THE INTERACTION OF THE SOLAR WIND WITH VENUS

Venus has been more frequently studied by space missions than any other planet. Consequently, we know more about the solar wind interaction with Venus than with any planet except the Earth. Unlike the Earth, Venus's ionosphere, not a magnetosphere deflects the solar wind flow However, as on the Earth, this deflection is accomplished with the formation of a bow shock, which heats and compresses the solar wind flow The shock is both closer to the planet and weaker than would be expected for an ideal gas dynamic interaction with a perfectly reflecting obstacle. The ionized flow of the magnetosheath can interact directly with the neutral atmosphere through charge exchange and photoionization. The former process removes momentum from the flow; both processes add mass to the solar wind, since the high altitude neutral atmosphere is mainly hot oxygen, not hydrogen. Finally, Venus, like Earth, has a magnetotail but not for the same reason. The mass loading of the flow in the magnetosheath slows the transport of magnetic flux tubes past the planet, while the ends of the tubes continue to travel rapidly in the solar wind. Thus the planet accretes interplanetary magnetic flux. This process is the dominant source for the magnetotail flux, not unipolar induction, although the latter process is present at least when the solar wind dynamic pressure is high. On the whole, the solar wind interaction with Venus is more comet-like than Earth-like.

Venus is the Earth's closest planetary neighbor, both in distance and in physical properties. The closeness in space has invited frequent exploration: the Venera series of spacecraft, three Mariner flybys and the Pioneer Venus orbiter and probe mission. The similarity in size, density, and a dense atmosphere has invited speculation that Venus is similar to the Earth in other ways as well, but we find many differences. The most important difference in interaction of the solar wind is the weakness or absence of a planetary magnetic field on Venus, which in this respect is more akin to a comet than to the Earth. We would expect Venus's magnetic field to be weaker than the Earth's simply because Venus's rotation rate is much lower (Busse 1976). Venus rotates on its axis with respect to the stars once in 243 d; the earth's rotational period is only 23 hr 56 min 4 s. However, Venus's field is much weaker than would be extrapolated from comparison with the terrestrial field, implying that there are probably important differences between the interiors of Earth and Venus, which affect Venus's ability to generate a self-sustaining magnetic dynamo.

Venus orbits the Sun every 225 days; the Earth orbits the sun every 365 days. The beat between these two periods is 586 d or 19.3 mon, so there are nearly two equi-spaced launch opportunities every three years. The first of these to be used was by Mariner 2 in 1962 which flew by Venus with a distance at closest approach of 6.6 Venus radii (R). The next Venus data were not obtained until three launch opportunities later, when in October 1967 both Venera 4 and Mariner 5 probed the planet. At the next opportunity, duplicate missions Veneras 5 and 6 were launched. Veneras 7 and 8 did not study the solar wind interaction and Mariner 10, whose primary mission was to explore Mercury, passed by Venus at too large a distance to provide new insight into the nature of the interaction. Hence four opportunities passed before new results on solar wind interaction were obtained in October 1975 with the Venera 9 and 10 orbiters. Repeated sampling with orbiters was a great advance over previous missions, and simultaneous data from two spacecraft provided the capability to unambiguously separate temporal from spatial effects. The next opportunity in May 1977 was skipped, but in December 1978 Venus was visited by a flotilla of spacecraft: the Pioneer Venus orbiter and probe, which consisted of a bus and four atmospheric probes, and Veneras 11 and 12, each having a flyby and a lander. Of the six U.S. spacecraft, however, only the orbiter provided data on the solar wind interaction and there have been no new results on the interaction from the four Soviet spacecraft, although the flyby craft did contain sophisticated plasma instrumentation for solar wind studies. These missions and their arrival dates and trajectory characteristics are listed in Table I.

The number of missions devoted at least in part to studying the interaction of the solar wind with Venus would suggest that much is now known about the interaction of the solar wind with Venus. While our understanding of some aspects of the interaction has in fact become quite extensive, in other aspects we have only rudimentary knowledge. Despite all this activity, each mission to Venus has involved some sort of compromise from the ideal mission for solar wind studies, leaving gaps in our knowledge. For example, the Pioneer Venus orbiter depends on the stability of solar wind conditions and the repetition of features to separate temporal from spatial variations. The Venera 9 and 10 orbiters permitted this separation of factors but had a limited number of simultaneous transmissions and no in situ ionospheric data. Throughout this review we caution the reader as to the limits of our present understanding in each area.

Fig. 1. Basic features of the solar wind interaction with Venus. The solar wind is deflected around Venus by the planetary bow shock. The obstacle is the planetary ionopause; the interface between the magnetic field which piles up in front of the planet and the ionosphere is called the ionopause. At times of low solar wind dynamic pressure the ionosphere is unmagnetized except for the occurrence of twisted filaments of field called flux ropes. Behind the planet a tail is formed by flux tubes which become filled with plasma on the day side and convect around the planet together with some flux tubes which penetrate the ionosphere and become hung up there.

The basic features of the interaction are now well understood, as shown in Fig. 1. The planetary ionosphere presents an obstacle to the magnetized solar wind flow. The supersonic, or more correctly, supermagnetosonic, solar wind is deflected about the planet by a collisionless bow shock in much the same way as the solar wind is shocked and deflected about the terrestrial magnetosphere. Finally, in apparent analogy with the Earth's magnetosphere but more correctly in analogy with cometary ion tails, Venus has its own magnetotail.

In this chapter we will consider fast the nature of the obstacle to the solar wind flow, and review the limits on the possible size of any planetary magnetic moment. We will also examine the early models of the interaction to compare with our present understanding, and discuss the physical processes occurring at the ionopause, the interface between the shocked solar wind and ionospheric plasmas. Then we will examine the wake and tail region behind Venus, including the question of ion pickup by the solar wind and mass loss. Finally, we will move out from the planet to consider questions such as the location and shape of the bow shock and the effects of Venus on the solar wind beyond the bow shock. Other viewpoints on some of these matters are represented in Chapters 26 and 27 by Cloutier and by Gringauz.

At the Earth the solar wind is deflected far above the atmosphere by a strong intrinsic magnetic field generated by a dynamo deep in the interior of the Earth. This shielding by the Earth's magnetic field does not prevent solar wind energy from coupling to the Earth's atmosphere. The aurora is one manifestation of this coupling. However, the solar wind does not directly interact with the atmosphere and current rates of loss to the solar wind are thought to play a minor role in the evolution of the terrestrial atmosphere. A similar type of obstacle to the solar wind is found at Mercury, Jupiter and Saturn; we can generally understand the behavior of their magnetospheres in terms of our terrestrial experience.

The earliest measurements suggested that Venus's magnetic field was far weaker than Earth's and possibly even absent, although the latter suggestion evoked some controversy. If the magnetic field were so weak that the solar wind interacted directly with the atmosphere, then a qualitatively different interaction would exist, one more like the interaction of a comet with the solar wind. The fact that the controversy on whether Venus had an intrinsic or induced magnetosphere raged for close to 20 yr shows the limitations of flyby missions for gathering in situ data, especially in situations involving variations in latitude, longitude, altitude, and time.

A. The Search for a Magnetic Moment

Mariner 2 was the first successful mission to another planet, passing within 6.6 radii (R) of the center of Venus in December 1962. It carried a triaxial fluxgate magnetometer (Smith et al. 1963, 1965) a solar wind probe (Neugebauer and Snyder 1965) and an energetic particle experiment (Frank et al. 1963). None of these instruments detected any evidence of the solar wind interaction with the planet, implying that the magnetic moment of Venus was < 1/20th that of the Earth.

On 18 and 19 October 1967 Venera 4 and then Mariner 5 reached Venus. Venera 4 carried a triaxial fluxgate magnetometer (Dolginov et al. 1968) and four charged particle traps (Gringauz et al. 1968). Venera 4 transmitted ion and magnetic field data until it reached an altitude of 200 kin. A bow shock was reported in both the magnetic field and ion data. Assuming that their data contained no offsets and noting that there was little change in the magnitude of the field with altitude, Dolginov et al. used the last field measurement to calculate an upper limit to the Venus moment of 1/3000th of the terrestrial moment. Later Dolginov et al. (1969) compared their entry data with simultaneous Mariner 5 data obtained
80 (R) away in the solar wind. Figure 2 shows this comparison.

Fig. 2. The magnetic field measured by Venera 4 and Mariner 5 during the Venera 4 encounter with Venus (Dolginov et al. 1969).

Both spacecraft are in the solar wind simultaneously at the beginning of this figure. The traces have been aligned for the expected delay time from Venera to Mariner. The fields at the two spacecraft are quite different. Such differences are rare but not unknown over similar baselines at 1 AU (Russell et al. 1980c), where the time delay is quite variable also. Thus we should not be surprised at these differences, but they could also be instrumental in origin. The principal point of this comparison is that the magnitude of the interplanetary magnetic field as measured by Mariner and the magnitude of the shocked magnetosheath field near the planet both stayed nearly constant as Venera 4 approached within 200 km of Venus. Thus Dolginov et al. reaffirmed their conclusions on the weakness of the magnetic moment of Venus. Note, however, that the vector components of the field change significantly as the planet is approached.

The next day Mariner 5 flew by Venus, coming within 1.4 R of the axis of the optical shadow and within 1.7 R of the center of the planet. Mariner 5 carried a triaxial vector helium magnetometer, a solar wind probe (Bridge et al. 1967) and an energetic particle detector (Van Allen et al. 1967). Figure 3 shows the trajectory of Mariner 5 and the plasma and magnetic field magnitude observed during the encounter period. Superimposed on these data are dashed lines corresponding to the numerical solution of the problem of solar wind interaction with an impenetrable object the size of the ionosphere when the interplanetary magnetic field is aligned with the flow (Rizzi 1971). Between points 2 and 3 the velocity observed is less than that predicted by the nondissipative hydrodynamic model, suggesting the presence of a viscous boundary layer. Such viscous boundary layers have been proposed by Perez-de-Tejada and coworkers (Perez-de-Tejada and Dryer 1976; Perez-de-Tejada et al. 1977). However, momentum is removed from the solar wind protons in other ways. The first is charge exchange, in which a fast ion becomes a fast neutral and dumps its momentum into the upper atmosphere; another example is photoionization of neutral oxygen in the boundary layer; the oxygen ions become accelerated while the protons become decelerated. These processes will be discussed at greater length in Secs. LB and II. C below. Later more sophisticated analysis of the solar wind data reaffirms these conclusions (Shefer et al. 1979).

Fig. 3. The number density, solar wind velocity, and magnetic field strength during the Mariner flyby. On the right is the Mariner 5 trajectory in solar cylindrical coordinates, in which the distance from the center of the wake is plotted versus the distance along the Sun-planet line (Rizzi 1971). The dashed lines show the hydrodynamic solution for a specific set of ionospheric and solar wind parameters.

The magnetic field strength at point 3 in Fig. 3 increases above the value of the field in the surrounding regions and becomes quite steady for a few minutes. This point is also the closest approach to the axis of the optical shadow.

