The Dynamics of Planetary Magnetospheres

Planetary and Space Science 49 (2001) 10051030www.elsevier.com/locate/planspasci

The dynamics of planetarymagnetospheresC.T. Russell

Department of Earth and Space Sciences and Institute of Geophysics and Space Physics, University of California, Los Angeles, 3845 Slichter Hall,MS 156704, Los Angeles, CA 90095, USA

Received 8 December 1999; received in revised form 9 October 2000; accepted 10 October 2000

Abstract

Mercury, Earth, Jupiter, Saturn, Uranus, Neptune, and the moon, Ganymede, have presently-active internal dynamos while Venus, Mars,at least two of the Galilean moons, the Earths moon, comets and asteroids do not. These active dynamos produce magnetic 5elds that havesu6cient strength to stand o7 the pressure of the exterior plasma environment. Because of changes in these exterior plasma environmentsthese magnetospheres are very dynamic. The jovian magnetosphere includes a strong time-varying energy source that adds to the dynamicsof its magnetosphere and produces a quite di7erent circulation pattern than that found at Earth and, presumably, Mercury. Not only intrinsicplanetary magnetic 5elds produce magnetospheres but also unmagnetized planets. Venus, Mars and comets have induced magnetospheresassociated with the solar wind interaction with their atmospheres. Cometary magnetospheres, parts of which can be remotely sensed,exhibit spectacular disruptions called tail disconnections. Even the atmosphereless bodies with weak magnetic 5elds can interact with thesolar wind. Small magnetic anomalies on the moon and possibly asteroids cause weak de>ections of the solar wind. The dynamics of thesevarious magnetospheres provide a rich spectrum of behavior which we review herein. c 2001 Elsevier Science Ltd. All rights reserved.

1. Introduction

At 5rst sight the bodies of the solar system can be classi-5ed into three types according to their plasma environments:the atmosphereless bodies, with no intrinsic magnetic 5eldsuch as the Earths moon; the bodies with atmospheres likeVenus and Mars but with a weak or absent intrinsic 5eldwhere the interaction is governed by the interaction of theplasma with the neutral atmosphere; and the bodies such asMercury, Earth and Jupiter, in which the intrinsic magnetic5eld is so strong that it de>ects the external plasma at alti-tudes far above the surface. The magnetic 5elds of the bod-ies whose 5elds are derived from internal dynamos or, inthe case of the moon, from magnetized rocks are steady onthe time scales of space exploration but their plasma envi-ronments are not. Thus each of the solar system bodies ex-ists in a very dynamical plasma and magnetic environmentbecause of changes external to the body.We have su6ciently explored these bodies to be able to

recognize a continuum of behavior over the spectrum ofplanetary interactions. Of course, we understand best thebehavior of just a few bodies: the terrestrial and jovian

Corresponding author. Tel. +1-310-825-3188; fax: +1-310-206-8042.E-mail address: [email protected] (C.T. Russell).

magnetospheres, the Venus and cometary magnetospheres,and the lunar interaction. We stress herein the informationgleaned from our studies of these bodies. However, we at-tempt to be complete by including the brief glimpses wehave gained into the behavior of other magnetospheres aswell. We emphasize that our purpose is to compare the dy-namical behavior of the solar system magnetospheres andthus our discussion of the basic structure of these magneto-spheres will be brief. By dynamical behavior we mean thecirculation and motion of the magnetized plasma on 5eldlines that thread the body and especially the sudden changesin that circulation and plasma content.We begin our review with the Earth because we under-

stand best the behavior of the Earth. We then examine thesolar wind interaction with the Moon and Mercury, two ba-sically atmosphereless bodies, and the magnetospheric inter-action with Ganymede, an atmosphereless moon in Jupitersmagnetosphere. Next we examine induced magnetospheres,including those of Venus and Mars and comets. Finally, weclose with a discussion of the jovian magnetosphere andwith a few words about the magnetospheric dynamics ofthe remaining outer planets. We stress that our goal is toprovide a tutorial overview of the dynamics of planetarymagnetospheres with an emphasis on the jovian magneto-sphere to provide context for the other papers associated withthe 1999 symposium on the Magnetospheres of the Outer

0032-0633/01/$ - see front matter c 2001 Elsevier Science Ltd. All rights reserved.PII: S 0032 -0633(01)00017 -4

1006 C.T. Russell / Planetary and Space Science 49 (2001) 10051030

Planets. This paper is not intended as a complete review ofthe jovian, terrestrial or other literature.

2. The Earth

The Earth has a dipole magnetic moment of 81015 Tm3that produces a magnetic 5eld strength at the equator on theEarths surface of about 30; 000 nT, and at 10 Earth radii(RE) of about 30 nT. The solar wind interaction slightly morethan doubles this value on the dayside so that the pressurein the magnetic 5eld is about 2 nPa. The sun emits a mag-netized plasma consisting of mainly protons and electronswith a density of about 7 cm3 at the orbit of the Earth (1astronomical unit or AU) at a velocity of about 440 km s1.The pressure exerted by this >owing plasma is also about2 nPa, thus balancing the pressure exerted by the magneto-spheric 5eld.Disturbances of three types propagate in this magnetized

solar wind plasma. The fast mode wave compresses the mag-netic 5eld and plasma; the intermediate mode wave bendsthe >ow and magnetic 5eld but does not compress it; andthe slow mode wave rare5es the 5eld while it compressesthe plasma and vice versa. The solar wind travels faster thanthe propagation speed of all three of these waves so whenit reaches the Earths magnetosphere the pressure wavesneeded to de>ect the solar wind plasma cannot propagateupstream into the solar wind without creating a shock front.The geometry of this shock, the de>ected >ow and the

magnetopause is shown in Fig. 1. The fastest wave is theaptly named fast mode wave. It does the yeomans work inslowing, de>ecting and heating the solar wind downstreamof the bow shock so that the plasma can >ow around the mag-netosphere. Nevertheless it cannot cause all of the changesin the plasma needed to move both the plasma and the mag-netic 5eld around the bullet-shaped magnetosphere and theintermediate and slow modes also play a role. The net resultof these standing waves is a >ow that bends to >ow paral-lel to the magnetopause, the boundary between the magne-tosheath and the magnetosphere. The pressure normal to thesurface is transmitted by the thermal motions of the plasmaand by the magnetic 5eld.

2.1. The size of the magnetosphere

In order to determine the scale size of the magnetospherewe need to understand the pressure applied to the magne-tosphere by the solar wind. Fortunately, we do not have tosolve the complex non-linear solar wind interaction problemto do so. We can obtain a quantitative formula for the dis-tance from the center of the Earth to the magnetopause in astraightforward manner. Conservation of the momentum ina stream tube of varying cross-section, S, gives us

(u2 + nkT + B2=20)S = constant; (1)

Fig. 1. The cross-section of the magnetosphere and the bow shock inthe plane of the incoming >ow to the stagnation point and the magnetic5eld. The dashed lines show the >ow and how it is de>ected at the shockand is bent around the magnetopause. The solid lines are magnetic 5eldlines that also are straight until they encounter the shock at which pointthey are bent around the magnetopause. The shaded region is the ionforeshock where ions move back from the bow shock into the solar wind.This backward re>ection can occur because of the inability of some ionsto penetrate the shock barrier and to re>ect back along the 5eld linesand due to the thermal heating in the shock=magnetosheath that resultsin some particles moving backward at greater than the incoming bulkvelocity. The geometry of this 5gure is patterned after early gas dynamicsimulations by Spreiter et al. (1966).

where ; u; n; T and B are the mass density, speed, num-ber density, temperature of the solar wind and magnetic 5eldstrength, respectively. This formula allows us to use the in-coming solar wind dynamic pressure, u2, which dominatesover the thermal and magnetic pressures, in front of the bowshock instead of having to calculate these pressures in themagnetosheath downstream from the shock front, given thatwe know the expansion of the cross-section of the streamtube, S.It is instructive to compare the size of the terms in (1).

The ratio of the 5rst two terms is

u2=nkT = u2=(kT=mi) = u2=c2s = M2s ; (2)

where is the polytropic index, cs is the sound velocity andMs is the sonic Mach number. The ratio of the second 5rstand third terms is

u2=(B2=20) = 2u2=(B2=0) = 2u2=v2A = 2M2A; (3)

where MA is the Alfven or intermediate Mach number.

C.T. Russell / Planetary and Space Science 49 (2001) 10051030 1007

Finally, the ratio of the second and third terms is

nkT=(B2=20) = = 2(MA=Ms)2=: (4)

Since, as stated above, the solar wind >ows faster than any ofthe three waves in the plasma and usually much faster, Machnumbers are much greater than unity. Thus Eqs. (2) and (3)tell us that the dynamic pressure dominates over the thermaland magnetic pressures in front of the bow shock. Eq. (4) in-dicates that, when the magnetic pressure dominates (low ),the speed of Alfven waves, vA, exceeds that of sound waves,cs, and magnetic forces dominate in the plasma frame. Theintermediate wave propagates at the Alfven speed along themagnetic 5eld. The fast mode wave propagates at a speed0:707{c2s + v2A + [(c2s + v2A)2 4c2s v2A cos2 ]1=2}1=2 where is the angle between the magnetic 5eld and the direction ofpropagation of the phase fronts of the wave. Perpendicular tothe magnetic 5eld this speed is equal to (v2A+c

2s )

1=2. The slowspeed is 0:707{c2s + v2A [(c2s + v2A)2 4c2s v2A cos2 ]1=2}1=2.The fast mode is the only mode that can transmit energyacross a magnetic 5eld.Returning to the question of the stando7 distance of the

nose of the magnetopause, we now know that we can approx-imate the solar wind pressure contribution by the momentum>ux u2 diminished by a factor accounting for the expan-sion of the stream tube. This e7ect is small, roughly 10%.The magnetospheric pressure is dominated by the pressurein the magnetic 5eld. The pressure equals (aB0=L3mp)

2 wherea is a shape-dependent factor, equalling 2:4 for the shapeof the Earths magnetosphere, and Lmp being the distanceto the magnetopause from the center of the Earth. Equatingthe pressure to the solar wind dynamic pressure we obtaina stando7 distance.

