the latter can be subdivided into four main families:

Pure nitrile polymers H2CN + HCN polymer CN + C4N2 ^ polymer CN + C2H3CN ^ polymer CN + C2N2 ^ polymer H2CN + CH2NH ^ polymer

Polyynes

C4H + C6H2 ^ polymer C6H + C4H2 ^ polymer C6H + C6H2 ^ polymer

Copolymers

HC3N + C4H polymer C4H2 + C3N ^ polymer HC3 N + C4H3 ^ polymer

Polyaromatics PAHs and N-PACs C6 H5 + C2H2 ^ polymer C6H5 + HC3N ^ polymer C6H5 + C6H6 ^ polymer.

The photochemical description of the atmospheric composition along with the above pathways, provide a vertical production profile for the haze monomers. This is presented in Fig. 12.19 where the contribution of each family is shown along with the cumulative production at each altitude. Below 300 km contributions from aromatics, pure nitriles and copolymers are comparable, while at higher altitudes copolymers dominate and nitriles become important above 800 km. The microphysical evolution of the produced monomers leads to the production of bigger particles, the final size of which is constrained only by the electrostatic repulsion among them. Figure 12.20 presents the calculated size distribution of the particles based on the latter production profile. The average size of the particles is constant below 200 km and between 0.1 and 0.2 ¡m.

Using measurements for the refractive index, the opacity of the produced particles can be calculated with the use of Mie theory (see §6.3). The resulting column haze opacity along with the vertical extinction profile at 0.5 ¡m, are presented in Fig. 12.21 and Fig. 12.22, respectively, and compared with the measured values retrieved by the DISR instrument on board the Huygens probe (Tomasko et al. 2005). The measurements correspond to Titan's atmosphere through which the probe descended, while the model results correspond to the disk average opacity necessary to match the geometric albedo. UV photons are absorbed at

Wavelength (nm)

flG. 12.21. The column absorption opacity for the produced haze particles based on the model of Lavvas et al. (2007b) along with the measured extinction by the DISR instrument on board the Huygens probe (Tomasko et al. 2005).

Haze extinction at 0.53 um (m"1)

flG. 12.22. Extinction profile of the haze particles at 0.53 pm compared with the DISR retrieved profile. (Lavvas et al. 2007b)

Temperature (K)

Fig. 12.23. Calculated temperature profile (solid line) from Lavvas et al. 2007b and retrieved temperature profiles of HASI (dashed line, Huygens probe) and CIRS (dotted lines, Cassini) instruments.

high altitudes while up to 0.6 ¡m haze opacity determines the altitude at which photons are absorbed. At longer wavelengths photons can reach the surface only inside the methane windows where photon reflection depends significantly on the surface albedo.

12.5.6 Thermal structure

The resulting temperature profile, based on the above description of haze production and evolution is presented in Fig. 12.23. For comparison the retrieved vertical temperature profiles from HASI (F. Ferri, private communication) and CIRS ( A. Coustenis, private communication) are presented. The bundle of CIRS profiles give the variation of the temperature structure at different latitudes from pole to pole. The variation of the temperature is rather small, only a few degrees for a large part of the atmosphere until the winter pole where the stratopause temperature increases and the stratosphere cools. This is related to the seasonal variation of the atmospheric structure, which enhances the production of haze particles in the winter pole. The calculated vertical profile, which corresponds to midlatitude conditions, provides a good fit to the observed profiles, close to the

Table 12.6 Species composition from Voyager and Cassini.

Voyager

Cassini

Species

IRIS (7°S)a

CIRS (15° S,160 km)b INMS

1200 km)c

Acetylene

C2H2

2.7 ±0.2x10-6

3.0x10-6

2.8x10"

4

Ethylene

C2H4

1.5 ±0.3x10-7

1.65x10-7

6x10-

3

Ethane

C2H6

1.65 ±0.1x10-5

1.5x10-5

1.2x 10-

4

Methylacetylene

C3H4

8.0 ±0.6x10-9

9.0x 10-9

4.0x 10-

6

Propane

C3H8

9.0 ±1.6x10-7

4.55x10-7

2.3x10-

6

Diacetylene

C4H2

1.7 ±0.2x10-9

1.8x 10-9

6.0x 10-

-5

Hydrogen cyanide

HCN

2.2 ±0.2x10-7

1.76x 10-7

2.0x 10-

4

Cyanogen

C2N2

<1x10-9

Cyanoacetylene

hc3n

<1.5x10-9

2x 10-

5

Ammonia

NH3

7x10-

6

Methyleneimine

CH2NH

<10-

5

Carbon dioxide

co2

1.5 ±0.15x10-®

a. Coustenis and Bézard 1995, b. Flasar et al. 2005, c. Waite et al. 2005, Vuitton et al. 2006

a. Coustenis and Bézard 1995, b. Flasar et al. 2005, c. Waite et al. 2005, Vuitton et al. 2006

HASI profile in the lower atmosphere and approaches the CIRS profile above the stratosphere.