Fig. 4. The magnetic field measured by Mariner 5 and Venera 4 plotted along their trajectories in solar cylindrical coordinates. For Mariner 5 the 3-dimensional trajectory information is used. For Venera 4 we have only the ecliptic plane projection. Closed triangles mark magnetopause traversals; dots indicate the boundary of a second current layer, possibly the rarefraction wave; the S marks the location of the Mariner 10 shock crossing.

Figure 4 shows the Mariner 5 and Venera 4 magnetic field plotted along the trajectory in a coordinate system that assumes cylindrical symmetry about the Sun-Venus line (Russell 1976a). Near closest approach to the tail axis the field lines are stretched out in a tail-like fashion both on Mariner 5 and Venera 4. The boundaries that one would identify as magnetopause crossings into a magnetotail-like field are consistent in location with the boundary of the tail seen in Pioneer Venus data (Russell et al. 1981). The direction of the tail field at Venera and Mariner is opposite. Mariner 5 passed through the northern-lobe of the Venus wake. However, the Venera 4 bus also entered in the northern hemisphere (Dolginov et al. 1978). Thus these data argue against the existence of an intrinsic planetary magnetic field, and are not consistent with the existence of an intrinsic field as large as the upper limit of 6.5 X 10²² Gauss-cm³ which had been advocated by Russell (1976b).

The next measurements, on Venera 6, returned data only on the distant bow shock (Gringauz et al. 1970). The next set of data came from the magnetometers on the Venera 9 and 10 orbiters which arrived at Venus on 22 and 25 October 1975. Their orbits had periods of slightly > 48 hr, periapsides of ~ 1500 km and inclinations of ~ 35^o These satellites made frequent passes through the wake region, crucial to probing for weak fields. The initial results suggested the intrinsic field to be weak (Dolginov et al. 1976).

Veneras 9 and 10 clearly observed a magnetic tail extending to large distances from the planet. The question to be settled was whether this long magnetic tail was due to sources intrinsic to the planet as concluded by Dolginov et al. (1978), or whether it was a manifestation of the interaction of the magnetized solar wind with the planet as concluded by Yeroshenko (1979). As laboratory simulations showed, it is certainly possible to have a long magnetic tail behind an unmagnetized obstacle (Dubinin et al. 1978). Despite the fact that simultaneous data from Veneras 9 and 10 could help distinguish the always vexing temporal and spatial variations, the amount of simultaneous coverage was limited. Furthermore, the Vcnera spacecraft were not triaxially stabilized in most cases when they were far from the planet; they were allowed to spin slowly about an axis close to the solar direction. Hence, when Pioneer Venus approached Venus in late 1978 the question of the possible size of any intrinsic Venus magnetic moment was still unsettled. Dolginov et al. (1978) felt Venus had an intrinsic moment of about 1/4000th of the terrestrial field. Yeroshenko (1979) felt that this number was an upper limit and that all observed field variations were due to the solar wind interactions.

The Pioneer Venus Orbiter (PVO) was placed into a highly elliptical near polar orbit on 4 December 1978. The initial periapsis altitude was 378 km at solar zenith angle 63^o . Periapsis altitude lowering began almost immediately and soon the satellite was probing deeply into the dayside ionosphere, where little magnetic field was observed (Russell et al. 1979a). Except for the occasional orbit during which the ionosphere was compressed to low altitudes by the solar wind, the field strength in the ionosphere was less than a few nT. One possibility was that the convection rate of the highly conducting dayside ionosphere was much greater than the diffusion rate of planetary field up into the ionosphere from below. Thus, it was of interest to see if there was a buildup of magnetic flux on the night side. In slightly > 2 mon periapsis had moved to the center of the nightside ionosphere. The nighttime field was generally similar to the day side. Often there were very low field values throughout the night ionosphere as in Fig. 5a; sometimes the field was moderately high (tens of nT) and steady as in Fig. 5b.

Fig. 5. Magnetic field in Venus wake plotted in solar cylindrical coordinates. (a) PVO orbit 65 2/7/79; (b) PVO orbit 72 2/14/79.

Occasionally there werelimited regions of radial field extending apparently through the base of the ionosphere, but the field mainly wrapped around the nightside ionosphere as if the solar wind magnetic field had draped around the planet and was closing behind the obstacle (Russell et al. 1979c, 1980b). There was no obvious sign of a planetary magnetic field. However, it was possible that there was some net flux out of the planet masked by the external currents and their fluctuations. Thus, the average magnetic field over the nightside ionosphere was calculated and inverted to find the best fit for magnetic moment. Using data from two nightside passes the upper limit to the magnetic moment was found to be 4 X 10^-5 of the terrestrial moment or <3 X 10²¹ Gauss-cm³ (Russell et al. 1980a).

The magnetic field strength necessary to deflect the solar wind at 0.72 AU is ~ 100 nT. If the solar wind doubles the planetary field by compression in the interaction, this deflection can be accomplished above the surface of the planet by a magnetic moment only 1/360th of the terrestrial moment. Small as this may be, the actual measured moment is < 1/25000th of the terrestrial moment, and we can conclude that at the present time intrinsic magnetic fields play no significant role in the solar wind- Venus interaction. Venus presents a very different type of obstacle to the solar wind than has been probed at the other planets.

B. Non-Magnetic Barrier Models

If a superconducting ball is placed in a uniform external magnetic field, the magnetic field lines bend so as to exclude the volume occupied by the ball as shown in Fig. 6.

Fig. 6. Magnetic field lines around a superconducting sphere.

This is the simplest model of the Venus solar windionosphere interaction. However, the region external to the highly (but not infinitely) conducting ionosphere is not a vacuum, but a flowing plasma. Thought of in magnetohydrodynamic terms (as an electrically conducting magnetized fluid) the solar wind has both mass and momentum and an electric field. This flowing plasma has to be deflected by the planet by some pressure wave. The flow introduces day-night asymmetry, and unless the field is flow-aligned there will be asymmetry about the flow direction as well. This superconductor analogy is useful but is only an approximation to the condition of the Venus ionosphere. Furthermore, there are the complications of nonmagnetohydrodynamic processes such as charge exchange and photoioniza- which we ignore for the moment, but will discuss later.

A superconducting or infinitely conducting ionosphere will deflect the solar wind flow around it. The nature of this interaction is quite complicated, involving the creation of a bow shock to heat and deflect the flow. This problem has not been solved analytically, although computer codes have long existed (Spreiter et al. 1966) which treat well the unmagnetized solar wind and the case of field-aligned flow (Rizzi 1971). The magnetization of the solar wind creates two problems. First, magnetic stresses affect the flow (cf. Walters [ 1964] for a discussion of these effects at the Earth); Spreiter et al. (1966) did not take into account these effects in their computations. Second, the flowing magnetized solar wind has an electric field. This was realized as important by all early investigators (Dessler 1968; Sonett and Colburn 1968; Cloutier et al. 1969; Cloutier 1970) but the effect of this electric field depended critically on the nature of the obstacle. It is important at this point to distinguish between two types of induction. The most familiar type of induction is that associated with a time varying magnetic field, as seen in our case in the reference frame of the planet. This can be caused by propagating waves or by convected structure in the solar wind. The second type of induction is unipolar induction, in which currents are driven in the planetary conductor by the solar wind electric field associated with the moving magnetized highly electrically conducting solar wind. The former type of induction is very important on the moon whose surface, being nonconducting, prevents the latter type from occurring. But it may be important at Venus, and the former is a complication to be avoided in our present observational state.

We would not need to worry about unipolar induction if the ionosphere were either perfectly conducting or nonconducting. In the former case, the surface of the obstacle would be an equipotential and all flow would be diverted around the planet. In the latter, the flow would all crash into the planet as it does on the lunar surface. In the real world of Venus, the true case is somewhere between. Early investigators proposed a variety of possible scenarios for the solar wind interaction with Venus. Michel (1971) has classified these as the direct interaction model, the tangential discontinuity model, and the magnetic barrier model. These are illustrated in Fig. 7 with altitude profiles of the electron number density, magnetic field strength and horizontal flow velocity.

Fig. 7. Summary of Venus solar wind interaction models by Michel (1971). Profiles of electron concentration, n, magnetic field B, and horizontal flow velocity V are shown as a function of height for the three models. Incident solar wind velocity and field are normalized to unity. (a) Direct interaction model; (b) tangential discontinuity model; (c) magnetic barrier.

In the direct interaction model (Fig. 7a) the inflowing postshock solar wind depresses the ionosphere until a rather rapid transition occurs from flow to diffusive dispersion of the inflowing plasma. The solid line in Fig. 7a shows a case in which flow rate is entirely controlled by photoion pickup. Collisional stagnation above the peak of the ionospheric density would be expected to modify this profile as sketched by the dotted line. In the tangential discontinuity model illustrated in Fig. 7b, the ionospheric pressure equals the solar wind pressure across the ionopause and the plasmas are noninteracting. In the magnetic barrier model of Fig. 7c, a tangential discontinuity occurs between the flowing solar wind and the magnetic field. The magnetic field then transmits the pressure to the ionosphere, decreasing with decreasing altitude so there is net pressure balance in the ionosphere. None of these models agrees well enough with observations that we could call it a successful prediction. Nevertheless each of these models has some correct features.

The first clues to the correct nature of the planetary obstacle to the solar wind came from the Mariner Stanford Group's (1967) radio occultation data. The interesting feature of these data was the sudden drop in density with increasing altitude which has been called the ionopause. The name "anemopause" was suggested instead by Bridge et al. (1967) because it was supposed that the ionopause was where the solar wind stopped, but never became popular. This lack of popularity had nothing to do with the fact that the anemopause concept, as proposed, was incorrect.

Fig. 8. Dayside electron density profiles from Venera 9 and 10 occultation measurements at solar zenith angles of 14o (profile 1), 46o (profile 2) and 85o (profile 3) (Aleksandrov et al. 1976b).

Mariner 10 provided another dayside occultation with somewhat lower ionopause and not as sharp as Mariner 5. Venera 9 and 10 followed with many more. Figure 8 shows three altitude profiles of the number density obtained by the Venera orbiters (Aleksandrov et al. 1976b) at solar zenith angles of (1) 14^o , (2) 46^o , and (3) 85^o, illustrating the variability of the ionopause position.

These provided support for the existence of an ionopause, but until in situ data were obtained the physical nature of this boundary was uncertain. As discussed in more detail in Chapter 23 by Brace et al., the ionopause can generally be considered as a tangential discontinuity with ionospheric thermal pressure on the inside balancing magnetic pressure on the outside (Elphic et al. 1980b). There is little plasma pressure immediately outside the ionopause. This is apparent from the near perfect pressure balance between the magnetosheath field and the ionospheric plasma (Elphic et al. 1980b) and the near equality of these two pressures with the preshock solar wind dynamic pressure (Elphic et al. 1980a; Brace et al. 1980; Vaisberg et al. 1980).

Fig. 9. Magnetic field strength as a function of time through periapsis on three successive Pioneer Venus orbits at low solar zenith angle (Elphic et al. 1980a).