Lmp = 107:4(nsw u2sw)1=6; (5)

where nsw is the solar wind proton number density incm3; usw is the proton bulk velocity in km=s, and Lmp isthe stando7 distance in RE.

2.2. Tangential stress

The dynamic pressure determines the overall size of themagnetosphere and to zeroth order its shape but tangentialstresses also a7ect the shape and cause momentum trans-fer across the boundary. Several di7erent mechanisms havebeen proposed for the source of tangential stress andmomen-tum transfer to the magnetosphere. Fig. 2 illustrates a num-ber of the popular mechanisms. The upper left-hand panelillustrates di7usive entry. An ion enters the magnetosphericmagnetic 5eld and instead of returning to the magnetosheathwith the same velocity with which it started, it becomes scat-tered and drifts within the magnetosphere carrying with itwhatever momentum parallel to the boundary it had initially.This process relies on scattering centers that seem not tobe present in the boundary layers inside the magnetosphere.The lower left panel shows a variant of di7usive entry in

Fig. 2. Sources of viscosity at the magnetopause. Schematic illustrationsof di7usive entry, impulsive penetration, surface wave induced momentumtransfer and the KelvinHelmholtz instability.

which an entire tube of magnetosheath plasma crosses themagnetopause into the magnetosphere. This mechanism isnot expected to be e7ective because as long as there is a 5niteangle between the magnetosheath and magnetospheric 5eldsthe tubes cannot penetrate one another. If they do becomealigned in some region, the three-dimensional geometry ofthe interaction causes them to be at some signi5cantly largeangle not far away from the point of alignment. Thus thetwo magnetized plasmas are kept separate.The top right-hand panel illustrates momentum transfer

by wave processes. In a dissipative medium, such as theionosphere to which the magnetospheric 5eld lines are con-nected, the eddies formed by the passage of the surfacewave are increasingly smaller with distance from the bound-ary. The net result is a >ow in the magnetosphere paral-lel to the magnetosheath >ow. Finally, the lower left panelshows the result of boundary-wave amplitude growth viathe KelvinHelmholtz instability when the magnetosheathvelocity passes an instability threshold. The boundary shapebecomes non-sinusoidal and the momentum transfer by theprocess discussed above proceeds at even a greater rate.This process takes place independently of whether the so-lar wind magnetic 5eld is parallel or antiparallel to that of


Fig. 3. Reconnecting magnetospheres for southward interplanetary mag-netic 5elds (top) and northward interplanetary magnetic 5elds (bottom)(Dungey, 1961, 1963). The diagrams are not to scale nor are the detailsof the solar wind interaction taken into account.

the magnetosphere, but the tangential stress on the magne-tosphere clearly depends on the magnetic 5eld orientation.There seems to be little momentum transfer that is indepen-dent of control by the north-south component of the inter-planetary magnetic 5eld.The mechanism by which the magnetized solar wind pow-

ers the magnetosphere was 5rst proposed by Dungey (1961,1963) as sketched in Fig. 3. In the top panel the interplane-tary magnetic 5eld is southward and becomes connected tothe terrestrial magnetic 5eld at the subsolar point in a pro-cess known as reconnection. The resulting V-shaped mag-netic 5eld accelerates plasma (whose origin is both in thesolar wind and the magnetosphere) as the 5eld lines straight-ens. Then the magnetic 5eld slows the plasma as the 5eldlines are stretched behind the terminator. Over the daysideenergy >ows from the magnetic 5eld into the plasma, but inthe tail there is a Poynting >ux of energy into the magnetic5eld from the solar wind plasma. This process results in thestorage of energy in the magnetotail. This energy is in turntapped at a reconnection point in the tail that causes the >owof plasma into the magnetosphere proper and back downthe tail. The plasma and 5eld continue to move toward thedayside reconnection point where the cycle repeats itself.In this way the magnetospheric plasma can be made to cir-culate even in a dissipative system as energy is continuallysupplied by the solar wind. Estimates of the rate of energyinput into the magnetosphere during active times range upto about 2 TW (2 1012 W).When the interplanetary magnetic 5eld is northward re-

connection can still occur but it has a quite di7erent e7ect onthe magnetosphere. The panel on the bottom of Fig. 3 showsthis situation. The interplanetary magnetic 5eld now recon-nects with the terrestrial magnetic 5eld above and behind

Fig. 4. A cut away diagram illustrating the three-dimensional magneto-sphere, its plasma regions and current systems.

the poles. The reconnected 5eld line is added to the daysideand a corresponding >ux tube is removed from the night-side and transported down the tail. E7ectively this transportsmagnetic >ux from the nightside to the dayside magneto-sphere. The reconnected dayside tube moves along the mag-netopause boundary (in three dimensions out of the planeof the page) and replaces the >ux tube that was lost downthe tail. This mechanism makes a boundary layer of plasmawithin the magnetopause and maintains circulation of theplasma for northward interplanetary magnetic 5eld. Whenthe magnetic 5eld is northward, but not due northward, thisprocess can still proceed, but the same >ux tube is unlikelyto reconnect at both ends. This results in plasma circulationbut no change in the magnetic >ux on open and closed mag-netic 5eld lines (those with 1 and 2 feet in the ionosphererespectively).The tangential stress in the outer magnetosphere must

ultimately apply stress to the ionospheric plasma as the iono-sphere is the ultimate site of dissipation in the magneto-spheric system. Although there are several ways for energyto be deposited in the ionosphere from the magnetosphere,such as particle precipitation whereby energetic charged par-ticles in the magnetosphere enter the atmosphere and arelost, or as wave transport in which waves generated in themagnetosphere propagate into the ionosphere and heat it,the most signi5cant energy dissipation mechanism is jouledissipation in electric currents. Fig. 4 illustrates where thecurrent systems in the magnetosphere >ow. When the forcesof the solar wind are directed along the magnetopause nor-mal and there are no tangential stresses, the magnetopausecurrent and the tail current are dissipationless currents >ow-ing on the surface of the magnetopause with j E=0 wherej is the current density and E is the electric 5eld. Whenthere is a magnetic 5eld normal to the magnetopause thesecurrents can accelerate or decelerate the >ow in the day-side and tail regions respectively. The current crossing the


plasma sheet, labelled neutral sheet current, is an extensionof the tail current system. It too can lead to plasma ac-celeration when there is a normal component of the mag-netic 5eld across the tail current sheet. The magnitude ofthe magnetopause current is such as to rotate the 5eld di-rection and increase its strength from that of the magne-tosheath to that of the magnetosphere. The neutral sheetcurrent strength is of magnitude to reverse the direction ofthe magnetic 5eld from the northern lobe of the tail to thesouthern lobe.

2.3. The ring current

The ring current consists of the current due to theeastward (electron) and westward (proton) drift in the ra-diation belts. Since this current does not pass through theionosphere, it too is basically dissipationless. The energycontent of the radiation belt is generally fairly constant ex-cept during periods known as geomagnetic storms. The ringcurrent causes a net decrease in the magnetic 5eld on thesurface of the Earth (Dessler and Parker, 1959; Sckopke,1966) as opposed to the magnetopause current that causesan increase. The energy of these circulating particles can beeasily calculated from their e7ect on the ground-level mag-netic 5eld. In major magnetic storms this energy can reach10 or more petaJoules (1015 J) and the energization rate canexceed several teraWatts (1012 W), equaling and possiblyexceeding the energy dissipation in the auroral ionosphere.The e7ect of the tail current system is to oppose the Earths5eld and has a stronger e7ect on the nightside thus, causinga daynight gradient in the 5eld due to external sources.Thus a sudden compression of the size of the magneto-sphere by an increase in the solar wind pressure will causea greater e7ect on the dayside of the magnetosphere thanon the nightside.

2.4. Magnetosphereionosphere coupling

The current system ultimately responsible for the majorityof the dissipation is labeled 5eld-aligned current in Fig.4. This current system closes in the magnetosphere on thepressure gradients in the plasma and thus is controlled bythe magnetospheric stresses. In the ionosphere these currentscross the magnetic 5eld in the collisional ionosphere. Thesecollisions heat the atmosphere and the cross-5eld currentsaccelerate the ionosphere plasma against the drag of the at-mosphere. These currents attempt to maintain the >ow in theionosphere to follow the >ow in the magnetosphere. Whenthe stress in the magnetosphere increases, the bend in themagnetic 5eld increases and the current >ow along the 5eldlines increase. Fig. 4 describes a steady-state magnetosphere.However, transient events also can cause stress-induced cur-rent systems that close in the ionosphere. Substorms that aredescribed below are particularly important in causing suchcurrents.

Fig. 5. The coupling of the magnetosphere to the ionosphere via5eld-aligned currents. The current >ows across magnetic 5eld lines inboth the magnetosphere and the ionosphere and along the magnetic 5eldin between. This causes the J B force applied to the magnetosphericplasma at high altitudes to be applied to the ionosphere. The electric 5eldassociated with the motion of the plasma times the magnetic 5eld per-turbation provides a Poynting >ux, S, into the ionosphere (Strangewayet al., 2000).

