The upper atmospheres of all three terrestrial planets are slowly evaporating into space, with the highest loss rate for the lightest atom, hydrogen. At Venus and Mars, solar UV radiation penetrates into the middle atmospheres, where photodissociation of H2O by solar UV radiation releases H and O, which diffuse into the upper atmosphere and eventually reach the exobase. The H atoms may be lost into space, while relatively fewer O atoms escape, mainly due to non-thermal processes. There are several lines of evidence for the loss of large amounts of water by this process over geologic time on Venus and Mars, including the elevated ratios of atmospheric D/H compared with terrestrial standard mean ocean water (SMOW). From the higher evaporation rate of H compared with D due to the atomic mass difference, the present D/H ratio can be related by models to the total amount of water lost into space. By contrast, on the Earth the atmosphere is replenished by geologic activity, and the SMOW D/H ratio is unmodified by escape . At Mars, with its weak gravity, a relatively large amount of gas is believed to have been lost over the last 4 Gyear, corresponding to a potentially substantial surface depth of water. In the case of Venus, a warm early greenhouse would have raised H2O into the middle atmosphere where it would have been subject to photodissociation by solar UV much more than on the Earth.
The importance of the history of water on Mars is indicated in NASA’s Mars research program, which is focused on the theme “Follow the Water”. A complete understanding of the history of water on the terrestrial planets requires an understanding of the evolution of water on all three planets, Venus, Earth, and Mars, and the very different conditions on the three planets provides a solid testbed for theories. Nonetheless, there remains a poor understanding of the present-day total rate of atmospheric escape from Mars and Venus, and even larger uncertainties in the relative importance and efficiency of various processes. Understanding the details of escape today is a requirement to be able to extrapolate into the past to learn the history of venusian and martian water. The selection of two Mars escape missions for the next Mars Scout reflects the importance of this scientific question, and we here propose an inexpensive way to perform a key measurement of the atmosphere of Venus.
The atmosphere of Venus is known to have undergone substantial evolution over geologic time. Evidence for this includes the present remarkable contrast between Venus’ atmosphere and the Earth’s: Venus has very little water, a 95% CO2 atmosphere, a surface temperature of 750 K, and a surface pressure of 90 bar. The early venusian atmosphere is thought to have undergone either a moist or runaway greenhouse heating episode to produce these conditions, and this would have included hydrodynamic escape of light gases from the upper atmosphere to deplete possibly as much as an ocean of water. Support for this scenario comes from the measured ratio of D/H in Venus’ atmosphere of roughly 1.6% from the Pioneer Venus mass spectrometer, orbiting ion mass spectrometer (OIMS) data, and IR spectra of the nightside atmosphere. This large enhancement over cosmic abundances is consistent with the loss of an ocean’s worth of H2O over geologic time, with the H escape exceeding the D escape due to the mass difference. By determining the abundance of different isotopes today we can place constraints on the nature of the escape process in Venus’ past. The present ratio of D/H and any variation with altitude are also important to understand the physics and chemistry of the present rate of escape from Venus’ upper atmosphere, which may be dominated by the solar wind/ionosphere interaction in the absence of an intrinsic magnetic field.
The general scenario for the escape of water involves details of the physics and chemistry in the atmosphere of each planet. Starting with photodissociation of water by solar UV radiation, H and D end up diffusing upward mainly in the form of H2 and HD, and are converted back to atomic H and D in the upper atmosphere. The ratio D/H is not expected to be constant in all forms of H and D, nor at all altitudes, due to a series of factors. For example, there are different rates of both condensation and photodissociation of HDO and H2O, which is well known from studies of the Earth’s atmosphere. It is important to measure the ratio D/H in different levels of each planet’s atmosphere to understand the details of how the atoms move from surface water to the exobase.
With a ratio of D/H of 2.5% in the bulk atmosphere, it should be relatively easy to detect the D Ly α emission from Venus’ upper atmosphere, although this has never been done before. The D line will be optically thin and resonantly scatter the broad solar emission 0.33 Å blueward of the H Ly α line. Based on a measured 20 kR sub-solar brightness for H Ly α, the D Ly α line is expected to be of the order of 1 kR. The best observational upper limit to date is 300 R at 2σ from IUE observations. This suggests that something is wrong with our picture of Venus’ upper atmosphere. With 3 independent measurements of the bulk D/H ratio, and now an enhancement in the middle atmosphere, it is surprising that D/H in atomic form would be underabundant. The D/H ratio clearly varies with altitude, and the previously assumed number for H density (based on Pioneer Venus measured H Ly α brightnesses and geometry) may have been too high or time variable. The bulk of Ly α resonant scattering of solar emission occurs within one scale height above 110 km, a lower boundary determined by absorption of UV sunlight by CO2, whereas the OIMS data can measure D and H densities only above 170 km altitude. A variation in D/H ratio with altitude seems the most likely possibility, but we must perform the observation to determine this. The correct altitude distribution of D and H is important to be able to understand the present nature of the escape processes from Venus’ upper atmosphere.