One might wonder what happens to the postshock solar wind flow. As Fig. 9 shows there is a gradual buildup in magnetic field strength as the planet is approached in the subsolar region, though the strength of this magnetic barrier varies from day to day. The data suggest that the postshock solar wind is mainly lost from the flux tubes of the magnetic barrier. Zwan and Wolf (1976) have suggested that this occurs because the hot plasma escapes along field lines as the tube slowly convects against the planet. Observationally, we obtain the picture shown in Fig. 10 as the post-Pioneer Venus magnetic barrier model.

Fig. 10 Post-Pioneer Venus magnetic barrier model. See legend to Fig. 7 for further explanation.

The first noticeable difference between Fig. 10 and Fig. 7c is that there is an abrupt increase in field strength with increasing altitude at the ionopause and a gradual decay, rather than the reverse. In conjunction with the abrupt increase in field strength the cold plasma density of the ionosphere abruptly decreases so as to maintain pressure balance with the magnetic field in the barrier. The density of the hot plasma decreases with decreasing altitude mainly due to the Zwan and Wolf effect. Outside the ionopause charge exchange occurs at a rate dependent on both the neutral atmospheric density and the postshock solar wind density. This further depletes the hot component and adds mass, if the charge exchange is between a light ion (e.g. H) and a heavy neutral (e.g. O). Photoionization also occurs here as it does throughout the atmosphere, at a rate dependent on the extreme ultraviolet (EUV) flux and the neutral density. This further adds mass to the postshock solar wind flow. The mass accretion slows down the flow; hence the velocity decreases as the planet is approached as does the electric field (not shown). In this model the anemopause (where the solar wind stops) is a gradual transition whereas the ionopause is a sharp transition, albeit the cold plasma density increase is somewhat spread out due to charge exchange and photoionization in the slow- moving magnetosheath just exterior to the ionopause.

The key assumption in this post-Pioneer magnetic barrier model is the equipotential nature of the ionopause, in contrast to the assumptions of the Cloutier model (cf. Cloutier et al. 1981 and Chapter 26). However, at the observed average altitude of the ionopause, and with the observed thickness of the ionopause current layer (Elphic et al. 1981) it is necessary to take into consideration the gradient terms in the generalized Ohm's law. When this is done the current is found to be overwhelmingly diamagnetic drift currents which are dissipationless, the Joule dissipation rate becomes extremely low and the potential drop across the entire planet is found to be 0.1 to 1V (Zeleny and Vaisberg 1981). There has been some evolution in the Cloutier model since its inception; originally there were no diamagnetic drift currents in the model (Cloutier and Daniell 1973; Daniell and Cloutier 1977). Later a diamagnetic drift current was added (Cloutier and Daniell 1979) but not selfconsistently. We believe that this non-self-consistency led them to the larger Ohmic current and potential drop across the planet. This discussion concerns only the high-altitude ionopause (altitude 300 km); when the ionopause current is pushed down to lower altitudes it does become significantly resistive and the Cloutier model becomes viable. We defer discussion of the low-altitude ionospheric obstacle for the moment.

The processes of charge exchange and photoionization were mentioned above without much discussion of what these processes do, and without much evidence that they take place. In a charge exchange an electron passes from a neutral atom to an ion. If the neutral atom is cold and the ion hot, the exchange produces a fast neutral and a slow ion. This can be an important process in the solar wind interaction with Venus, because it directly removes momentum from the postshock solar wind and deposits that momentum in the atmosphere. It also decreases the need for solar wind to be deflected by the planet and thus can weaken the shock and change its location, as we discuss in Secs. IV. A and IV. B. Charge exchange occurs most readily, i.e., has the largest cross section between members of the same species, such as H-H. However, proton-oxygen charge exchange is accidentally resonant and Venus has a significant hot neutral oxygen exosphere (Nagy et al. 1981). This charge exchange reaction also adds mass to the postshock solar wind.

Photoionization, by extreme ultraviolet emissions from the Sun, adds mass to the postshock solar wind flow. Photoionization of oxygen is very important because the exosphere is dominantly oxygen in the region of interest and because oxygen ions are much heavier than protons. These processes are important in the formation of the plasma mantle and magnetotail of Venus and will be discussed again in Sec. II.C.

C. The Low-Altitude Barrier

There is a noticeable change in the ionopause at low altitudes 300 km. This change can be seen in Fig. 11, which shows the thickness of the ionopause current sheet as a function of altitude.

Fig. 11. Thickness of the ionopause current sheet as a function of altitude. Thin lines show 2 and 10 gyroradii (Elphic et al. 1981).

The ionopause is typically ~ 30 km thick above 300 km but increases to 90 km thick at ~ 250 km (Elphic et al. 1981). We attribute this to the change in conductivity and Joule dissipation with altitude. When the ionopause is low, currents can be driven in the ionosphere as in the Cloutier model.

Fig. 12. Altitude profiles of the magnetic field strength and electron density measured by the electron temperature probe for the three passes shown in Fig. 10.

In fact, as shown in Fig. 12, the low-altitude ionosphere is magnetized during these periods of low ionopause altitude, which of course are associated with periods of high solar wind dynamic pressure (Luhmann et al. 1980). These high field regions are also more prevalent at low solar zenith angles; this fact is consistent with these high field regions lying along a belt and not covering the entire ionosphere, and is also in accord with Cloutier's model. Suffice it to say that at times when the solar wind dynamic pressure is high the nature of the solar wind interaction with Venus may be significantly different than under typical solarwind conditions.

D. Summary

We now know that for all practical purposes Venus is a nonmagnetic planet and the ionosphere is responsible for deflecting the solar wind flow. At times when the solar wind dynamic pressure is low and the ionopause altitude is above ~ 300 km, a magnetic barrier forms which deflects the solar wind before it directly encounters the ionosphere. At higher solar wind pressures, the ionopause moves to low altitudes, the current layer thickens, and a more direct interaction seems to occur in which currents are driven in the ionosphere by the solar wind electric field, i.e., by unipolar induction.

The region of postshock solar wind flow at Mercury, Earth, Jupiter, and Saturn is called the magnetosheath. As the physics of the postshock flow at Venus is basically the same as at the other planets, we will call this region the magnetosheath even though the term ionosheath has at times been popular. In this section we examine what is known about the flow properties of the magnetosheath, especially near the ionopause in the region known as the plasma mantle. Are the properties of this region those of a viscous boundary layer transferring momentum from the magnetosheath to the magnetosphere, or are they the properties of a mass accretion layer as sketched in Fig. 10?

A. Magnetosheath Flow

In the terminator regions of the magnetosheath, plasma measurements show that the plasma flow direction is satisfactorily described by the hydromagnetic model.

Fig. 13. Direction of flow (arrows on the upper part), ion temperature (in the middle) and velocity (below) as measured by Venera 10 RIEP on 29 October 1975. Location of the shock and the boundary of obstacle as well a flow direction (solid lines in the upper part) are from HD model of Spreiter et al. Dashed arrows show the low-energy ion flow, and closed dots (below) their parameters. Drafting errors in the originally published figure have been corrected. In particular, there is no component of solar wind flow into the wake anywhere on this pass.

Figure 13 shows Venera 10 RIEP analyzer measurements of the flow direction across the terminator compared with gas-dynamic calculations (Vaisberg et al. 1976). Note that the solar wind flow does not penetrate the wake region but the cold component does. Note also the heating of the cold component as the center of the wake is approached. Similar measurements have been obtained from a much larger data set on Pioneer Venus by Mihalov et al. (1980).

Fig. 14. Magnetosheath flow around Venus as measured by the Ames plasma analyzer in solar cylindrical coordinates - Vectors measured north of Venus are plotted in the upper half of the figure while the remaining vectors were measured south of Venus's orbital plane. Model ionopause and shock surfaces for Mach 8 and Mach 2 (dashed) are shown together with a Mach 8 streamline.

As shown in Fig. 14, over the terminators where the vast majority of the data have been obtained, the flow near the ionopause is tangential to it, as expected. There are only three reported measurements in the range of solar zenith angles (SZA) 40-60^o. Of these three, two have substantial components of the flow into the ionopause.

The postshock plasma parameters have been compared with theory by Mihalov et al. (1980), Mihalov (1981), and Smimov et al. (1981). While there is rough agreement between the observed and expected decrease in velocity across the shock, Mihalov et al. find that the observed decrease in velocity is too small in 80% of the cases for solar zenith angles <70^o. Furthermore, the observed magnetosheath ion temperature is close to a factor of two too low everywhere. These observations are supported by detailed case history models using the full gas dynamic modeling of Spreiter and Stahara (1980) in the study by Mihalov (1981). Deeper in the magnetosheath the observed ion temperature is found to be up to an order of magnitude too low. This observation is consistent with the significant decrease in ion temperature in the same region found earlier by Venera 10 (Vaisberg et al. 1976a), shown in Fig. 13. This could in part be due to the lack of information on the electron temperatures. However, in the terrestrial magnetosheath the electron temperature is seldom equal to the ion temperature, let alone greater than it. Furthermore, the ion densities are found to be less than predicted. Smirnov (personal communication, 1981) has reported that the observed ion temperature jumps were equal or somewhat lower than theoretical ones.

B. The Boundary Layer

A principal feature of the venusian magnetosheath is the boundary layer. This layer was characterized on different occasions by an unusually flat ion distribution function, by the depletion of the energetic tail of the ion spectra, and sometimes is seen as a rarefaction region (Vaisberg et al. 1976a; Romanov et al. 1978). Depletion of the plasma by the loss of energetic (1.5-3.0 keV) ions is clearly seen on successive spectra as measured on two revolutions of Venera 10 (Fig. 15). Magnetosheath spectra (12-5 for 10/31/75 and 10-6 for 11/6/75) show a gradual loss of the energetic tail of the ions while approaching the obstacle. This depletion is apparently due to the loss of energetic ions in the planetary atmosphere. The flow lines on which depletion occurs pass close enough to the atmosphere of Venus (comparable to twice the gyroradii of the energetic ions) for this depletion to take place.

Fig. 15. Successive ion energy spectra as measured by 4 electrostatic analyzers of RIEP on (a) 29 October 1975: and (b) 31 October 1975. Insert in the upper part shows the orientation of analyzers. Numbers on the orbit show the location of spectra. Scale on the bottom shows the energy intervals; time is given on the right.

On the night side of the planet, quite far downstream, the boundary layer is particularly apparent in the kinetic parameters of flow. Velocity and ion temperature profiles obtained during the crossing of the venusian interaction region from the bow shock to the central part of the wake, at distances 5-8 R from the planet, are shown in Fig. 16 (Romanov et al. 1978). The boundary layer is seen as a decrease of velocity and rise of temperature, as in a viscous boundary layer. The velocity profile, calcuated from the ion energy spectra under the assumption that the flow consists of protons everywhere, changes monotonically across the wake boundary that was crossed at 0255 Moscow time. The temperature drops at the boundary of the wake. As we will see later, this seems due to change of the ion composition across the boundary of the wake.