Fig. 5 illustrates the physics involved in the coupling ofthe magnetosphere to the ionosphere via 5eld-aligned cur-rents. At the top of the 5gure is the magnetopause wherethe stress applied to the magnetosphere has pulled the topof a bundle of magnetic 5eld lines into the page. These 5eldlines are now slanted with respect to their neighbors. Thisshear in the magnetic 5eld is the equivalent of a parallel cur-rent. This current that began in the generator region of themagnetosphere closing across 5eld lines in a pressure gra-dient and acting in the direction to slow the solar wind, nowexperiences a load in the ionosphere where it applies astress to the ionospheric plasma in the same direction as thestress in the magnetosphere so that the ionosphere begins to>ow in the same direction as the solar wind. If the plasma>ows in response to this stress, then the product of the mag-netic 5eld distortion, B, and the electric 5eld, V B, is aPoynting >ux, S, into the ionosphere. It is this energy thatis dissipated in the ionosphere as Joule dissipation.In the Earths magnetosphere there are many types of au-

roral processes. Any process that causes the precipitation ofenergetic charged particles from the magnetosphere into theatmosphere in signi5cant numbers is likely to cause di7useaurora. However, the most intense discrete aurora is causedby electrons accelerated down into the ionosphere by paral-lel electric potential drops along magnetic 5eld lines in up-ward parallel electric currents. Thus most intense auroras arein regions of velocity shear in the magnetospheric plasma.

2.5. Geomagnetic storms

A geomagnetic storm occurs when the energy content ofthe radiation belts increases to unusually large values. Theconditions in the solar wind that lead to the generation ofgeomagnetic storms are rare. The interplanetary magnetic


Fig. 6. Development of a geomagnetic storm. The top panel shows thedynamic pressure of the solar wind as measured by the WIND spacecraft.The second panel shows the solar magnetospheric Z component of theinterplanetary magnetic 5eld as measured by the WIND spacecraft. Thethird panel is the interplanetary electric 5eld computed from the productof the radial solar wind velocity and the north-south component of theinterplanetary magnetic 5eld. The 5fth panel contains the AU index andAL index the di7erence between which is the AE index. These indices arecomputed from the maximum (upper) and minimum (lower) horizontalcomponent of the magnetic 5eld in the auroral zone and measure thestrength of auroral processes. The bottom panel shows the Dst indexcomputed from the average near equatorial surface 5eld.

5eld must be strong and steadily southward for several hours(Russell et al., 1974). Fig. 6 illustrates the solar wind con-dition and magnetospheric response during the build up of ageomagnetic storm that has been chosen because it clearlyillustrates the di7erent types of responses of the magneto-sphere to geomagnetic activity. The top panel shows thesolar wind dynamic pressure. This pressure compresses themagnetosphere and increases the 5eld strength on the sur-face of the Earth but has little other e7ect. In particular theDst index, or the average worldwide surface magnetic 5eldin the near-equatorial regions with the quiet-day 5eld re-moved, shows an increase in association with the pressureincrease but this increase signals not a change in the ringcurrent but a change in the magnetopause current. The sec-ond panel shows the north-south or solar magnetospheric Zcomponent of the interplanetary magnetic 5eld during theinterval. The magnetic 5eld is slightly southward at the be-ginning of the interval plotted and becomes strongly south-ward in the second half of the plot. This leads to currentsin the auroral zone as seen in the fourth panel from the top.This activity remains roughly constant over the entire pe-

riod plotted even though the character of the interplanetarymagnetic 5eld changes drastically.After the pressure pulse arrives the interplanetary

magnetic 5eld strength increases and begins to have strong>uctuations in the northsouth direction. This hardly a7ectseither the auroral currents as measured by the AE indexin the fourth panel or the ring current as registered by theDst index. Eventually the interplanetary magnetic 5eld be-comes steadily southward and strong and lasts this way forseveral hours. Now the ring current builds up (Russell etal., 1999a, b) and the Dst plunges to values more negativethan 200 nT.We can understand the injection of energetic plasma into

the ring current with a fairly simple model of the solar windmagnetosphere interaction (Burton et al., 1975). The basicpremise of this model is that the energy >ows into the ringcurrent from the solar wind at a rate proportional to theinterplanetary electric 5eld in the dawndusk direction butnot when the electric 5eld is in the opposite direction. Oncein the ring current, the energy exponentially decays with a5xed time constant so that a 5xed percentage of the ring cur-rent energy is lost per unit time. The Dst index, which is theaverage horizontal component of the magnetic 5eld aroundthe near equatorial region is a good measure of the ring cur-rent, but there are other contributions to Dst. In particularthe magnetopause currents contribute to Dst as well. Burtonet al. (1975) proposed the following recipe for calculatingDst from the solar wind parameters:

dDst=dt = F(Ey) aDst0 ;Dst0 = Dst b(P)1=2 + c;F(Ey) = d(Ey 0:5); Ey 0:50 mV=m;F(Ey) = 0; Ey 0:50 mV=m;

where a = 3:6 105 s1; b = 15:8 nT=nPa1=2; c = 20 nT;and d= 1:5 103 nT=(mv m1 s).This recipe states that the rate of change of the energy in

the ring current increases due to an energy coupling func-tion F(Ey) and decreases by a 5xed percentage each minutebecause of loss processes. The ring current itself is Dst0 anddi7ers from the surface 5eld disturbance by contributionsproportional to the square root of the solar wind dynamicpressure, P, and the quiet day ring current, c. The couplingfunction F(Ey) extracts energy from the solar wind when thesolar wind electric 5eld exceeds 0:50 mv m1 in the dawndusk direction and is zero otherwise.The energy contained in the ring current according to

the DesslerParkerSckopke formula (Dessler and Parker,1959; Sckopke, 1966) is

ERC (J ) = 2:8 1013B (nT)so that the 200 nT storm shown in Fig. 6 contains about 61015 J (6 PJ), which it was losing to the Earths atmosphereat a rate of about 2 1011 W or 0:2 TW. At the peak ofthe injection of power into the ring current the injectionrate was about 1012 W or 1 TW. We note that this buildup


of energy was not predictable from the currents in the au-roral zone as can be seen by the lack of a clear relationshipbetween the quantities plotted in the fourth and 5fth panels.We note also that during disturbed solar wind conditions theconvected magnetic energy in the solar wind, that is the so-lar wind Poynting >ux, is about a teraWatt integrated overthe entire dayside magnetopause. Thus reconnection wouldhave to be 100% e7ective for the Poynting >ux to power amagnetic storm. Instead the magnetosphere taps a small frac-tion of the 60 TW of mechanical energy >ux that the solarwind convects toward the dayside magnetopause under dis-turbed conditions. We also note that studies of the e6ciencyof the interplanetary electric 5eld for causing geomagneticactivity suggests that the rate of reconnection diminishes forvery high solar wind Mach numbers when the beta of themagnetosheath becomes large and the magnetic 5eld weak(Scurry and Russell, 1991). This could be important at Earthduring times of extremely intense solar wind disturbancesbut should be more important in the outer heliosphere wherehigh Mach number conditions are more prevalent.

2.6. Substorms

Strictly speaking the Dungey model for the solar wind in-teraction sketched in Fig. 3 is a steady state model, but it canreadily be converted (e.g. Russell and McPherron, 1973) tobe a time-varying model as shown in Fig. 7. The top panelshows the Dungey model for southward interplanetary mag-netic 5eld. The magnetic 5eld has just turned southward and>ux is being eroded from the dayside magnetosphere andmoved into the polar cap or tail lobe region. Thus, at the be-ginning of the substorm period, the merging rate, M , goesup, decreasing the >ux !day in the bottom panel. Until re-connection begins between the two tail lobes, the magnetic>ux in the tail, !lobe, will increase. When the reconnectionrate, R, in the center panel climbs rapidly, then the sub-storm onset has occurred and that >ux is returned 5rst tothe plasma sheet and then to the dayside. Almost by def-inition the substorm involves the release of energy that ismuch more rapid than its accumulation time. Finally, wenote that even though its name suggests that a substorm isa small storm or a process that leads to a storm, there ap-pears to be little connection between the processes that leadto substorm storage and release of energy into the auroralzone illustrated in Fig. 6 and the storm process. The sub-storm appears to intimately involve the tail for storage andrelease of energy. The storm appears to be associated withthe penetration of the solar wind induced plasma circulationdeep into the magnetosphere.

3. The moon and asteroids

The Earths moon and the asteroids are too small tohave currently active dynamos because any internal liquidcore has long since cooled and ceased its dynamo action.

Fig. 7. Schematic illustration of a substorm. The top panel shows theDungey magnetosphere when the IMF has just turned southward. In thecenter panel the time variation of the reconnection and convection ineach of the three locations: dayside reconnection, tail reconnection andconvection to the dayside. The bottom panel shows the variation of totalmagnetic >ux in each of the three regions.

Nevertheless, such bodies could have remanent magnetism,created when the surface material cooled while an early dy-namo was operating. Even those bodies that are too smallto ever have had a core may be part of an earlier largerbody that possessed a core. Thus a priori we cannot ruleout the presence of some magnetized regions on any ofthese bodies. The Moon proves this point. Measurementsfrom the Explorer 35 and Apollo subsatellite magnetome-ters show that, when magnetized regions interact with thesolar wind >ow past the Moon, disturbances are created inthe solar wind >ow called lunar limb compression (Col-burn et al., 1971; Russell and Lichtenstein, 1975). Recently,these results were con5rmed by the Lunar Prospector mis-sion (Lin et al., 1998). Thus the solar wind does inter-act strongly with magnetic regions at least down to thesize of the proton gyro radius. Reports of analogous distur-bances from the asteroid, Gaspra, suggest that these distur-bances can be created for even smaller scale-size features(Kivelson et al., 1993).It is worth reviewing the behavior of a large unmagnetized

body, such as most of the Moon appears to e7ectively be,that is, too weakly magnetized to prevent the solar windhitting the surface. This interaction is shown in Fig. 8 for an


Fig. 8. The solar wind interaction with the Moon when the interplanetarymagnetic 5eld is perpendicular to the solar wind >ow. The solar windis completely absorbed on streamlines that intersect the Moon, leaving acavity on the downstream side that 5lls by ion motion along the magnetic5eld at the ion thermal velocity. Because of the charge neutrality conditionin the plasma the electrons move with the ions. In MHD terms the regionin which the plasma is moving toward the wake is called an expansionfan (Spreiter et al., 1970).

interplanetary magnetic 5eld perpendicular to the solar wind>ow. Not shown is the >ow-aligned case that occurs muchmore rarely. In both cases the >owing plasma is absorbedby the moon leaving an empty wake behind the Moon. Inthe aligned->ow case the plasma cannot >ow into the cavitybehind the moon but the wake does narrow to a diameterless than that of the moon. In the case with the interplanetarymagnetic 5eld perpendicular to the >ow, the plasma closesbehind the Moon at the ion thermal velocity. Since the ionsare much more massive than the electrons and since chargeneutrality requires electrons and ions to stay together in thesolar wind, ion motion governs the electrons as well.An important aspect of this interaction is the electric 5eld.