In the stratosphere, the strong temperature increase is due to absorption of solar radiation by methane and most importantly, by the haze, similarly to the Earth's ozone layer. The haze has an 'antigreenhouse' effect as it is optically thin for thermal radiation but reflects solar radiation. In this way the heating from the absorption of the solar radiation is balanced partially by the escape to space of the thermal radiation. In the mesosphere, the strong negative temperature slope is due to thermal emission, mainly by methane and other hydrocarbons present in smaller amounts (C2H2, C2H6 and HCN) but enough to affect the thermal structure to a smaller degree. In the thermosphere, the photodissociation by solar UV radiation of the main species produces the temperature increase. The observed temperature oscillations are still under investigation.

12.5.7 Atmospheric chemistry

Table 12.6 presents a comparison of the chemical composition of Titan's atmosphere, based on measurements from the Voyager missions and the recent results of the Cassini/Huygens mission. The latter has verified the results of the former but also has provided new nitrile species that make Titan's atmospheric chemical laboratory even more interesting. The simulated chemical composition of the atmosphere from Lavvas et al. (2007a,b), for some of the main hydrocarbons (Fig. 12.24) and nitriles (Fig. 12.25) found in Titan's atmosphere, are presented together with the latest results from the CIRS and INMS instruments.

The most abundant hydrocarbon found in Titan's atmosphere after methane, is ethane (C2H6), which is mainly produced in the collisional addition of two methyl radicals

Fig. 12.24. Calculated vertical profiles of the main hydrocarbon species found in Titan's atmosphere. (Lavvas et al. 2007b)

INMS (Vuitton et al., 2006) • • IRAM (Marten et al., 2002)

— CIRS/Ltab (15 S) . . ,„„„ , ,_____ (Vinatier et al., 2006)

Boxes CIRS/Nadii (33 N) (Coustenis et al., 2007)

INMS (Vuitton et al., 2006) • • IRAM (Marten et al., 2002)

— CIRS/Ltab (15 S) . . ,„„„ , ,_____ (Vinatier et al., 2006)

Boxes CIRS/Nadii (33 N) (Coustenis et al., 2007)

10* 10* 10-7 Mixing Ratio

10* 10* 10-7 Mixing Ratio

10"3

Fig. 12.25. Calculated vertical profiles of the main nitrile species found in Titan's atmosphere. (Lavvas et al. 2007b)

The chemical destruction of methane molecules by the methylene radicals enhances the production of methyl radicals and hence the total ethane production

Because ethane's absorption cross-section falls in the same wavelength region as that of methane, its photolysis is constrained by the high amounts of the latter. Hence its destruction is mainly due to chemical reactions with radicals, of which the most important are with excited methylene

1CH2 + C2H6 ^ CH3 + C2H5, and ethynyl (C2H) that is produced by the photolysis of acetylene (C2H2)

Acetylene and ethylene (C2H4) are the next most abundant hydrocarbons after ethane and their formation is initiated in the upper atmosphere directly from the products of methane photolysis

2 1'3CH2 ^ C2H2 + 2 H 1,3CH2 + CH3 ^ C2H4 + H CH + CH4 ^ C2H4 + H .