Fig. 16. Flow velocity and ion temperature along the orbit of Venera 10 on 18, 19 April 1976. Orbit of Venera 10 (solid line) is shown above orbit of Mariner 5 (dashed line). Boundary layer is clearly seen by velocity decrease and temperature increase. Temperature drops upon entry into magnetotail plasma. Circles on Mariner 5 trajectory correspond to events I through 5 in Fig. 3. Number 1, 2 and 3 on Venera 10 trajectory and time series indicate similar events on Venera 10 pass.

The viscous interpretation of the solar wind interaction with Venus has been championed by Perez-de-Tejada and colleagues (Perez-de-Tejada and Dryer 1976; Perez-de-Tejada et al. 1977; Perez-de-Tejada 1979; 1980 a, b).

This viscous interaction transfers momentum to the ionosphere causing a decrease in velocity and an increase in temperature as the planet is approached; Fig. 16 agrees with this view. In the region of velocity decrease the temperature rises. On the other hand, these viscous boundary layer treatments of the interaction sweep under the rug the underlying plasma instabilities and wave-particle interactions responsible for the observed plasma behavior. The analogy may be valid, but it does not show how the momentum transfer or the heating occur.

Two-component flow structure near the planet is often evident in the energy spectrum; in addition to the broad spectrum of shocked solar-wind plasma a narrow peak appears in the low-energy part of the spectrum. This peak can be seen in Fig. 15a (spectra 6-8) and Fig. 15b (spectra 5-6). This component is sometimes seen at larger distances from the planet. The number flux of these particles reaches ~ 10⁶ cm^-2 s^-1, with temperature on the order of several eV and kinetic energy ~ 100-200 eV. The most probable origin of these ions is photoion pickup. With the probability of photoionization of venusian exospheric ionsf ~ 2 X 10^-7 s^-1, neutral number density 10⁴cm^-3, and path length in which ion pickup occurs L ~ 10⁹ cm, we obtain number flux of photoions F = f n L ~ 2 X 10² cm^-2 s^-1. This compares favorably with the observed flux of low-energy component.

C. The Plasma Mantle

One of the most important results of the plasma measurements on Veneras 9 and 10 was the detection of plasma flow in the venusian wake with quite different parameters from the external flow (Vaisberg et al. 1976c; Romanov et al. 1978). As Shown by Romanov et al. the jump in flow parameters occurs at the boundary of the wake. No definite identification of the source of these ions was suggested until recently, though their planetary origin was evident. The same plasma layer (mantle) was identified from measurements of electrons on Pioneer Venus Orbiter by Spenner et al. (1980), who defined the mantle as the region in which the suprathermal electron energy distribution changes continuously from one characteristic of the magnetosheath to one characteristic of the ionosphere. The solar wind was inferred to stop above the mantle, i.e., the mantle boundary rather than the ionopause is the anemopause.

The magnetic barrier found on the day side of Venus (Russell et al. 1979a) helps explain the origin of the plasma mantle. The magnetic barrier develops as a result both of the current excluding the ionosphere from the magnetic field of external flow (for moderate solar wind pressure), and of magnetohydrodynamic (MHD) effects in the external flow. This barrier transfers the solar wind pressure to the ionosphere (Russell et al. 1979a).

Consideration of the pressure balance across the venusian ionopause shows that the magnetic pressure of the magnetic barrier is slightly smaller than both the solar wind pressure and the ionospheric pressure (Vaisberg et al. 1980) (see Fig. 17). This pressure deficiency could be provided by plasma pressure in the barrier (Vaisberg et al. 1980). An even smaller pressure deficiency was found in the analysis of data of Pioneer Venus orbiter (PVO) (Elphic et al. 1980a). In fact, the different estimations of the ratio of plasma pressure to magnetic barrier pressure are probably all consistent with unity within the calibration uncertainty of the plasma detectors.

Fig. 17. Comparison of solar wind pressure with (a) ionospheric pressure; with (b) magnetic pressure in the magnetic barrier ; as well as (c) between the magnetic pressure and ionospheric pressure, according to PVO measurements on first 18 revolutions. Solid lines show the best least-square fit. The dashed lines show the variation if the pressures were equal.

Observations of the magnetic barrier allow us to determine the topological connection of the venusian wake with the dayside structure. Since the boundary of the wake is the internal boundary of the shocked solar wind and since wake plasma is much hotter and dilute than ionospheric plasma, the magnetic barrier is evidently the dayside counterpart of the wake (Zeleny and Vaisberg 1981). The similarity of the electron components of the mantle on the day side and on the night side of the planet according to PVO data (Spenner et al. 1980) confirms this supposition.

There are three possible sources of plasma in the magnetic barrier: relative diffusion of the ionospheric plasma and the magnetic field; photoionization; and charge exchange in the barrier. To understand some of the properties of the plasma mantle, it is not important how these ions enter the magnetic barrier region, but only that they accrete in the barrier at a specific rate. Zeleny and Vaisberg (1981) have considered a model of the plasma mantle using the assumption that the source of the ions found there is photoionization, together with some simplifying approximations about the dynamics of the magnetic tube in the solar wind approaching ionosphere. Part of the hot plasma filling the flux tube that decelerates on the day side will move along the tube, leading to its depletion (Zwan and Wolf 1976). At the same time the new photoions will fill the tube, and while it sinks into the magnetic barrier the solar wind plasma is replaced by new cold photoplasma. Subsequent motion of the tube containing the photoplasma is determined by the magnetic pressure gradient within the magnetic barrier (the magnetic tension term is comparable to, but smaller than, the pressure term). Part of the new plasma will move along the tube, but most of the new ions will move with the magnetic tube towards the terminator, since the parallel velocity of the ions is small. Solving the equation of motion of the magnetic tube with a photoion source determined by the density of neutral oxygen in the exosphere of Venus gives the following parameters of mantle plasma near the terminator (Zeleny and Vaisberg 1981): velocity ~ 60 km s^-1, number density 4-5 cm^-3 number flux ~ 2.5 X 10⁷ cm^-2 s^-1. Figure 18 shows a sketch of the physics of the interaction.

Fig. 18. Model of solar wind-Venus interaction (Zeleny and Vaisberg 1981). Subsequent locations of magnetic field tube are shown. Spirals are trapped photoions; arrows show the flow of low-energy ions in the wake.

Observations of wake plasma at the distance ~ 0.5-1 R behind the terminator give a number flux smaller by an order of magnitude. This does not contradict the model, as the field tube should expand in the Wake by the factor of the ratio of wake radius to magnetic barrier thickness, i.e., by a factor ~10. Yet the energy of the directed motion of oxygen ions with velocity ~ 60 km s^-1 obtained in the model of Zeleny and Vaisberg is almost twice as large as the energy of motion of the tail ions, ~ 150 eV (Vaisberg et al. 1976b; Romanov et al. 1978). Taking the exospheric: helium photoions as the source of the mantle, we obtain an even larger discrepancy between the model and the observed energy of wake ions. This discrepancy suggests that anomalous ionization processes investigated recently by Formisano et al. (1982) may be important in the formation of the mantle. The anomalous ionization process will lead to the earlier loading of the magnetic field tube by photoions and diminish the final velocity of plasma. Using the anomalous ionization criteria by Formisano et al. (1982); Zeleny and Vaisberg (1981) were able to obtain agreement between the model of Venus's mantle and observations. On the other hand, recent measurements of the critical ionization velocity phenomenon in weak magnetic fields (Brenning 1981) show that the critical ionization velocity effect disappears when the electron gyro frequency is much less than the electron plasma frequency, as it is everywhere in the vicinity of Venus. Fortunately, there are other candidates for increased ionization. Charge exchange certainly takes place. Furthermore, the photoionization rate is greater than that which was assumed, since the authors used the neutral atmosphere of Niemann et al. (1979a) rather than the more recent Nagy et al. (1981) calculations of the oxygen exosphere. Finally, recent studies by Mihalov and Barnes (1982) of the Pioneer plasma data suggest that O⁺ is moving at speeds close to but less than that of the ambient magnetosheath plasma near to and sunward of the terminator plane. The most intense fluxes occur at low altitudes, as do the lowest speeds. At altitudes above a few thousand km, the two speeds are comparable.

D. The Venusian Wake

Energy spectra of the low-energy ion component found on Veneras 9 and 10 are shown in Fig. 15 (two lowest frames on each of Fig. 15a and 15b). The temperature of these ions typically lies between 1 and 10 eV, with energy of directed motion ~ 150 eV. Closer to the center of the wake both temperature and energy diminish. The velocity of low-energy motion has a strong component towards the center of the wake (Fig.13). Immediately behind the planet there is a cavity where no ions with energy 50 eV were observed on Veneras 9 and 10 (Vaisberg et al. 1976c). Similar results illustrated in Fig. 19, were obtained by Verigin et al. (1978), who described their results in terms of an umbra and penumbra. Low-energy ions usually occupy a significant part of the wake, yet regions free of E 50 eV ions were observed as far as 5-6 R downstream.

Fig. 19. Ion flux obtained 1 November 1975 by Venera 9 and 10 wide angle plasma detectors through the wake, into the magnetosheath, and through the shock (Gringauz et. al, 1976a).

The layer of plasma in the wake is thick at the high magnetic latitudes, where the vector to the spacecraft measured from the center of the wake is nearly orthogonal to the interplanetary magnetic field direction projected on the terminator plane. The layer of wake plasma is thin at low magnetic latitudes, where the spacecraft position vector relative to the center of the wake is almost parallel to the interplanetary magnetic field direction (Vaisberg 1980). Measurements of the layer of plasma in the wake at different magnetic latitudes are shown in Fig. 20.

Fig. 20. Dynamic spectra of ions for two passes of Venera 10: (a) 31 October 1975 and (b) 6 November 1975. Length of the bar is proportional to logarithm of the RIEP counting rate in specific energy interval (scale is on the left). (1) solar wind, (2) magnetosheath, (3) wake; (a) low magnetic latitudes, (b) high magnetic latitudes.

The cross section of the mantle at a distance ~ 1 R behind the terminator is shown in Fig. 21. These observations agree with the model of Zeleny and Vaisberg (1981) in that high-latitude plasma is convected within the flux tubes. Stronger deviation of plasma flow towards the center of the wake is due to the magnetic tension of bent magnetic-field lines. Low-latitude plasma is squeezed out along a relatively narrow layer of field lines (Zeleny and Vaisberg 1981).

Fig. 21. Projections of the parts of orbits in which low-energy wake plasma was observed. Every orbit was rotated until YZ-component of interplanetary magnetic field was coincident with horizontal (axis). Ring is Venus; dashed oval is suggested boundary of the cavity.

Comparison of the directions of flows near the boundary of the wake, where two plasma components (external and internal) mix, allows us to estimate the mass of the ions in the wake flow. From the measurements of directions of flow and energies of mass motion of these components the velocity of internal flow can be calculated, taking into account that the axisdirected convection velocities of the two components are equal. This estimate gives an ion mass in the internal flow between He⁺ and O⁺, closer to O⁺. In this case the typical velocity of O⁺ ions in the wake is ~ 40 km s ^-1 with typical number density ~ 0. 5 cm^-3.