The solar wind is a >owing, magnetized plasma and hencehas an electric 5eld in the frame of reference of the Moon.Thus ions produced on one side of the moon by photoion-ization of its tenuous atmosphere will be accelerated downon to the surface, while on the other side ions will be re-moved from the moon (Freeman and Ibrahim, 1975). In thisway the solar wind electric 5eld both implants ions into thelunar surface and removes them from the lunar atmosphere.However, the currents through the body of the Moon, drivenby this electric 5eld, are very, very small because of the ex-tremely low electrical conductivity of the lunar surface. Thesolar wind does cause currents in the interior of the moonby carrying a spatially varying magnetic 5eld past the moonthat the moon sees as a time varying magnetic 5eld and thatinduces a voltage across the moon. These currents >ow en-tirely within the moon and do not penetrate the crust. Fi-nally, we note that Mars tiny moons Phobos and Diemoshave been reported to cause disturbances in the solar wind(Riedler et al., 1989; Dubinin et al., 1990; Sauer et al., 1998)but since these moons orbit close to the bow shock whenthey are in the solar wind it is di6cult to separate lunar fromplanetary e7ects.

Fig. 9. The average con5guration of the magnetic 5eld in the Mercurymagnetosphere as drawn in the noon-midnight meridian based on theMariner 10 >ybys. (Russell et al., 1988).

4. Mercury

To the non-specialist Mercury looks much like the Moon.It has a cratered surface and no signi5cant atmospherebut unlike the Moon it has a magnetic 5eld that de>ectsthe solar wind well above the surface. The magnetic 5eldcon5guration in the noon-midnight meridian is shown inFig. 9 as inferred from two >ybys by Mariner 10 in 1974and 1975. Some recon5guration of the magnetosphere wasdetected on the 5rst >yby and interpreted in terms of amagnetospheric substorm as on Earth (Siscoe et al., 1975),but, since Mercury has no signi5cant ionosphere, stressesmight be communicated much more rapidly in the Mer-cury magnetosphere than in the terrestrial magnetosphere.Under the assumption that Mercurys magnetosphere wasresponsive to the interplanetary magnetic 5eld orienta-tion in a manner similar to that on the Earth, Luhmannet al. (1998) modi5ed Tsyganenkos (1996) terrestrial mag-netic 5eld model to apply to Mercury. Fig. 10 shows theequivalent magnetic 5eld models for three IMF conditionsobtained by Luhmann et al. (1998). They then assumed thatthese model 5elds were immediately attained when the IMFchanged and calculated what IMF conditions would createthe magnetospheric conditions observed. Their conclusionwas that the dynamics of the Mercury magnetosphere couldbe directly driven with little or no storage of energy in themagnetic tail, unlike the terrestrial magnetosphere.


Fig. 10. The Mercury magnetic 5eld in the noon-midnight meridian if the5eld is responsive to the IMF in a manner similar to that of the terrestrial5eld (Luhmann et al., 1998).

5. Ganymede

The last intrinsic magnetosphere in a terrestrial-sized bodythat we discuss is that of Jupiters moon, Ganymede, thatsits well inside the jovian magnetosphere at 15 jovian radii(RJ) in a plasma of density about 4 cm3 >owing about130 km s1 relative to Ganymede. Its magnetic moment is1:4 1013 Tm3 (Kivelson et al., 1997a) and it sits in abackground magnetic 5eld of strength 100 nT. The staticpressure of the Jovian magnetic 5eld is about 4 nPa andthat of the >owing plasma is 2 nPa so that the Ganymede

Fig. 11. The con5guration of the Ganymede magnetosphere in Jupitersmagnetic meridian when Ganymede is at high jovian magnetic latitudes(after Kivelson et al., 1996a).

magnetosphere is more a7ected by the external (jovian)magnetic 5eld than the corotating magnetosphere plasma.This 5eld rocks back and forth mainly in the meridian planeas Ganymede moves back and forth across the magneticequator. This alters the high altitude magnetic 5eld directionshown in Fig. 11, but the low altitude magnetic 5eld remains5xed. There are no reports yet of substorms or storms in thistiny magnetosphere.

6. Venus, Mars and comets

Venus and comets are completely devoid of intrinsic mag-netic 5elds. Mars has some signi5cant regions of remanentmagnetic 5eld (Acuna et al., 1998) but not so strong thatit dominates the solar wind interaction with the planet. Thesolar wind interaction for Venus and Mars is basically assketched in Fig. 12 and described in detail in the volumeVenus Aeronomy (Russell, 1991) and Venus and Mars: At-mospheres, Ionospheres and Solar Wind Interaction (Luh-mann et al., 1992). Solar EUV shining on the planetaryatmosphere creates an ionosphere. While the ions are pro-duced from the atmosphere over a wide altitude range, theycan recombine only at low altitudes where there are colli-sions. This sets up a circulation pattern in the ionospherethat is downward in the neighborhood of the subsolar pointand toward the antisolar point at other solar zenith angles.The thermal pressure of the ionosphere is generally greater

than the dynamic pressure of the solar wind so that it can de->ect the solar wind prior to the solar wind hitting the atmo-sphere. When the solar wind pressure is low this de>ectionoccurs well above the collisional regime in the ionosphereand there is a thin layer of current separating the solar windand the ionosphere so that the magnetic 5eld penetrates very


Fig. 12. The con5guration of the magnetic 5eld in the Venus and Marsionospheres for low and high dynamic pressures (after Saunders andRussell, 1986).

little into the ionosphere. In the case of a high solar winddynamic pressure the interaction takes place at a lower col-lisional altitude and a thick current layer is formed withmagnetic 5eld deep into the ionosphere. In both cases thesupersonic solar wind is de>ected with a bow shock muchlike the Earth but, when the EUV is weak at solar minimum,the solar wind apparently is able to interact directly withthe atmosphere and is absorbed. This absorption allows theshock to move closer to the planet.Some ions are created in the >owing solar wind itself

from a hot oxygen exosphere whose high thermal velocityis su6cient to carry neutral atoms to high altitudes. As atthe Moon the solar wind electric 5eld is responsible foraccelerating these ions. Some of them get accelerated backto the planet. Others are accelerated away from the Venusand have been detected as far away as Earth (Grunwaldtet al., 1997). The geometry of the magnetic 5eld lines inthe tail is in the sense to accelerate ions via the J Bforce away from the sun (McComas et al., 1986). Someevidence for substorm-like activity in the tail or at least theproduction of energetic ions in the tail of Venus has beenpresented by Vaisberg et al. (1994). The loss rate of ionsfrom Venus is poorly known but it may be close to 1024

oxygen ions per second. To accelerate these ions to the solar

wind velocity requires about 3 GW, a number that is a smallpercentage of the over 100 GW incident on Venus in solarwind mechanical energy >ux.The solar wind interaction with Mars much resembles

that with Venus. However, the solar wind magnetic 5eld isweaker and the radius of Mars smaller than at Venus. Hencewhile MHD provides a suitable model for the solar windinteraction at Venus, kinetic e7ects become important atMars (Brecht, 1997). Also, while there is no global planetary5eld there are strong localized 5elds that also cause somedi7erences with the Venus interaction (Acuna et al., 1998).A more detailed discussion of the solar wind interactioncan be found in the article by Vaisberg et al. (1990) andthe volume edited by Luhmann et al. (1992). Evidence forsubstorm-like phenomena have not yet been reported forthe Martian tail but the Mars plasma environment continuesto be studied. Data are still being returned from the MarsGlobal Surveyor and the Japanese Nozomi spacecraft is dueto enter Mars orbit in late 1993 (Yamamoto and Tsuruda,1998).The European Giotto and Russian Vega 1 and 2 space-

craft directly probed the coma of comet Halley and theISEE-3 mission, renamed ICE, has probed the coma ofcomet Giacobini-Zinner. In addition Giotto encountered asecond comet, Grigg-Skjellerup. The cometary plasma envi-ronment has been described in detail in the volume CometaryPlasma Processes (Johnstone, 1991). The cometary interac-tion covers a larger volume of space than the interaction withVenus and Mars because of the lack of signi5cant gravita-tional pull on the neutral gas that escapes from the nucleusat over 1 km=s and is ionized in about a day. These mea-surements represent only snapshots of the region around thecomet and to study the dynamics of the cometary plasmaenvironment one must rely on remote sensing data.These remote sensing data provide a strong case for a

substorm-like phenomenon. Fig. 13 shows an example of acometary tail disconnection event where it appears that thetail has been pulled out of the cometary coma. The brighthead of the cometary coma is thought to be much like theVenus ionosphere and magnetosheath but on a larger scalebecause of the absence of a signi5cant gravitational 5eld.As at Venus and Mars a magnetic-5eld-free region can formclose to the nucleus if there is su6cient outgassing andEUV to form an ionosphere that has su6cient pressure tostando7 the solar wind dynamic pressure. The ions createdin the solar wind >ow form a ring around the magnetic 5eldin velocity space that is unstable to the formation of bothion cyclotron waves and mirror mode waves. The formeroscillate at the gyro frequency of the heavy ions created.The latter form depressions in the magnetic 5eld strengthwith a diameter of several gyroradii.There are two outstanding theories of cometary tail dis-

connection events. The 5rst model assumes that the phe-nomenon involves reconnection in the cometary coma whenthe interplanetary magnetic 5eld reverses (Niedner andBrandt, 1978). The second assumes that reconnection in the


Fig. 13. A cometary tail disconnection event, seen in comet Morehouse1908a at 1943 UT, October 1, 1908 (Niedner and Brandt, 1978).

tail causes the disconnection as sketched in Fig. 14 (Russellet al., 1986). Just as with terrestrial substorms, such a taildisconnection may be triggered by a solar wind event.