Destruction of ethylene through photolysis acts as the major source of acetylene in the upper atmosphere

C2H4 + hv ^ C2H2 + 2H/H2, while in the lower atmosphere acetylene recycles back to ethylene through the vinyl radical (C2H3) according to the following scheme

C2H2 + H + M ^ C2H3 + M C2H3 + H + M -> C2H4 + M. C2H2 + 2 H ^ C2H4

The main loss of acetylene in the upper atmosphere is from reaction with methylene that leads to the formation of propargyl radicals (C3H3, which is the main precursor of benzene formation)

1CH2 + C2H2 ^ C3H3 + H, while in the lower atmosphere the main loss is due to photolysis. Acetylene's absorption extends well beyond methane's dissociation limit, up to ^230 nm and due to its high concentration it is the main absorber in this part of the spectrum (along with the haze). The photolysis rate is significant down to the stratosphere. Hence, there is large production of the radicals C2 and C2H, by acetylene's photolysis in this region, which, as in the case of ethane, leads to the catalytic destruction of saturated hydrocarbons through hydrogen abstraction. This mechanism affects methane

C2H2 + hv ^ C2 + H2 C2 + CH4 ^ C2H + CH3 C2H + CH4 —>■ C2H2 + CH3 CH4 CH3 + H, or

C2H2 + hv ^ C2H + H C2H + CH4 —>■ C2H2 + CH3 CH4 CH3 + H .

The overall destruction of methane is dominated by its catalytic destruction by the above two schemes along with loss due to excited methylene radicals. The three molecules C2H6, C2H4, and C2H2 together with molecular hydrogen correspond to the most abundant hydrocarbons observed in Titan's atmosphere since the Voyager missions. From these, more complex species are produced, such as propane (C3H8), diacetylene (C4H2) and benzene (CeHe), which finally lead to the formation of the haze precursors.

The basis of the nitrile chemistry is the hydrogen cyanide molecule (HCN) that is formed directly by the reaction of atomic nitrogen with the methyl radical produced by methane photolysis

This pathway is not dominant as the H2CN radical, also produced in the above reaction, reacts readily with atomic hydrogen to give HCN with an overall larger rate

N + CH3 ^ H2CN + H H2CN + H -> HCN + H2 N + CH3 HCN + H2 .

Since the majority of nitrogen atoms is produced in the upper atmosphere, the production of HCN has a maximum in this region. The photolysis of HCN leads to the production of cyano radicals (CN)

HCN + hv ^ CN + H , which have a double role in the atmospheric chemistry. Like C2H radicals, they enhance the catalytic destruction of molecular hydrogen and saturated hydrocarbons through the hydrogen-abstraction mechanism

CN + H2 ^ HCN + H CN + CH4 ^ HCN + CH3 CN + C2H6 ^ HCN + C2H5, leaving a free radical and a HCN molecule, but also lead to the production of other nitrile species in reaction with unsaturated hydrocarbons through the abstraction/addition mechanism. In the latter, a hydrogen atom is abstracted from the unsaturated hydrocarbon and the CN radical is added in its place. Typical examples are the formation of cyanoacetylene (HC3N) and acrylonitrile (C2H3CN)

Although acrylonitrile has not been observed in Titan's atmosphere before the Cassini/Huygens mission, it was believed to be formed and was taken into consideration in photochemical models that included nitrile chemistry. Recently, the analysis of the ionospheric composition by the Ion Neutral Mass Spectrometer (INMS) on the Cassini spacecraft has verified the presence of this species in the upper atmosphere together with other nitrile species.

Excited nitrogen atoms produced in the upper atmosphere from the photoion-ization of N2, lead to the production of other nitriles and amines. Reaction with methane has two possible products, imidogen (NH) and methylenimine (CH2NH) with the second having the largest yield

Methylenimine has been detected by the INMS instrument and is an intermediate in HCN production and is a possible candidate for haze formation.

Ammonia (NH3) is another species recently detected in Titan's atmosphere but also in the solid core of the haze particles. Its presence was anticipated due to the large abundance of nitrogen in the atmosphere and its formation was believed to be controlled by ionospheric chemistry. Yet, the measured abundance is significantly higher than suggested by models. Lavvas et al. (2007b) suggested another possible contribution to its formation from the neutral chemistry originating from the photodissociation of ethyleneimine (C2H5N), which is a possible product from the reaction of N(2D) with ethane

Finally, reaction of excited nitrogen atoms with ethylene and acetylene leads to the production of acetonitrile (CH3CN) and the CHCN radical, respectively

N(2D) + C2H2 ^ CHCN + H, while production of the latter is dominated in the same region by the replacement of H by CH in HCN

Yung (1987) suggested that the CHCN radical could be the basis for the production of cyanogen (C2N2) and dicyanoacetylene (C4N2)

N + CHCN ^ C2N2 + H CHCN + CHCN ^ C4N2 + H2, from which the first has been detected in Titan's atmosphere but the second, although not detected, is included in photochemical models since it is believed to be a possible precursor of haze formation.