Cases of quasi-periodic acceleration of ions were observed by Romanov et al. (1978); the periodicity of these events is 5 to 10 min. These events were interpreted by Romanov et al. (1978) as venusian substorms, the existence of which was suggested earlier by Russell (1976c) from the analysis of magnetic data.

The transition from external to internal flow at the wake boundary is quite distinct, although sharper and more gradual transitions were observed. The thickness of the transition layer ranges from 100 to 500 km. Strong fluctuations in the magnetic field with B B are usually recorded near the wake boundary. An example of a wake boundary crossing (Romanov et al. 1978) is shown in Fig. 22,in which a sharp transition from one flow regime to another one is seen, as well as strong magnetic field fluctuations. A burst of ions with energy ~ 2 keV is also observed at the boundary. Bursts of accelerated ions were observed at almost all crossings of the wake boundary. They usually lasted from 10 to 20 s in the spacecraft time frame. It appears that the acceleration of ions to an energy of several keV occurs permanently at the wake boundary and that it is associated with strong magnetic field fluctuations.

Fig. 22. Crossing of the wake boundary on 3 December 1975. Vertical dashed line marks the wake boundary. Magnetic field components are shown on the upper panels. The lower five panels show the measurements of the five electrostatic analyzers of RIEP: I3, (2.0-20 keV); I4, I6, (0.05-0.5 keV); and I7, (0.3-3. 0 keV).

Fig. 23. Spectrum synthesized from several spectra of the accelerated ions upon crossing the wake boundary.

E. Plasma Clouds

As seen in both the magnetic field and plasma measurements, the Venus ionopause is very dynamic, often varying dramatically in height from day to day (Elphic et al. 1980a, b; Brace et al. 1980). Often what appear to be multiple crossings of the ionopause suggest that the ionopause has a wavy structure. However, sometimes a pair of ionopause crossings is so far from the rest of the ionosphere that it seems to be a detached plasma cloud (Brace et al. 1982a). If these plasma clouds are moving at the solar wind velocity or any significant fraction thereof, they represent an important loss mechanism for the Venus ionosphere and ultimately for the Venus atmosphere. At minimum these clouds must be moving with a velocity greater than the escape velocity if they are detached from the ionosphere. Figure 24 shows an example of a plasma cloud occurrence during a periapsis passage of Pioneer Venus, and Brace et al.'s interpretation of this feature.

Fig. 24. Electron density profile through periapsis (SZA 77o) on Pioneer Venus orbit 138 containing a plasma cloud. Lower panel shows interpretation of these density enhancements (Brace et al. 1982a).

The magnetic field in a plasma cloud generally has the peculiar variation shown in Fig. 25 (Russell et al. 1982b).

Fig. 25. Magnetic field in solar ecliptic coordinates, plasma wave amplitude and electron density through a plasma cloud (orbit 601) (Russell et al. 1982b).

The magnetic field essentially reverses in a plasma cloud as if the field were pulling on the cloud. The location of this feature and the orientation of the field relative to the interplanetary magnetic field orientation is that which would be expected if the field lines were hung up near the subsolar point and bent into a hairpin shape. This feature is shown projected on the surface of the planet in Fig. 26. The plasma cloud and the bend in the field are along the noon-midnight magnetic meridian, the great circle through the subsolar point which is perpendicular to the interplanetary magnetic field. This observation suggests that the plasma cloud is, instead, a ridge of plasma on the ionosphere, presumably of flowing plasma. The plasma is accelerated and prevented from spreading by the magnetic field. Note the strong 100 Hz plasma waves occurring in the cloud. Increases in the plasma wave amplitudes are common at the ionopause but these have a different spectrum.

Fig. 26. Magnetic field on pass (orbit 601) shown in Fig. 25 projected on to surface of planet. Black bars on trajectory show plasma cloud and entry into ionosphere. Magnetic pole is intersection of noon-midnight magnetic meridian with terminator. Noon- midnight .nagnetic meridian is perpendicular of Y-Z ecliptic projection of interplanetary magnetic field, and passes through subsolar point.

The Maxwell stress on this plasma cloud is ~ 4 X 10^-8 dyne cm^-2 If we approximate the plasma cloud with a vertical slab 3^o thick assumed to be oxygen, it should be accelerated by this stress at 0.7 km s^-2. If the plasma started at rest at the subsolar point, it would reach velocity 90 km s^-1 at the point of observation with transit time 130 s if the acceleration were constant. If we assume that this accelerating cloud has height 500 km, the distance down to the ionopause as observed later on this pass, and that there is a similar cloud in the south, we would obtain a loss rate of 2 X 10²⁵ ions s^-1. We can also estimate the plasma temperature from the diamagnetic depression observed at the center of the cloud in Fig. 25. The sum of electron and ion temperatures is 7.2 eV. Since the electron temperature was observed to be 1.2 eV at this time, the ion temperature must be 6 eV. If the ions are O⁺, this temperature corresponds to a thermal speed of 8.5 km s^-1 close to the escape velocity. This low temperature compared to that expected for creation in the solar wind electric field indicates that the ions were created either in the ionosphere or in a low-velocity zone outside the ionopause.

The frequent occurrence of these clouds in the Pioneer data suggests that they are an important contributor to the loss of ions from the Venus ionosphere. The magnetic structure suggests that the dayside ionopause is not quasi hemispherical but rather ridged along a noon-midnight magnetic meridian. The two possible sources of these clouds or ridges are the ionosphere itself and the mass accretion region or mantle outside the ionopause. The first possibility leaves unexplained how the field penetrated the plasma; likewise, the second possible source does not explain how the field lines gathered so much plasma in their passage from the subsolar region. On the other hand, the cloud loss rate estimated here is not much different from the expected loss due to photoionization and charge exchange above the ionopause, and the field lines in the cloud must have swept across the entire dayside ionosphere before being observed. Nevertheless, the cloud formation is still somewhat of a mystery.

F. Summary

It is clear that in the immediate postshock flow, especially near the terminators, Venus has a magnetosheath much like the magnetosheath of any other planet. But at low solar zenith angles, where little data have been gathered, the situation may be different; the flow directions may differ from gas-dynamic calculations and the observed velocity decrease across the shock seems too small. The ion temperature also seems too low, sometimes by almost an order of magnitude. This region begs for further exploration.

Beyond the terminator the flow does not expand into the wake. The wake boundary is essentially impenetrable and the flow in the wake is different from the external flow. Close to the planetary ionopause the flow slows, the hot plasma density decreases, and the cold plasma density increases in the magnetic barrier or plasma mantle. This cooling of the flow, whether due to the loss of the high gyroradii ions from the flow because of their interaction with the atmosphere, or due to loss of the hot component due to charge exchange, presumably causes the rarefaction wave observed by Veneras 9 and 10 (Vaisberg et al. 1976b). Evidence for ion pickup is seen with the plasma analyzers in this region. As is evident in Fig. 13 this low-velocity component of the flow is very cold compared to the magnetosheath protons. Another ion phenomenon perhaps even more closely associated with the physics of the ionopause, is the formation of plasma clouds. We must admit, however, that the understanding of how these various features are related and what dominant physical processes control them is still rudimentary.

All planets with an intrinsic magnetic field have a magnetotail, which arises mainly because of tangential stresses applied by the solar wind. If there were only normal stresses on a planetary magnetosphere, it would assume a teardrop shape and would not have a long tail. Comets also have long tails; one type consists of dust, but the other consists of flowing ions presumably constrained by the magnetic field of a cometary magnetotail. A comet is not thought to have an intrinsic magnetic field; instead, the cometary magnetotail is thought to arise by capture of interplanetary magnetic flux by mass-loading of field lines which enter the cometary atmosphere (Alfven 1957). Processes such as charge exchange, photoionization, and possibly even the critical ionization velocity process (Alfven 1954) will add mass to the field lines in the region of closest approach to the comet. This mass accretion slows the motion of the magnetic flux tube. The plasma near the comet moves more slowly while that farther away moves at the solar wind velocity. The field line then bends around the comet and stretches out in the antisolar direction. Eventually these draped field lines will slip through or around the comet. In analogy with Dungey's (1965) estimate of the length of the terrestrial magnetotail, we can estimate the length of a cometary tail as the distance the solar wind travels as the field line drifts through or around the comet. If this slippage past the comet takes about half a day, the resulting tail would be ~ 0. 1 AU long. However, the situation is not perfectly analogous. Once the field line has left the massloading region it will attempt to straighten itself, but can only do so slowly because of the added mass at the crook of the field line.

We would expect this same cometary process to occur at Venus, though we cannot see any visible evidence for a venusian tail. The expected atmospheric loss rates at Venus, on the order of 10²⁵ ions s^-1, are many orders of magnitude below the loss rates for a large comet. However, our discussion in earlier sections demonstrates that processes analogous to those occurring in comets occur at Venus. We expect the field lines to be slowed as they move across the front of the planet. The plasma clouds discussed in Sec. II.E seem to be a manifestation of this process. The field lines then slip over the poles, or more precisely over the magnetic poles, of the planet, all the while trying to straighten up and thereby accelerating the plasma contained in the crook of the field line. This is what we expect, but do the observations support this picture?

A. First Observations

As shown in Figs. 3 and 4 Mariner 5 made a close approach to the axis of the Venus wake although it did not pass directly behind the planet. Close to point 3 of Fig. 3 the magnetic field became steady and was roughly aligned with the solar wind flow. This behavior suggests that Mariner 5 entered a Venus magenetotail though it could not be distinguished whether this tail was due to intrinsic field sources or to 'hung up' interplanetary field lines. Figure 4 illustrates the alignment of the field more clearly.

Fig. 27. Ion fluxes as a function of energy and time recorded by the narrow angle detectors in the distant tail on 17 April 1976 (Romanov et al. 1978). The top panel shows 20 s averages of the magnetic field; the bottom panels show the ion fluxes. The instrument does not sample all energies simultaneously but rather stays at each energy for 20 s and then steps to the next lowest energy; after 160 s it repeats the cycle.

The Venera 9 and 10 orbiters probed the wake region often and at various distances. The properties of the tail region as seen in the magnetorneter and plasma data much resembled those of Earth's magnetotail. Figure 27 shows ion fluxes recorded by the narrow angle plasma detectors in the distant tail (Romanov et al. 1978). The top panel shows magnetic field strength; the bottom panels, ion flux versus time and energy. The instrument does not sample all energies simultaneously, but samples one energy for 20 s, and then steps to the next lowest energy for 20 s, etc. Thus, time runs both upwards and to the left in this diagram. The feature of interest here is the periodic increases and decreases in the magnetic field strength accompanied by energizations of the plasma and sometimes followed by dropouts of the flux. Romanov et al. (1978) interpret these as substorms.

Fig. 28. Distribution of magnetic field vectors near a laboratory model of Venus (wax sphere 6 cm in diameter) (Dubinin et al. 1978).