7. Jupiter

Jupiter is the planet of solar system records. It hasthe fastest rotation, the most mass, the largest radius, thestrongest surface magnetic 5eld, largest magnetic moment,most intense radiation belts, strongest radio emissions,and the largest moon, Ganymede. Fittingly it also has thelargest magnetosphere, one so large that, if it could be seenfrom Earth, it would appear larger than the Earths moon.The size of the magnetosphere, like that of the Earth, isdetermined by a balance between the solar wind dynamicpressure and the pressure exerted by the magnetosphere.Part of the reason for the enormity of the jovian magneto-sphere is that the mass density of the solar wind and henceits dynamic pressure decreases as the inverse square of theheliocentric radius. By the time the solar wind has traveledfrom 1 AU to Jupiter the dynamic pressure has decreased afactor of 27. Thus, if Jupiter were like the Earth, this e7ectalone would increase the radius of the magnetosphere bya factor of 4. Under most conditions the Earths magnetic5eld can be approximated by a dipole. A dipole magnetic

Fig. 14. A tail reconnection model for the disconnection of cometary tails.

5eld is a force-free magnetic con5guration in which the out-ward magnetic pressure is balanced by the inward curvatureforce. Its 5eld strength decreases as the cube of the radius.The jovian magnetic 5eld is decidedly not force free andits interior pressure falls o7 more slowly than an inversecube in its outer portions. This e7ect increases the size ofthe magnetosphere relative to that of the Earth over andabove its factor of 18,000 greater magnetic moment. Theweaker outward pressure gradient also makes the size of themagnetosphere more sensitive to changes in the solar winddynamic pressure than the Earths magnetosphere (Slavinet al., 1985; Huddleston et al., 1998a). Thus the nose of thejovian magnetosphere has been found at distances from 40to over 100 R jovian radii (RJ).As we discuss in greater detail below, the additional

force in the jovian magnetosphere that is not present in theterrestrial magnetosphere is centrifugal force due to a verystrong source of plasma at the moon Io. This additionalstress stretches out the magnetic 5eld forming a magne-todisk beyond about 24RJ (Smith et al., 1975). The resultingmagnetic con5guration resembles the sketch of the jovian


Fig. 15. Magnetic 5eld lines in the noon-midnight meridian of the jovianmagnetosphere showing the current sheet in the magnetodisk region (afterRussell et al., 1998a, b).

magnetic 5eld in the noon-midnight meridian shown inFig. 15. As can be seen in this 5gure the nose of the mag-netosphere is sharper than that of the Earth. Just as theaerodynamic shape of a supersonic airplane allows the bowshock to form very close to the nose of that airplane, themore streamlined shape of the jovian magnetopause allowsthe bow shock to be formed closer to the magnetospherethan at Earth (Stahara et al., 1989).The existence of a variable source of mass in the inner

jovian magnetosphere provides an extra dimension to thedynamics of the jovian magnetosphere. There is possiblecontrol by the rate of mass addition as well as by the solarwind and the interplanetary magnetic 5eld. This mass addi-tion could a7ect the size and the shape of the magnetosphere.We do not yet know how variable is this mass-loading rate,so we cannot yet estimate how important this e7ect is on thesize of the magnetosphere. If mass loading were to totallycease we estimate that the magnetopause stando7 distancewould be only about 40RJ which is similar to the smalleststando7 distances seen, but these conditions also most prob-ably correspond to periods of higher than usual solar winddynamic pressure.As we discussed above, the Earths magnetosphere is very

much a7ected by the strength and orientation of the inter-planetary magnetic 5eld, or more correctly, the product ofthe solar wind velocity and the component of the magnetic5eld perpendicular to the solar wind >ow. While the mag-netic 5eld strength is almost a factor of 10 smaller at Jupiterthan at the Earth, the enormous size of the magnetospheremight compensate for this decrease. We can estimate the im-portance of the solar wind electric 5eld on a magnetosphereby comparing the solar wind electric 5eld, the product of

the magnetic 5eld perpendicular to the solar wind >ow andthe solar wind speed, with the corotational electric 5eld ofthe planetary magnetosphere that is equal to the corotationalspeed !R times the north-south component of the magnetic5eld. Since the corotational speed increases as R and themagnetic 5eld decreases as R3 (in a dipole) the electric 5eldof a rotating dipolar magnetosphere decreases as L2. Thusthe terrestrial corotational electric 5eld is 14L2 mV m1

and that of Jupiter 4900L2 mV m1 where L is the dis-tance in planetary radii. The solar wind electric 5eld at 1and 5:2 AU respectively is typically 3 and 0:4 mV=m. If allof this 5eld were able to penetrate the terrestrial and jovianmagnetospheres, the interplanetary and corotational 5eldswould be equal at 2:1RE and 100RJ respectively. Since atEarth only about 10% of the solar wind electric 5eld pene-trates the magnetosphere, the typical distance at which theelectric 5elds balance is 6RE. If the same rule applied toJupiter the balance point would be about 300RJ. In fact, wehave reason to believe that reconnection is even less e7ectiveat Jupiter than at Earth. While >ux transfer events, one man-ifestation of magnetopause reconnection, were observed atthe jovian magnetopause they were typically smaller and lessfrequent than on Earth (Walker and Russell, 1985). More-over, the reconnection is apparently less e6cient for highbeta conditions that occur behind high Mach number shocks(Scurry et al., 1994), and the jovian shock has a signi5-cantly higherMach number than the terrestrial shock. Finallyand most importantly, jovian auroral phenomena behave dif-ferently than terrestrial aurora (Clarke et al., 1996; Prangeet al., 1998). Jovian aurora rotate with Jupiter and are asso-ciated with the inner magnetodisk portion of the magneto-sphere. Unlike terrestrial auroras they do not cluster aboutthe boundary between open and closed 5eld lines. It is clearthat the jovian magnetosphere works much di7erently thanthe terrestrial magnetosphere.The electric 5eld associated with corotation arises be-

cause the ionosphere rotates with the atmosphere and the at-mosphere rotates with the planet. Since electrons can movequite freely along the magnetic 5eld, the magnetic 5eld linesare equipotentials and transmit this electric 5eld to the equa-tor regions. It is, of course, possible that this electric 5eldis altered in some way. If some process held the >ux tube5xed in the equatorial plane, it would either have to bendbecause it was also 5xed to the ionosphere, or it wouldhave to slip with respect to the ionosphere. If it slipped withrespect to the ionosphere, a potential drop would have toappear across the point where the >ux tube slipped. As dis-cussed for the Earth this velocity shear leads to intense au-rora. Thus, to zeroth order, auroral pictures of Jupiter maysimply show us where this slippage is taking place.

7.1. Mass addition at Io

Io is the engine that drives the jovian magnetosphere andmass addition is the fuel that powers the magnetosphere.


Fig. 16. Schematic diagram of the addition of ions to the Io torus (after Huddleston et al., 1998b).

In the following sections we treat 5rst the immediate vicin-ity of Io and then move out into the Io torus and then intothe middle and distant magnetosphere. The interaction of Iowith the jovian magnetosphere is shown in Fig. 16. Io or-bits Jupiter at 17 km s1 whereas corotating plasma wouldbe travelling at 74 km s1. Ions that are produced from neu-tral atoms at rest with respect to Io must be accelerated by57 km s1. If the mass loading rate is large and the gen-erally accepted 1 ton per second is large, the plasma mustslow down. Where there is good coupling of the ionosphereto the magnetospheric >ux tube, the slow down at the equa-tor can be taken up by a bending of the 5eld. In steady statethe bend in the magnetic 5eld can remain but the velocityof the plasma must reach the corotation level if there is noslippage. If the bend is to be erased the >ow must exceedthe corotational velocity for a while to catch up with theionosphere. The existence of an auroral spot, not only at themagnetic footpoint of Io, but extended in the direction of ro-tation suggests that the >ux tubes that pass near Io slip withrespect to the ionosphere for a long distance downstream ofIo (Prange et al., 1996).Fig. 16 also illustrates some important aspects of the dis-

tribution of the picked up charged particles about the mag-netic 5eld line. The initial motion of the charged particle isin a cycloid about the 5eld line. The thermal or gyro ve-locity of this motion is equal to its bulk velocity. BecauseIo is close to the magnetic equator, but not usually at theequator, there is a small velocity of the ion along the mag-netic 5eld. Initially the distribution of the ions can be rep-resented as two delta functions perpendicular to the 5eldbut with time the particles scatter and eventually becomeMaxwellian.