flg. 12.26. Low-resolution infra-red spectra of Venus, Earth and Mars, based on actual observations from spectrometers on NASA spacecraft. In the next few decades, we should have spectra of this quality from planets orbiting other stars, and from them we should be able to say whether ozone (hence oxygen) and water are present, and possibly to estimate the surface pressures and temperatures.

flg. 12.26. Low-resolution infra-red spectra of Venus, Earth and Mars, based on actual observations from spectrometers on NASA spacecraft. In the next few decades, we should have spectra of this quality from planets orbiting other stars, and from them we should be able to say whether ozone (hence oxygen) and water are present, and possibly to estimate the surface pressures and temperatures.

flG. 12.27. The calculated infra-red spectrum of an Earth-like planet around a Sun-like star, as viewed by a large, but feasible, multimirror space telescope from a distance of 10 parsecs. The dashed line is the blackbody spectrum of a planet with no atmosphere at 300 K, the solid line is the signal from the planet, and the stepped line is the simulated spectrum, with noise sources included, binned at the resolution of the spectrum. (After ESA)

flG. 12.27. The calculated infra-red spectrum of an Earth-like planet around a Sun-like star, as viewed by a large, but feasible, multimirror space telescope from a distance of 10 parsecs. The dashed line is the blackbody spectrum of a planet with no atmosphere at 300 K, the solid line is the signal from the planet, and the stepped line is the simulated spectrum, with noise sources included, binned at the resolution of the spectrum. (After ESA)

12.6 Extrasolar planets

Because our own Solar System exists, and given the size of the Universe with an estimated 50 billion galaxies, some containing thousands of billions of stars, there is a high probablility that planetary systems exist in large numbers. Some of these have already been detected, more than 100 at the time of writing, although with a bias towards those that contain very large planets - more like Jupiter than Earth - that are easier to detect. It is only a matter of time until we are able to detect planets similar to the Earth, orbiting stars that resemble the Sun in size and luminosity. Very large telescopes, floating in space, have already been designed that have the capability to obtain infra-red spectra of these planets, allowing techniques like those described in Chapter 9 to be employed to infer the existence of atmospheres, and their compositions and temperatures. Assuming reasonable technical progress, and the expenditure of sufficient funds, simple models of the climate on a number of these new Earths may be available to us in a few decades from now. It is likely that these will be sufficiently detailed to reveal the frequency with which Earth-like life accompanies an Earth-like environment.

Figure 12.26 shows what Venus, Earth and Mars would look like when viewed from outside the Solar System with a spectrometer attached to a suitably large telescope. The dominant feature in all three is the carbon dioxide band at 15 ¡m, while the 6.3 and 9.8 ¡m water and ozone bands only show up strongly on the Earth. Figure 12.27 is a simulation of how the terrestrial spectrum would appear from a distance of 10 pc through the DARWIN telescope currently planned by the European Space Agency. An 800-h (33 days) exposure by Darwin is required to obtain a spectral resolution of 0.5 pm per pixel. This is very low resolution, but is just enough to permit the detection of water and ozone with reasonable certainty. This in turn might be a strong indication of the presence of active, oxygen-producing life on the planet.

12.7 Bibliography

12.7.1 Notes

For the origin and evolution of planetary atmospheres see text by Lewis and Prinn, and for the the Earth's atmospheric evolution see texts by Holland, and by Walker.

For the early work on the origin of life on Earth see Miller, and Miller and Urey.

For anoxic conditions on the primitive Earth see early work of Cloud; Holland; and review of Kasting and references therein. For the role of organic burial on atmospheric oxygen levels see Holland et al.; Des Marais et al.

See the early work on chert and phosphate measurements by Knauth and Epstein; Karhu and Epstein that suggested 'warm' surface temperatures on the primitive Earth.

For isotopic and geochemical analyses that support growth models of continental crust formation see references in Carver and Vardavas.

For geological/biological conditions in the Late Proterozoic see Knoll; and Des Marais. For the evolution of the continental crust see Taylor and McLennan.

For references related to Titan's atmosphere, see for example; Yung et al., Couste-nis et al., McKay et al., Rannou et al., Yelle et al., Imanaka et al., Fulcignoni et al., Israel et al., Wilson and Atreya. See also the text by Coustenis and Taylor.

12.7.2 References and further reading

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