At this point we should give credit to the guidance of laboratory simulations in our interpretation of the Venus data. While the solar wind interaction with Venus cannot be duplicated in the laboratory, qualitative scaling can be performed in which ratios of scale lengths have the right ordering if not the proper size. Venus has been simulated with a wax sphere having a conducting core in a magnetized supersonically flowing plasma (Dubinin et al. 1978). The results of this simulation are shown in Fig. 28. The interplanetary magnetic field, as we expect from the cometary analogy, becomes draped around the obstacle. It is not clear, of course, whether the hanging up of the field lines is due to unipolar induction in the planetary ionosphere or mass-loading ahead of the obstacle. Dubinin et al. (1978) point out that these results show that the mere existence of a tail is no proof of an intrinsic magnetic field. The dependence of tail field polarity on the interplanetary magnetic orientation must be investigated.

The Venera 9 and 10 wake data, however, appear so much like similar observations in the Earth's magnetotail that Dolginov et al. (1978) interpreted them in terms of an intrinsic magnetic field. The surprising aspect of this interpretation is that the field lines do not radiate from the polar regions of the planet, but are confined to a conical region expanding away from the core. As we discuss below, the tail observed by PV flares slightly but shows no evidence of converging on to the planet this rapidly. On the other hand, there are small regions of radial field lines on the night side of Venus associated with holes in the nightside ionospheric density. Perhaps observations of the magnetic field in these nightside holes led to Dolginov et al.'s interpretation.

The interpretation is, of course, incorrect; Venus does not have an intrinsic field large enough to supply the magnetic flux in such a tail. Yeroshenko (1979) correctly pointed out that the Venus tail polarity was controlled by the interplanetary magnetic field. Currently, none of the various principals involved are holding out for an intrinsic magnetotail at Venus.

B. Pioneer Observations

Fig. 29. Pioneer magnetic field measurements in the Venus wake region (orbit 189, 6/11/79). Coordinate system is spacecraft coordinates, which is very close to solar ecliptic.

Because of the high inclination of the Pioneer Venus orbit, the spacecraft cuts through the wake in two regions of limited radial extent: one near the planet at 2000 km altitude; and one far down the wake from ~ 7 to 12 R. Figure 29 shows the magnetic field on a pass through the distant wake. In the center of this period there is an interval of enhanced though fluctuating field strength, during which the field is mainly oriented along the solar wind direction (the X-direction). This interval has been defined as the magnetotail (Russell et al. 1981). The frequent reversals in the X-component signal tail current sheet crossings. Their coincidence with the low-field regions imbedded in the high-field region suggests the current sheet has an associated plasma sheet, in analogy to the Earth's plasma sheet and tail current sheet. Since the interplanetary magnetic field is usually close to the ecliptic plane, the tail current sheet should be mainly in the north-south plane, rather than the ecliptic plane as the Earth's current sheet is. The high inclination of the Pioneer orbit causes the spacecraft to fly roughly parallel to this sheet. Perhaps flapping of the tail causes the multiple current sheets observed. However, it is also possible that the Venus tail is actually striated.

Fig. 30. Pioneer magnetic field measurements in the Venus wake together with the plasma wave amplitudes at 5.4 kHz (orbit 188, 6/10/79) (Russell et al. 1981).

Figure 30 shows another passage through the tail as seen in the Pioneer magnetic field data. In this figure we have added the amplitude of plasma waves seen at 5.4 kHz. The striking observation here is that the region immediately outside that identified as magnetotail has strong plasma wave emissions. Some special process occurs here in the plasma, but at this time we do not know what it is.

Fig. 31. Terminator plane projection of orbit during Pioneer tail crossings (heavy lines). Orbit numbers are given at the start of each pass (bottom of panel). Numbers along trajectory give the x distance at tail entry and exit (Russell et a]. 1981).

Figure 31 shows the location of these tail crossing intervals as projected in the plane perpendicular to the solar wind. The region is much larger than the optical shadow of Venus, ~ 4 R across, and it is shifted from the optical shadow because of the 35 km s^-1 planetary motion of Venus perpendicular to the solar wind. The circle we have drawn does not order the data perfectly, probably because of the variability of solar wind direction relative to the solar direction.

Fig. 32. Location of distant bow shocks and tail encounters in solar cylindrical coordinates (Russell et al. 1981). Boxes show near tail encounters by Veneras 9 and 10 (Rornanov et al. 1978); triangles show Venera 4 and Mariner 5 near tail encounters.

Figure 32 shows these locations as a function of distance along the Venus-Sun line, together with the distant bow shock locations. Also shown are corresponding observations with Mariner 5 and Veneras 4, 9, and 10. The tail is seen to expand slightly behind the planet and is consistent with the source of field lines being in the dayside magnetic barrier. However, some of the tail field lines certainly penetrate the night ionosphere. The nighttime ionosphere field is generally weak, but there are two magnetic features which may be related to the magnetotail (Luhmann et al. 1981a,b). First are radial magnetic fields associated with large plasma density dropouts or ionospheric holes (Brace et al. 1982b; see also Chapter 23). These regions are limited in extent but are very interesting from both plasma physical and also atmospheric viewpoints, since they allow whistler mode electromagnetic waves to propagate through the ionosphere, thus permitting the mapping of lightning source locations (W. W. L. Taylor et al. 1979b; Scarf et al. 1980a). Ile second feature is the occurrence of large horizontal low-altitude magnetic fields when the solar wind dynamic pressure is high. This field configuration also leads to a disappearing ionosphere in which the electron density is suppressed far below its usual value throughout the night ionosphere. Thus, the tail may be somewhat different at times of high solar wind dynamic pressure.

If our estimate of the transit time of a field line through a plasma cloud is appropriate for typical conditions at Venus, i.e. ~ 2 min, then our introductory tail formation scenario would give a Venus tail length 8 R since the solar wind moves 4 R min^-1. A similar time estimate (~ 140 s) is obtained by Zeleny and Vaisberg (1981). However this estimate is only an approximation, since the field lines which do slip by Venus do not immediately straighten up. Thus we are not surprised to see a well-developed magnetotail at 12 R. Another implication of this slippage is that all the field lines in this tail do not close across the day side of Venus, nor in the night ionosphere. Instead, many field lines close across the tail current sheet. We note that the magnetic flux in the tail is on the order of 3 megaweber. Even if we assume that the field in the magnetic barrier is 100 nT and 500 km thick, which are large estimates, there is only 1 megaweber across the front of the planet. Holes have < 0.5 megaweber. Again, this calculation suggests that there must be much closure across the current sheet.

Fig. 33. Magnetic field in the Venus tail at current sheet crossing. Left, time series in spacecraft coordinates. Right, hodogram of tip of field vector between intervals marked by vertical lines on right, expressed in minimum variance coordinates (orbit 182, 6/4/79) (Russell et al. 1982c).

Examination of the magnetometer data at high resolution across a current sheet in the center of the tail is shown in Fig. 33. The crossing seems to be that of a MHD tangential discontinuity with little normal component at this time. The direction of the normal to this current sheet was in the solar ecliptic Y-direction, as expected (Russell et al. 1982c).

Fig. 34. Magnetic field in Venus tail at magnetopause crossing (orbit 182, 6/4/79) (Russell et al. 1982c).

The comparison of the Pioneer plasma and magnetic field data in this region is still at an early stage. Heavy ions are supposedly seen in this region (Mihalov et al. 1980) but sometimes there are no heavy ions, and sometimes there is no detectable plasma at all. Much work remains to be done.

C. Summary

There is agreement now that Venus has a magnetic and plasma tail. However, much of that field does not originate in the ionosphere but comes from the magnetic barrier and some of it, the older field lines, close across the center of the tail itself. It is perhaps incorrect to call such a tail induced, because it does not arise from either of the processes usually called induction. It is mainly an accreted tail. The tail in many respects seems similar to the Earth's tail, but this may be deceptive. It will be interesting to discover the exact interrelationship of the proton and oxygen flow regions to the magnetic structure. Furthermore, much work remains to be done on the microphysics of the tail. What is the cause of the plasma wave emissions seen bounding the tail? What is the usual hydromagnetic character of the tail current sheet and magnetopause? Is the formation of holes in the night ionosphere related to tail formation?

Fig. 35. Gas dynamic computer simulation of interaction of solar wind with an impenetrable object (Spreiter et al. 1966). Streamlines and wave patterns for supersonic flow past the magnetosphere are shown (M = 8, = 2).

It is clear from the observations reported above that Venus presents a fairly hard obstacle to the solar wind flow. Only a small fraction of the solar wind appears to be absorbed. Since the solar wind flows much faster than the speed of pressure waves, i.e. magnetosonic waves, a shock wave must develop in the flow. This shock wave both heats and deflects the flow. One might not expect the strength of this shock or its dependence on solar wind conditions to differ markedly from the terrestrial bow shock, since the plasma conditions in the solar wind at Venus are similar to those at Earth. However, it is of interest to see how the Venus shock differs from the terrestrial shock. Disparities that could affect the shock are: the scale size of the Venus shock, one tenth of the size of the Earth's; and momentum removal and mass addition by such non-MHD processes as charge exchange and photoionization. There are a variety of tests we may employ. We can examine the position of the shock, the jump in field strength and number density at the shock, the dependence of the structure of the bow shock on plasma conditions, and the wave phenomena seen upstream from the bow shock.

Fig. 36. Gas dynamic computer simulation of flow past an impenetrable obstacle. Velocity and temperature contours for supersonic flow past the magnetosphere are shown. (M = 8, = 2) (Spreiter et al. 1966).

As a point of reference for our observations we use the numerical computations of Spreiter et al. (1966) who studied the interaction of a supersonic gas-dynamic flow with a nonabsorbing nonviscous boundary. Figures 35, 36, and 37 show the streamlines and wave patterns, the velocity and temperature contours, and the mass flux contours for supersonic flow past the Earth's magnetosphere for Mach number 8 and ratio of specific heats 2. We use a specific heat ratio of 2 because such a ratio has been repeatedly observed to give the best fit in studies of the bow shock position (Fairfield 1971; Zhuang and Russell 1981) even though = 5/3 is most frequently used in simulations. We also note that a slightly lower Mach number would be more appropriate; it would result in a bow shock slightly farther from the planet.

Fig. 37. Gas dynamic computer simulation of flow past an impenetrable obstacle. Mass flux contours for supersonic flow past the magnetosphere are shown (M = 8, = 2) (Spreiter et al. 1966).

Spreiter et al (1970) have presented similar diagrams specifically for a nonmagnetic planet. However, all the figures are given for a ratio of specific heats of 5/3 rather than 2. Their paper presents a correspondence rule that relates diagrams derived for the Earth's magnetospheric interaction to ionospheric interactions. In particular, if the scale height of the Venus ionosphere is 1500 km, then the curves are identical with the subsolar ionopause coincident with the subsolar magnetopause. If the scale height is 750 km, the same curves can be used by shifting the center of the nonmagnetic planet away from the Sun by 840 km. If the scale height is 75 km, the same curves can be used by shifting the center of the nonmagnetic planet by 2430 km away from the Sun. We assume that the correspondence rule for a ratio of specific heats of 2 would be quite similar.