If the mass loading rate is large enough, then the additionof plasma measurably slows the >ow. This is seen at Io(Frank et al., 1996) and is illustrated in Fig. 17. The plasma isheated at the edge of the wake where newly created ions areadded to the >ow with a velocity equal to the >ow velocity.The decrease in the magnetic 5eld strength in the wake isdi6cult to explain simply by plasma diamagnetism. Themissing magnetic energy in the wake is up to 4:5 MeV=cm3

which would require a density of 22; 500 ions cm3 at atemperature of 200 eV. While the density in the wake is atits peak over 104 ions=cm3, the temperature is very cold,about 10 eV and does not explain the depression. Moreover,any slow down of the >ow and pile up of magnetic 5eldshould increase the 5eld in the near wake, not decrease it.The middle panel shows estimates of the source density

of new ions (Bagenal, 1997) based on the plasma measure-ments, and of speci5cally, N+SO2 , based on the amplitude ofthe ion cyclotron waves. The ions produced on >ux tubesthat move across Ios poles and into the wake region is asmall fraction of the total production, only a few tens ofkg s1 (Russell et al., 1997). The ions produced in the >owthat is de>ected by Io, as crossed by Galileo, is between180 and 580 kg s1 (Bagenal, 1997) signi5cantly less thanthe 1000 kg s1 usually quoted for the mass loading rate(Hill et al., 1983). Although re5ned estimates of the plasmaparameters in the ion torus now exist (Frank and Paterson,1999a, b), they have not changed these values substantially.The error bars simply represent the uncertainty in the shapeof the mass loading region around Io. Unless the Galileomeasurements occurred during a quiescent period of Io ac-tivity, and the density data suggest if anything the oppositeis true, then we come to either of two conclusions: the mass


Fig. 17. Averaged pro5les of plasma parameters from the Galileo-Io >yby.From top to bottom the panels show magnetic 5eld magnitude (Kivelsonet al., 1996), ion velocity and temperature from the plasma analyzerinstrument (Frank et al., 1996) total iogenic source density (Bagenal,1997) and SO+2 component pickup estimated by Huddleston et al. (1998b)and a comparison of the rms amplitudes of compressional and transversecomponents of B (Huddleston et al., 1999).

loading rate is smaller than generally believed, or the massloading region extends well beyond the near vicinity of Io.As we discuss in the next section, the latter may well betrue. If a 1000 kg s1 were accelerated to corotational ener-gies from orbital energies by the Io-torus interaction, then2:5 TW (2:51012 W) must be provided. This is equivalentto the power introduced for short periods into the Earthsmagnetosphere during the largest storms but at Jupiter thisenergy source is active continually. If the mass loading ratewere as low as 200 kg=s, then the energy supplied to thetorus is 0.5 TW. Because of its enormous cross-sectionthe mechanical energy >ux in the solar wind interceptedby the jovian magnetosphere is much larger than that inter-cepted by the Earths magnetosphere and it would dwarf theenergy >ux into the jovian magnetosphere at Io, if it couldbe tapped. Reconnection, that draws energy from the so-lar wind mechanical energy at Earth, appears to be weak atJupiter and in any event would not a7ect the rapidly rotat-ing plasma of the magnetosphere whose corotational elec-tric 5eld much exceeds that of the solar wind. Thus onlyover the polar cap and near its boundaries do we expect tosee activity powered by solar wind coupling. Fig. 17 also

Fig. 18. Trajectory of Galileo on its 5rst >yby of Io illustrating thephenomena observed and the inferred >ow pattern. Vectors along thetrajectory indicate >ow direction as seen by the plasma analyzer (afterRussell et al., 1997).

emphasizes that the main mass loading region is about 2 Iodiameters across. The amount of magnetic >ux convectedinto the 4RIo wide region is about 80; 000 Wb s

1. Thus,about 80; 000 Wb s1 accompany the newly added plasmauntil they reach a region in which they can be separated.Such a separation is needed in steady state because the mag-netic >ux in the magnetosphere is set by the jovian dynamoand changes at most slowly. In contrast ions need to escapefrom Jupiter to avoid reaching an in5nite density.A 5nal aspect, illustrated by Fig. 17, is the growth of

waves in the neighborhood of Io. On either side of Io thesewaves are principally ion cyclotron waves, while at theedges of the wake there are mirror mode waves (Kivelsonet al., 1996a, b; Russell et al., 1998b; Huddleston et al.,1999; Russell et al., 1999b). The appearance and amplitudeof the waves is consistent with the observed properties ofthe plasma and adds credence to interpretation of the natureand the strength of the interaction.Fig. 18 shows a schematic of the interaction of the torus

plasma with Io during the 5rst Galileo >yby illustratingthe probable >ow streamlines and where various phenom-ena were observed along the trajectory. One phenomenonnot relevant to the mass content of the torus but possiblycaused by the mass loading is the presence of electron beamsin the wake (Frank and Paterson, 1999a; Williams et al.,1999). These beams possibly mark >ux tubes with electricpotential drops in the ionosphere where the mass loaded


Fig. 19. Three dimensional plot of the reduced distribution function versusenergy per charge for spectral measurements made in Voyager 1s C cupbetween 0730 UT (7RJ) and 1145 UT (4:9RJ) on March 5, 1979. Theback panel shows the total positive-ion elementary-charge concentrationas a function of time from 5ts to the corresponding spectra (Belcher,1983).

>ux tube is being accelerated back to corotational veloc-ity. These beams are important for the electron heat balancebecause electron-electron heat transfer is more rapid thanion-electron heat transfer.

7.2. The Io torus

While Voyager did not pass as close to Io as did Galileo,it provided measurements of the Io torus that have proveninvaluable to our understanding of the energetics of theJovian engine. A good summary of our pre-Galileo under-standing can be found in the review by Spencer and Schnei-der (1996). Fig. 19 shows energy per charge readings onMarch 5, 1979 as Voyager 1 cut inbound through the Iotorus, prior to its outbound pass underneath Io. Before 0800UT the torus plasma had a low density, 900 cm3, andwas warm. As the radial distance of the Io orbit is reached,the plasma density rises and the plasma remains warm. Af-ter 1000 UT, when the spacecraft is well inside the Io orbita very cold population of plasma is seen with a composi-tion that indicates its iogenic origin. (The varying densityand charge state has been explained in terms of centrifugaland mirror forces by Bagenal (1985).) The reason for thelow temperature in this region is poorly understood. Eitherthe plasma is produced cold or it cools in the inner torus.The plasma in the wake directly behind Io (Frank and Pa-

terson, 1999b) that was picked up at very low velocities iscold but it is not clear how to transport this plasma into theinner torus. Neutral gas from Io will get picked up in theinner torus with a much smaller gyro velocity than in theouter torus. This also leads to a lower temperature. Finally,the residence time in the inner torus may be long enoughfor Coulomb collisions to cool the ions.Two mechanisms may be responsible for maintaining Ios

neutral atmosphere than in turn supplies the Io plasma torus.The 5rst is solar heating of the SO2 frost on Io. This wouldproduce an atmosphere centered on Io noon. The secondprocess is sputtering by the Io torus plasma. This would pro-duce an atmosphere centered on the upstream site of Io. Toexplain the appearance of the neutral sodium cloud, Wil-son and Schneider (1999) involve both these sources. How-ever, sodium is a minor constituent of the atmosphere andthe photochemistry of sulfurdioxide is di7erent than that ofsodium and we cannot assume that the neutral SO2 cloud isthe same as the Na cloud. One way to determine the extentof the SO2 neutral cloud is to determine where the ion cy-clotron waves arise when the neutral cloud becomes ionized.There is very little evidence for an extensive spherical ex-osphere at Io as there are essentially no waves upstream ofIo (Russell et al., 2001a). Rather the neutral cloud appearsto be a disk extending perhaps 0:5RJ inside and outside ofIos orbit and far downstream but not upstream.The radial extent of the neutral torus around Io is at 5rst

surprising because the atmosphere of Io is relatively cold.However, photoionization followed by acceleration of ionsin the outward electric 5eld associated with the corotation ofthe torus plasma followed by charge exchange can producefast neutrals close to Io where the neutral density is large(Wilson and Schneider, 1999; Wang et al., 2000). The newions travel in cycloidal paths in roughly the direction ofcorotation. If the local density of neutrals is su6ciently large,these ions can charge exchange to become a neutral cloudmainly outward from Jupiter but with an inward extension aswell. This allows the total mass-loading associated with Io tobe large even though close to Io the mass-loading is smallerthan expected. The Galileo data also allow a test of whetherthe sulphur dioxide atmosphere responsible for the toruscomes from a dayside Io atmosphere or a wake-sputteredatmosphere. If the former is the source the neutral cloudwould vary with the orbital phase of Io. In the latter case itwould not. Galileo data show a very strong variation withthe orbital phase of Io (Russell et al., 2001a).Turning now to the origin of the hot outer torus, we show

in Fig. 20 isodensity contours of the torus derived from theVoyager 1 data by Bagenal (1994). We have integrated thenumber density along the 5eld and in azimuth to give a torusdensity per jovian radius (RJ) in the radial direction. If weadopt a simple model of mass conservation with a one ton=ssource at 5:9RJ and radial expansion we obtain the velocitiesshown above the plot that is required to maintain this pro-5le. A source strength of 200 kg s1 would give numbersone-5fth of these values. At the higher mass loading rate,


Fig. 20. Io torus density as observed by Voyager 1 (Bagenal, 1994)Numbers at top of 5gure indicate the vertically and azimuthally integratedmass of the torus in radial bins of extent 1RJ and the estimated radial>ow velocity for iogenic mass loading of a ton=s.

the plasma would take about 3 months to reach a distanceof 7RJ where the density drops more rapidly and the radialvelocity presumably increases. This is very similar to thescale time for the rise in S+ and its decay after a Na outburststudied by Brown and Bouchez (1997). Thus it is possiblethat the 1 ton=s is achieved occasionally but is not a perma-nent rate. Long term torus studies indicate much variabilityof its emissions (Thomas, 1993).An alternative loss mechanism to radial transport is loss

down the magnetic 5eld lines into the ionosphere. This isdesirable because it removes the particle from the magnetic5eld lines. If signi5cant particle precipitation is occurringthen our estimate of radial velocities based on mass conser-vation are over estimated. Radial transport in a collisionlessplasma transports the magnetic 5eld too. While Jupiter needsto shed these ions to maintain a steady state, it does not needto shed magnetic >ux. At radial distances close to that ofIo, the torus is quiet but the >uctuation level increases withincreasing radial distance. Because the background 5eld isstrong, the fractional amplitude of these waves is small. Thetransverse and compressional wave amplitudes are less than0.1% of the background 5eld at frequencies near the ioncyclotron frequency throughout most of the torus (Russellet al., 2001b). Whether these waves can reduce the masscontent of the >ux tubes depends both on the size of the losscone and on the radial transport times. Estimates of the lossrate as the >ux tubes convect outward through the dipolarregion of the magnetosphere indicate that this is only a mi-nor loss mechanism (Russell et al., 2001b). This is in accordwith the observed low auroral activity in this region. In themagnetodisk region the loss cone is small and the plasma istrapped in the current sheet region.