A. Position of the Bow Shock

Fig. 38. Location of Venus bow shock up to initial Venera 9 reports (Russell 1977).

The Venus bow shock was not crossed by Mariner 2 but was observed on all succeeding missions. The bow shock was roughly where it was expected to be for an obstacle the size of the planet. However, it was somewhat closer on the average than would be expected by scaling from our knowledge of the terrestrial shock and model calculations (Russell 1977). In fact, if the average Venus shock surface were scaled down so that it passed through the Mariner 10 shock crossing (Ness et al. 1974), the resultant shock surface would intersect the planetary surface. This observation prompted Russell (1977) to question whether, on occasion, the Venus shock could become attached to Venus as shown in Fig. 38. Even if the anomalous Mariner 10 crossing could be explained, there is still the implication that on the average 30% of the solar wind streamlines incident on Venus passed into the ionosphere, if we naively use the gas dynamic computations of Spreiter et al. (1966) to map the flow. This suggestion was criticized by Verigin et al. (1978) who claimed that there was nothing unusual about the Venus bow shock position as measured by Veneras 9 and 10. However, later analyses of an extended Venera 9 and 10 data base showed unequivocally that the subsolar bow shock, as well as the bow shock at the terminators, was well below the expected location (Smirnov et al. 1980). The subsolar shock is expected to occur at an altitude of ~ 3000 km, but the subsolar altitude of the bow shock is close to half this value. It lies only slightly > 1300 km above the surface of Venus, in surprising agreement with the initial estimate of 1450 km altitude, using only 7 data points and an approximate aberration correction (Russell 1977).

A further surprise in the study of the shock location came with the availability of Pioneer Venus observations. The bow shock near the terminator (where the Pioneer Venus observations occurred because of the orbital inclination) was farther from the planet than it was during the Venera 9 and 10 mission (Slavin et al. 1979a). The extrapolated subsolar point was also higher, ~ 2270 km rather than ~ 1500 km, but still much closer to the planet than would be expected for complete deflection of the solar wind. However, when aberration corrections are made this difference is much less than originally reported (Smirnov et al. 1981).

Slavin et al. attributed the difference between the Venera 9 and 10 positions and the Pioneer orbiter positions to solar cycle effects. Veneras 9 and 10 obtained data at solar minimum and Pioneer at solar maximum. However, latitudinal asymmetries controlled by the interplanetary magnetic field (IMF) direction had been predicted by Cloutier (1976) and reported in the Venera data by Romanov (1978). A search for these asymmetries revealed no such effect in the Pioneer data (Slavin et al. 1979b). In order to reconcile the differences between the Venera and Pioneer data, Smirnov et al. (1981) proposed a slightly different mechanism: the additional aberration of the shocked solar wind caused by magnetic stresses in the magnetosheath (Walters 1964) tilted the shock nose in the direction of planetary motion, moving the shock in, on the average, above the evening (-Y solar ecliptic) terminator and out, on the average, above the morning terminator. The Venera 9 and 10 shock crossings were recorded mainly on the west side of the planet and were strongly affected by any shock tilt, whereas the Pioneer data at high latitudes would be little affected by this process. The available data support this interpretation, as shown in Fig. 39. Note that the terminator projection of the shocks as determined in this analysis has only a 5% asymmetry, presumably due to either solar cycle effects or latitudinal differences as proposed by Romanov. We may be able to test the possibility that the difference is related to the solar cycle with the Pioneer data as the solar cycle declines.

Fig. 39. Location of Venus bow shock in terminator plane as observed by Pioneer Venus and Veneras 9 and 10 (Smirnov et al. 1981).

Thus, the study of the location of the Venus bow shock has led to two surprises: (1) the bow shock is closer to the planet than expected for complete deflection of the flow; and (2), while the shock may show additional aberration by the Walters effect, it shows little asymmetry about this aberrated flow direction. We examine the implications of these two observations, taking the second one first. There are two reasons we might expect an asymmetry about the solar wind flow which is ordered by the IMF (cf. Cloutier 1976). First, the effective obstacle may be asymmetric. Ions created in the postshock solar wind flow will have large gyroradii of several thousand km, because they feel the solar wind electric field. Depending on the orientation of the IMF these ions will either initially spiral away from the planet and escape, or spiral into the ionosphere and be lost. If this mass-addition process is an important contributor to the deflection of the solar wind, then there should be a north-south asymmetry. However, the magnetic barrier is a region of weak solar wind and slow flow. The electric field is not expected to be large near the ionopause, where the rate of production of these ions would be greatest. The second proposed source of asymmetry is the asymmetric propagation of magnetosonic waves in a magnetized plasma. This effect may be detectable far downstream from Venus, but the shock shape over the front of the shock is principally determined by the shape of the obstacle. This fact can be appreciated by noting the relative insensitivity of the shock stand-off distance to changes in Mach number for usual solar-wind Mach numbers (Spreiter et al. 1966), and by referring to Fig. 35 in which the Mach lines indicate what portion of the shock is sensitive to what portions of the ionopause. In short, the shock is quite symmetric because it is only weakly sensitive to obstacle shape and furthermore, because the obstacle is quite symmetric.

The unexpected closeness of the Venus bow shock is less easy to explain. The Venus ionopause generally appears to behave as an equipotential surface. If only 10% of the solar wind flow directly interacted with the ionosphere, a potential drop of ~ 5 kV would be applied across the planet (400 km s^-1 times 10 nT times 1200 km). This is to be compared with the 40 V of the model of Daniell and Cloutier (1977) and the ~ 1 V of Zeleny and Vaisberg (1981). The solution of this paradox must lie in the nature of the interaction. If momentum is removed from the flow via charge exchange, then the direct interaction with the neutral atmosphere occurs via fast charge-neutral particles. The charged particles can flow around the planet in the usual way, in which the single streamline intersecting the stagnation point coats the entire ionopause. The direct loss of momentum to the atmosphere obviates the need to bend the flow as sharply around the planet, and the shock can turn more normal to the incident flow throughout the subsolar region. This blunter shock will be closer to the subsolar ionopause than the normal shock. The loss of momentum to the planet will also alter the mass flux contours seen in Fig. 37. Charge exchange does not alter the number density (although it could alter mass density) but it will drastically reduce the velocity of the plasma and thus reduce the mass flux crossing the terminators. This will lead to further reductions in the height of the shock in the terminator region, even though the shape might not be affected here. With recent advances in numerical simulations these ideas could be readily tested, but this has not yet been done; until it.has, we must treat these ideas only as a plausible scenario. On the other hand, observations of shock strength and calculations of processes that could lead to weakening the shock do support these ideas. Furthermore, studies of mass addition at the subflow point, a process which differs in sign from the effect of charge exchange, show the shock to move away from the obstacle (Cresci and Libby 1962).

Finally, Slavin et al. (1980) have studied the dependence of the terminator crossing of the bow shock on the ionopause altitude, the Alfvenic Mach number, and the solar wind pressure. They find that the position of the shock near the terminators depends little on the ionopause altitude, and weakly on solar wind pressure. The correlation with Alfvenic Mach number is much higher (0.7), corresponding to a more distant shock at low Mach numbers. This correlation is presumably due to two factors: a change of stand-off distance, especially at the lowest Mach numbers; and a change in the flaring angle of the shock.

B. The Strength of the Bow Shock

Strictly speaking, the strength of the bow shock is measured in terms of its Mach number. However, we operationally determine how strong a shock is by measuring the jumps in magnetic field density and temperature or the decrease in velocity across the bow shock. At low Mach numbers 3 the jumps in magnetic field and density are very sensitive to the Mach number but they asymptotically approach upper limits as the Mach number increases above 3 (Tidman and Krall 1971). Mach number 3 is also the approximate dividing line between laminar and supercritical shocks. Supercritical shocks have abundant downstream ion heating, are more turbulent, and are presumably associated with upstream ion beams although this latter conjecture has not yet been adequately tested with terrestrial data. We would naively expect that since the bow shock stands in the flow, and therefore its velocity is equal to that of the solar wind but oppositely directed, the Mach number would be specified by knowledge of the upstream solar wind conditions. However, mass addition and charge exchange complicate the picture, so we cannot make this simple assumption, as we will demonstrate from observations.

The first investigator to point out something strange about the Venus bow shock was Wallis (1972), who proposed that because of mass accretion due to photoionization there might be no shock directly in front of Venus. He felt that the failure to detect upstream energetic particles at Venus supported this view. However, as we will discuss below, the structure of the Mariner 5 shock has another explanation and there certainly are upstream particle beams at Venus even though the evidence is indirect. Wallis (1972) proposed that this neutral exosphere from which the solar wind accreted mass through photoionization was constituted of helium, whereas we now know there is a hot oxygen exosphere. In a later paper Wallis (1973) admits to the existence of a Venus bow shock but proposes that it is weakened by the mass accretion process. Wallis also discusses the effects of symmetric charge exchange (proton-hydrogen) which has zero mass accretion, and points out that this process reduces the required expansion and so weakens the shock. In fact, the immediate downstream cooling rate can be so severe that the flow does not need to expand laterally. Although there is much to criticize in these two papers based on our present observational knowledge, the physics is sound and was very innovative for its time.

Fig. 40. Ratio of downstream to upstream magnetic field strength as a function of solar zenith angle at Venus, compared to computer simulation (Russell et al. 1979b).

Wallis's evidence that the Venus shock was weak was the supposed lack of upstream particles and the behavior of the Mariner 5 shock. Other explanations exist for these observations, but it is difficult to dismiss the fact that the magnetic shock jump at Venus is less than expected from theory, as shown in Fig. 40 (Russell et al. 1979b). The theoretical number would be lowered somewhat by using a ratio of specific heats of 2, but not as far as the observed ratio. Nine terrestrial shocks were examined at median solar zenith angle 72^o and found to have average shock jump of 3.1, i.e., 25% greater than their Venus counterparts. Further evidence for a weak bow shock lies in the attempts to model the interaction with a gas-dynamic code. Spreiter and Stahara (1980) have attempted to duplicate the parameters along two Pioneer Venus periapsis passes with their computer code. Using a specific heat ratio of 2 they obtain good qualitative agreement with the field and plasma behavior, but the observed field jump is less than the predicted jump in every case. In one case for which ion temperatures were available downstream, the temperature was ~1/2 the expected value. We note that when a ratio of specific heats of 5/3 was used, an unrealistically low Mach number was required to get the predicted shock position out to where it was observed, and the predicted shock jump became too low. Mihalov (1981) followed this work with a study of six additional examples; he measured the solar wind conditions and hence knew what the Mach number should be. However, he allowed the Mach number to vary to get the best fit in the shock jump and shock position. This best fit Mach number was close to 1/2 the observed solar wind Mach number. Even choosing a Mach number much lower than expected for the calculations, Mihalov was unable to approximate the postshock ion temperatures. In every case examined the ions were factors of 2 to 10 colder than calculated. The missing heat could be in the electrons, but if so the electrons would have to be much hotter than the ions and much hotter than electrons observed in the terrestrial magnetosheath.