Fig. 21. Linearly detrended time series near outer edge of torus at 7:7RJobtained on December 7, 1995 from 1521 to 1531 UT showing step-likechanges in the magnetic 5eld that might be produced by the interchangein stability (Russell et al., 1999a, b).

Radial transport is usually ascribed to the interchange in-stability in which a heavily laden >ux tube interchanges ra-dially with a lightly laden tube so that the heavily laden tubemoves outward. This preserves magnetic >ux and transportsions outward. This process has been treated by many authors(Southwood and Kivelson, 1987, 1989; Siscoe and Sum-mers, 1981; Pontius and Hill, 1989; Vasyliunas, 1989; Hill,1994). Observational evidence for >ux tube interchange canbe found in the Io torus. Kivelson et al. (1997b) have re-ported an empty inward moving >ux tube and the >ux tubesnear the outer edge of the Io torus, as shown in Fig. 21, ap-pear to be divided into tubes of di7erent plasma content asone might expect from the interchange instability (Russell etal., 2000b). If this process is not accompanied by scattering,tubes will retain their total content and be full until theyget emptied at some very large radius, e.g. by loss of theions down the tail. Then the interchanging tubes would con-sist of two types, full and empty. The variation seen in themagnetic 5eld in Fig. 21 suggests that there is a continuumof >ux tube contents at 7:7RJ. Plasma observations showthat this region at the outer edge of the torus has unusuallyintense 5eld-aligned electron beams with directional energy>uxes of up to 90 erg=cm2 s sr parallel and anti-parallel tothe 5eld, Frank and Paterson (1999a). If these >uxes extendall around the edge of the Io torus and extend over a 1RJband, then taking a >ux of 15 erg=cm2 we obtain a 0:3 GWenergy source if the beams each have a width of 0:5 sr. Thusthese beams represent a signi5cant heating source for theplasma. We note that the appearance of these beams coin-cides with a steep gradient in the plasma number density,i.e. the outer edge of the torus, reinforcing the interpreta-tion of these >uxes as being associated with the interchangeinstability. A second but possibly related phenomenon, isthe occurrence of transient energetic charged particle injec-tions deep in the middle magnetosphere (Mauk et al., 1999).


Fig. 22. Evidence for a plasma plume from Europa (Intriligator and Miller,1982). Top panel shows the >ux of oxygen and presumably sulfur ionsobserved by Pioneer as it moved outward past the orbit of Europa. Thebottom panel shows a sketch of the trajectory 5xed in a magnetic dipoleordered coordinate system.

While they envision an external trigger and possibly radialtransport inward, an internal source, possibly the interchangeinstability is also a candidate for these injections.

7.3. Transport in the middle magnetosphere

Our examination of the transport in the Io torus left us withtwo possible loss mechanisms for the ions, precipitation andtransport. We can help decide between the two mechanismswith the aid of the Europa plume (Intriligator and Miller,1982). Pioneer plasma observations shown in Fig. 22 (top)have been interpreted as a plume emitted by Europa andwrapping around Jupiter dragged by the corotational >owbut slowly moving out at about 400 m s1 as sketched in thebottom panel of Fig. 22. Disturbances seen in the Galileomagnetometer data (Russell et al., 1999a) also have beeninterpreted as due to a Europa wake. Fig. 23 shows a pairof these disturbances. If these were emitted by Europa onsuccessive rotations of Jupiter, then the radial separation ofthese two observations is equivalent to an outward speed ofthe plume of 500 m s1, quite consistent with the inferenceof Intriligator and Miller (1982). Unfortunately, if the twoevents detected were not produced on succession orbits ofEuropa then the speed would be reduced by some integral

Fig. 23. Magnetic 5eld observed by the Galileo spacecraft just outsideof the orbit of Europa showing wake-like disturbances that appear to beconvecting outward from Europa. The two disturbances seen at 0900 and1100 UT can be used to estimate the radial >ow velocity (Russell et al.,1999a, b).

factor. Nevertheless, this outward motion, no matter whatits absolute value, suggests to us that radial transport is op-erative. Nevertheless, we recall that for this transport to bee7ective we need to 5nd a mechanism that empties the >uxtube of ions and returns empty >ux tubes to the inner mag-netosphere.

7.4. The inner edge of the magnetodisk

The journey outward is not steady. Evidence of thetime-varying behavior of the magnetosphere comes 5rstabout 24RJ where the magnetic 5eld switches from nearlydipolar inside of this distance to decidedly non-dipolar out-side of this radial distance. Fig. 24 shows the change in theradial component of the magnetic 5eld between the regionsabove and below the current sheet and the component nor-mal to the current sheet (Russell et al., 1999c). The normalcomponent varies from pass to pass and thus the J Bforce in the plane varies. We can conjecture that the forcebalance in this region is principally between the magneticforces, the centrifugal force of the cold nearly co-rotatingplasma and the pressure gradient in the hot plasma thatcauses the depression around the current sheet. Two ofthese components can be estimated from the magnetometerand used to solve for the centrifugal force and if corotating,for the mass density. This is done in Fig. 25 for the 5rstfour inbound passes. These show some consistency in theradial pro5le of density from pass to pass but with occa-sional large decreases in the density. As before with thetorus plasma we can invoke conservation of mass to obtaina velocity pro5le. These points together with our earlierderived speeds are shown in Fig. 26. The left-most line isderived from the assumption that Io provides 1 ton=s. A


Fig. 24. Magnetic 5eld components (radial and normal to the current sheet)measured across and at the current sheet respectively. These componentsallow the radial J B force to be calculated. We note that while thereis little variation in the radial magnetic 5eld and hence the azimuthalcurrent, the normal component is variable (Russell et al., 1999a, b).

very consistent pattern of acceleration with distance is seenwith almost stagnant plasma in the inner torus, increasingin radial velocity to 500 m s1 at 10RJ to about 8000 m s1

at 20RJ to about 25 km s1 at 30RJ. If we instead use the

200 kg s1 value of Bagenal (1997) we displace this lineby a factor of 5 downward. This value is consistent withthat derived from the LECP data by Kane et al. (1995), butit becomes less consistent with the Europa plume estimatebased on the once-around assumption. The estimate basedon the Europa wake detection at radial distances beyond theEuropa orbit are higher than these two curves. This suggeststhat the Europa wake sightings occur during periods ofhigher than usual outward transport. This is consistent withtheir rare occurrence. We emphasize that these two curvesbracket our inferred radial pro5le of the average outwardvelocity under the assumption of no particle precipitation.If particles are lost along the 5eld line so that part of the

Fig. 25. The mass density of a ring of the magnetodisk current sheet ofradial extent, 1RJ , estimated from the radial force balance (Russell et al.,1999a, b).

Fig. 26. The radial variation of the out>ow velocity estimated fromconservation of mass and Voyager observations, the two independentobservations of the Europa plume by Pioneer 10, in December 1973 andGalileo in September 1996, by the stress balance in the jovian current sheetas deduced from Galileo magnetic observations and from the VoyagerLECP anisotropy measurements (Russell et al., 1999a, b).


Fig. 27. Four crossings of the magnetodisk current sheet in current sheet ordered coordinates illustrating the multiple crossings of the current sheetassociated with surface waves of about 10 min period as well as the evolving structure of the normal component to the current sheet (Russell et al.,1999a, b).

density decrease is due to particle loss along the 5eld, therequired radial out>ow speed would be further reduced.Another means to estimate the radial velocity is to use

magnetic >ux conservation. We do this despite the fact thatas much magnetic >ux must be convected inward as is beingconvected outward. The third curve in Fig. 26 shows thevelocity required to give the radial fall-o7 in the magnetic5eld crossing the equatorial plane if the plasma is convectingthe magnetic >ux outward. For this calculation we assumethat there is slow out>ow over nearly 360 with a narrowregion of rapid in>ow. If the >ux were to return at the samespeed over half the circumference, then the speed would betwice that shown. As we discuss later we believe that themagnetic >ux does return rapidly over a narrow region, atleast in the middle magnetosphere and torus, so that ourestimate is approximately correct.There is a di7erent slope in the inner magnetosphere and

torus than in the magnetodisk. This could be due to the

cessation of particle losses along the 5eld once themagnetodisk is encountered. We further note that themagnetic-5eld based calculation is about a factor of twoless than the lower of the two density-based calculations.While both calculations were made assuming 360 sym-metry and so both su7er from lack of knowledge of thetrue >ow pattern, the magnetic estimates especially at highradial distances push the accuracy of the magnetic measure-ments because the normal 5elds are very small, approaching0:1 nT. Thus the lower two curves should be viewed asmutually consistent. Based on the 200 kg s1 curve it takesabout 2.5 years for the torus plasma to reach Europa fromIo and about 0.5 years from Europa to Ganymede. FromGanymede to Callisto takes about one month and then fromthere to 50RJ takes about a week.As our conjectured >ow moves outward through the mag-

netodisk, the current sheet becomes less stable. As shownin Fig. 27 the magnetic 5eld crossing the current sheet


(component Bb) is moderately strong near 25RJ averagingabout 5 nT in the center of the sheet with some high fre-quency noise. At about 40RJ the normal component has di-minished to about 1 nT with >uctuations that may in partbe due to surface waves that cause changes in the currentsheet orientation. These waves carry the current back andforth over the spacecraft. At 50RJ the motion of the currentsheet is even greater and the possible e7ect on the apparentnormal component across the average current sheet is evengreater but the reversals seen in the normal component maynot be real. Their correlation with the >uctuation in currentsheet location suggests that the current sheet orientation isrocking as the surface waves move by. Finally, at 55RJ wesee pulses of reversed 5eld components at the current sheetthat cannot be explained by current sheet motion. These nar-row pulses of reversed 5eld seem to be small tearing islandsembedded within the current sheet. We suggest that thesesmall tearing islands may act as the seeds for reconnectionwhen this region rotates to the nightside but that while it ison the dayside the tearing islands remain quiescent withinthe current sheet.The multiple crossings of the current sheet, evident on

three of the four current sheet encounters displayed here,are quite fascinating as they indicate a magnetodisk that isconstantly in motion, agitated by some source, possibly deepinside the magnetosphere and possibly due to the solar windinteraction.