Fig. 41. Absorption of solar wind as a function of distance of streamline from planet at terminator. Top curve includes hot oxygen exosphere in calculation (Gombosi et al. 1981).

Gombosi et al. (1980) proposed and tested a new idea of removing solar wind from the postshock magnetosheath flow, that is, scattering of ions into the ionosphere by magnetic fluctuations. However, the dayside magnetic barrier has high field strengths and is generally quite steady. Computed loss rates using measured fluctuation amplitudes indicated that only 0.3% of the solar wind would be absorbed on the average, an entirely negligible amount. Since the same code could be used to calculate the charge exchange loss rate, Gornbosi et al. (1980) also treated this problem and found that the loss rate was much greater, from 1% to 7% for ionopause altitudes from 1000 to 200 km respectively. The loss rate depends greatly on the atmospheric model used, and when the existence of a hot oxygen exosphere was realized by Nagy et al. (1981) the calculations were immediately redone because protons have an accidental charge exchange resonance with oxygen atoms. As shown in Fig. 41 the inclusion of hot oxygen doubled the loss rates (Gombosi et al. 1981).

In summary, the observed inward displacement of the bow shock from its expected location, the observed weakness of the bow shock and the theoretical expectation of solar wind loss by charge exchange consistently show non-MHD processes to play a significant role in the postshock flow. However, we emphasize that to date these processes have not been incorporated in a selfconsistent model which calculates shock position and strength in a charge exchanging flow.

C. Shock Structure

Fig. 42. Magnetic field data during Mariner 5 encounter, plotted in solar cylindrical coordinates.

The structure of the Earth's bow shock depends on the Mach number of the solar wind flow, the ratio of solar wind flow velocity to the magnetosonic velocity, the of the solar wind (the ratio of thermal to magnetic energy density) and the local angle between the upstream magnetic field and the shock normal (cf. Greenstadt and Fredricks 1979). The dependence of the structure of the Venus bow shock on the IMF direction was first pointed out by Greenstadt (1970) in the Mariner 5 data, the inbound shock being quasi perpendicular and the outbound shock quasi parallel. This outbound shock crossing is shown in the magnetic field as displayed in solar cylindrical coordinates in Fig. 42. The interplanetary magnetic field in the region labeled UW (Upstream waves) is Closely aligned with the shock normal. When the field direction changes abruptly at the end of the UW interval to a direction not so nearly aligned with the shock normal, the upstream waves cease. This behavior is very similar to that observed at the terrestrial shock. The Mariner 10 outbound shock crossing shown in Fig. 43 was also a quasi parallel shock and was similarly disturbed (Ness et al. 1974). Due to the very turbulent nature of the quasi parallel shock it is difficult to characterize these shocks or to look for subtle differences between them and their terrestrial counterparts; in the following discussion we will therefore examine mainly quasi perpendicular shocks in which the interplanetary field lies mainly in the shock plane.

Fig. 43. Magnetic field data (6 s averages) during Mariner 10 encounter 5 February 1974 plotted in solar ecliptic coordinates (Ness et aL 1974) ----- S/C magnetic field;mn B larger amplitude, lower frequency: B smaller amplitude, higher frequency.

One of the features of supercritical quasi perpendicular shocks evident in terrestrial data is an overshoot in field magnitude immediately behind the shock. After a few tens of seconds the magnitude returns to an equilibrium value, often by passing through an undershoot (Russell and Greenstadt 1979). The plasma also participates in this phenomenon with electrons following the field closely (Bame et al. 1979). The overshoot appears to be associated with ion gyration at the shock front; the ions do not become thermalized until downstream of the overshoot.

Fig. 44. Magnetic field magnitude across bow shock at Venus for varying solar wind conditions (Russell et al. 1982a).

Figure 44 shows a set of Venus shock crossings as measured in the magnetic field strength (Russell et al. 1982a). Well-developed overshoots such as those seen in the second and third panels from the bottom are somewhat rare in the Venus data. The overshoots also seem somewhat weaker at Venus than would be expected under the same solar wind conditions at other planets. Figure 45 shows the overshoot magnitude as a function of magnetosonic Mach number for Venus, Earth, Jupiter, and Saturn, where the Mach number is derived from measured solar wind conditions. The Venus points are lower than those from the other planets. If, as we have speculated above, charge exchange in the postshock flow indeed weakens the Venus shock so that its effective Mach number is less than its nominal Mach number, then the Venus points on Fig. 45 would move to the left and move into agreement with the measurements at other planets.

Fig. 45. Magnitude of the overshoot at the bow shocks of Venus, Earth, Jupiter, and Saturn as a function of magnetosonic Mach number (Russell et al. 1982b).

Note that there are significant differences in the spectrum of plasma waves as seen at Venus, Earth, Jupiter, and Saturn, shown in Fig. 46 (Scarf et al. 1981).

Fig. 46. Very low frequency electric plasma wave spectrum at the bow shocks of Saturn, Jupiter, Earth, and Venus (Scarf et al. 1981).

The waves at the Venus shock are stronger at low frequencies and have a monotonic decrease with increasing frequency. Scarf et al. (1980b, 1981) attribute this to the difference in Mach number with heliocentric radius. This may be true, but another apparent difference is the ratio of ion to electron temperature, which seems much lower for the Venus magnetosheath than the other planets. From these differences Scarf et al. (1980b) conclude that the plasma instability conditions at Venus are not the same as those in the Earth's bow shock region.

D. Upstream Waves and particles

Fig. 47. Successive energy spectra (right hand columns in each pair) measured by the electrostatic analyzers of the RIEP plasma spectrometer on Venera 10 during a near planetary pass on 31 October 1975, through: (1) the wake, (2) rarefaction region, (3) magnetosheath, (4) foreshock, and (5) undisturbed solar wind. Left hand columns are the mean square fluctuations of the counting rate at each energy step during the 20 s measurement interval. Scale on the bottom is energy; scale on the right is time.

None of the Venus missions have included particle measurements with either sufficient sensitivity or the proper look directions to see particles emitted along the field lines upstream from the bow shock. However, the absence of these observations should not be taken as absence of the phenomenon. The waves associated with these ion beams as Earth (cf. Hoppe et al. 1981) are also seen at Venus, both in the magnetic field (as in Fig. 42 in the region labeled UW), and as fluctuations in the solar wind ion fluxes. The reason for this latter effect is that the waves driven by the upstream ion beams are hydromagnetic in nature. When the magnetic field twists, so does the flow velocity; when the magnetic field is compressed, so is the solar wind density. This effect is clearly seen in Fig. 47, in which the left-hand column of each pair shows the fluctuation in the ion current. On the right of each column pair is the 20 s average ion current as a function of energy. These data were obtained with the Venera 10 electrostatic analyzers on an outboard pass on 31 October 1975 starting in the planetary wake and moving out into the solar wind. Region 1 is the wake, 2 the rarefaction region, 3 the magnetosheath, 4 the foreshock, and 5 the undisturbed solar wind.

Fig. 48. Upstream waves at low frequencies observed by the Pioneer Venus magnetometer, compared with similar waves detected by the ISEE-1 magnetometer at Earth.

While this shows the existence of upstream hydromagnetic waves in front of Venus, it does little to answer the question of whether the venusian waves are different from the terrestrial waves. To do this we examine the Pioneer magnetometer data which are available at higher time resolution and for a greater time interval. Figure 48 shows four different intervals of upstream ultra-low-frequency waves at Venus and their terrestrial counterparts (M. M. Hoppe, personal communication, 1981). At the Earth in the IS EE data all four of these wave types are associated with energetic ions flowing back from the bow shock. These ion beams have properties that differ as a function of position. At the leading edge of the foreshock region, the region in which the ion beams are seen, the ion beams are narrow in both energy and pitch angle. Here the waves are weak or not seen. If waves are seen they are similar to those in the left-hand panels. The high-frequency waves seen at Venus are larger than at Earth; this is a general relationship, of which the examples shown here are typical.

Deeper into the foreshock region low-frequency transverse waves are often seen with high-frequency (~ 1 Hz) waves superimposed on the low frequencies, as in the second pair of panels. Still deeper into the foreshock the beams become spread in pitch angle and energy and the low-frequency waves steepen, as in the third pair of panels. Finally, far behind the foreshock boundary the beams become very diffuse and the waves very steep with leading (trailing in the spacecraft frame) wavetrains called discrete wave packets. We caution that this comparison was made from a very large data base without regard to frequency of occurrence of events. Thus, while these examples are typical of the waves at Venus and the Earth when they occur, this does not imply that the occurrence rates are the same; an occurrence rate study has not been done.

Fig. 49. Observed frequency of upstream low-frequency waves upstream of the bow shocks of Mercury, Venus, Earth, and Jupiter as a function of interplanetary field strength (Hoppe and Russell 1982).

The upstream waves occur at slightly higher wave frequencies at Venus than at Earth, but slightly lower than at Mercury. In Fig. 49 we show the wave period of the low-frequency waves, shown in the middle two panels at Mercury, Venus, Earth, and Jupiter as a function of the interplanetary field strength (Hoppe and Russell 1982). The linear dependence of wave frequency on field strength is as would be expected if the generating ion beams had the same energy at each planet. This relationship says nothing about the flux in the ion beams, only that the generation mechanism generates roughly the same energy at each planet. Note that the calculated energy of the beams is slightly but not significantly less for Venus and Mercury than for Earth and Jupiter.

Finally, we note that in the upstream solar wind VLF plasma waves associated with the electron and ion beams fiequently occur (Scarf et al. 1980b). These have potential to help probe the properties of the beams, but as yet a morphological survey of the waves in this region has not been performed.

E. Summary

The observations clearly show that Venus has a well-developed bow shock, in many respects similar to the bow shocks of the other planets. It apparently generates upstream particle beams and their associated waves. It depends on interplanetary parameters and the orientation of the interplanetary magnetic field direction much like any other bow shock. Nevertheless, there are important differences. The Venus shock does not stand off from Venus sufficiently to allow all the postshock flow to pass around the planet. The shock jump in the magnetic field strength is not as large as expected. The best inferred Mach number in fitting the shock position and jump is less than the observed solar wind conditions imply that it should be. The overshoot at the Venus shock is somewhat smaller than that seen under similar solar wind conditions at the other planets. These differences imply that something else is occurring at Venus in addition to magnetohydrodynamic processes, probably charge exchange with and photoionization of the neutral atmosphere in the postshock flow.

Our understanding of the solar wind interaction with Venus has come a long way since the initial Mariner 5 and Venera 4 measurements in 1967. In many ways we understand the Venus interaction as well or even better than the terrestrial interaction. The reconnection process between the magnetized solar wind and the Earth's magnetosphere has been difficult to understand and has been a source of controversy for two decades. The deflection of the solar wind by Venus seems much more straightforward, although our initial guesses were incorrect. We also might claim to understand the origin of the magnetic field in the Venus ionosphere better t