7.5. Separating the ions from the magnetic 7eld

When the Galileo spacecraft reached the distant midnightmagnetosphere, it was able to detect the process separat-ing the ions from the magnetic 5eld. Fig. 28 shows thetrajectory of the Galileo spacecraft in May and June 1997when it had a long sequence of continuous data through thepost midnight sector. Fig. 29 shows the magnetic measure-ments from June 1 to June 21 when Galileo moved fromapogee to about 50RJ at a local time of about 4 a.m. Thissection of data reveals many features of the dynamics ofthe jovian magnetosphere. The total 5eld shows the evi-dence of crossing the current sheet in its periodic decreasesbut it has longer modulations as well. These variationson the period of days could well re>ect changes in thesolar wind dynamic pressure as the solar wind is seen tovary on such time scales. The current sheet also appearsto move mainly above and below the spacecraft for daysat a time. This also appears to be due to the solar wind,in this case being due to variations in the direction of thesolar wind >ow. Finally, there are variations in the compo-nent of the 5eld across the current sheet. These variationswhile brief are the most dramatic of all. They rise rapidlyto 5eld strengths even greater than the originally existing5eld outside the current sheet. We interpret these aslarge-scale reconnection events (Russell et al., 1998a,2000c).

Fig. 28. The Galileo orbit used to study the substorm-like phenomenapresented in Figs. 29 and 30 (Russell et al., 1998a, b).

Fig. 29. The magnetic 5eld observed by Galileo on orbit G8 from apogeein to about 50RJ (Russell et al., 1999a, b).

Fig. 30 shows a very strong event in which the 5eldstrength increases nearly a factor of four when the 5eldturns southward. We interpret such a strong resultant 5eld asdue to explosive reconnection when the reconnection pointreaches a region of very low density out of the current sheet.Here the Alfven velocity is high and the plasma can be ac-celerated to high speeds and reconnection can proceed veryrapidly. This process creates an X-point in the magnetic 5eldin a swept-back meridian phase as shown in Fig. 31. Theselargest events contain about 1010 Wb or about 10 times themagnetic >ux in one of the tail lobes of the Earth but only


Fig. 30. An example of explosive reconnection as detected by Galileo.Magnetic 5eld components are directed outward, southward and in thecorotation direction.

Fig. 31. Interpretation of the measurements shown in Fig. 29. At pointA the magnetic 5eld becomes strongly southward and strong as if themagnetic 5eld had suddenly reconnected outside of the radial distanceof Galileo. Events with northward turnings are also seen (Russell et al.,1998a, b).

about 0.1% of the >ux in the jovian magnetotail. At the rateof events seen in Fig. 28, once every 4 hours, an average ofabout 70; 000 Wb per second can be emptied of ions and beprepared for return to the inner magnetosphere. It is reas-suring to note that the amount of magnetic >ux that is massloaded at Io every second, assuming that Ios mass loadingcan be approximated by a step function four Io radii across(in the radial direction perpendicular to the >ow), is about80; 000 Wb per second. Thus the substorm process and themass loading process are in approximate balance.Putting these observations in a more global context we

show in Fig. 32 the Vasyliunas (1989) prediction for the cir-culation in the jovian magnetosphere. Like Dungeys mod-els of the reconnection process at Earth, this model powersthe magnetosphere in a steady-state manner. The fact thatwe 5nd a temporally varying magnetosphere at Earth or at

Jupiter does not negate the model. The observed circula-tion is in fact quite consistent with the qualitative pictureof Vasyliunas. The Galileo, Voyager and Pioneer observa-tions used in this review have merely made this model morequantitative. We have not yet completely solved our earlierposed dilemma, because we have not returned the emptied>ux tubes into the inner magnetosphere. We examine thatprocess next.

7.6. Returning the 8ux to the inner magnetosphere

The above observations strongly support a model of thecontinued, inexorable, outward >ow of plasma, movingslowly at 5rst but accelerating rapidly until a radial veloc-ity of close to 50 km=s is reached about 40RJ. The radialplasma >ow carries magnetic >ux with it because we cannotseparate the plasma from the magnetic 5eld until recon-nection takes place beyond about 50RJ in the midnight to3 a.m. sector. Reconnection that reaches the lobes aboveand below the current sheet, as these powerful reconnectionevents do, create nearly empty magnetic >ux tubes that arelighter and have a much lower beta than the mass-loaded>ux tubes. In the rapidly rotating jovian magnetospherelight >ux tubes are buoyant and should be able to >oat intoward Io. If so then where are they? In fact we do seesuch depleted >ux tubes. One such tube was discussed byKivelson et al. (1997b). Others, but smaller, are seen in theIo torus (Russell et al., 2000e) as shown in Fig. 33. Theydistinguish themselves by their increased 5eld strength. Theincrease in magnetic energy density matches that expectedfor a tube drained of plasma at this location in the magne-tosphere. Such high-resolution data are rare in the Galileorecords and thus we cannot provide an extensive survey ofthe occurrence of such tubes but we have covered the Iotorus quite thoroughly (Russell et al., 2000e). One mightwonder at the paucity of such depleted >ux tubes in thelimited records on hand as they occur only about 0.4% ofthe time. Naively one might expect that the depleted re-turning >ux tubes would occupy half of the magnetosphere.However, since they are not transporting plasma, the >uxtubes can move rapidly. Thus if the depleted >ux tubes inthe Io torus move at about 200 times the outward velocityof the outward laden >ux tubes then there is >ux balance.If the outward velocity here is about 10 m s1 then theinward velocity might be 2 km s1. In fact Kivelson etal. estimated an inward velocity much greater than this,over 100 km s1.Depleted >ux tubes have also been observed between Eu-

ropa and Ganymede (Russell et al., 1999a). These tubes lastlonger than the tubes in the Io torus. Perhaps they are mov-ing more slowly but because the outward >ow is more rapidhere we would expect the empty tubes to occupy a largerfraction of the records unless they also moved inwards morerapidly in this region. It would be desirable but is very dif-5cult to get an inventory of the amount of inward moving


Fig. 32. The circulation pattern and magnetic topology in Vasyliunas (1989) reconnecting jovian magnetosphere.

Fig. 33. Empty >ux tubes in the Io torus. An average magnetic 5eldmagnitude has been removed in each panel leaving only the change in5eld magnitude. The 5eld strength during these periods is about 1500 nT.

>ux. In fact there may be much >ux moving inward withscale sizes too small for Galileos magnetometer to detectbecause of its generally poor ( 20 s) resolution.

8. Saturn, Uranus and Neptune

Our entire body of knowledge of the magnetospheres ofSaturn, Uranus, and Neptune rests on the information gainedfrom the >ybys of Pioneer 11, Voyager 1, and Voyager 2for Saturn, and Voyager 2 only for Uranus and Neptune.Expressed in terms of planetary radii the magnetospheres ofthese three planets are similar in size although in absoluteterms the magnetospheres are over a factor of two smaller.

Because as we stated earlier, the Mach number of the shocksstanding in front of these magnetospheres are very strong,the plasma behind the shocks is very hot. This produces ahigh beta plasma in which magnetic pressure is low andthe dynamical processes in the magnetosheath are domi-nated by pressure gradients in the plasma and not magneticforces. Thus we do not expect reconnection to be impor-tant at these magnetospheres. We do, however, see evidencefor some reconnection in terms of magnetic 5elds alongthe normal to the magnetopause (Huddleston et al., 1998a).We also note that as the magnetic 5eld of the solar windweakens with distance from the sun with a constant solarwind velocity, the electric 5eld of the solar wind becomeseven weaker with respect to the corotational electric 5eldof each of these planets that have moderately rapid rotationand moderately strong intrinsic magnetic 5elds. Neverthe-less, it is possible that the polar regions and their extensionsthe magnetotails are linked to the solar wind though recon-nection, and exhibit some solar wind associated phenomena.In fact such a dynamic event has been reported for Uranus(Kane et al., 1991). It is unlikely that we will learn moreabout the magnetospheres of Uranus and Neptune for quitesome time but we should soon have data from the Cassinimission in orbit about Saturn beginning in July 2004 andextending for four years. This mission includes a full com-plement of instruments to study the magnetosphere and itsinteraction with its moons. At this writing the spacecraft andits payload are well on their way to Saturn and performing>awlessly.

9. Conclusions

Dynamic pressure and the strength of planetary magneticmoments principally control the size of a planetary magne-tosphere. A secondary factor is the presence of an internalmass or energy source such as provided by Io deep in the


jovian magnetosphere. No other planet has a strong sourceof plasma so deep in the magnetosphere. Reconnection playsa signi5cant role in the dynamics of most magnetosphereswith intrinsic magnetic 5elds and possibly for some inducedmagnetospheres, most notably for comets. Field-aligned cur-rents are the pathways by which stresses are communicatedfrom one part of a planetary magnetosphere to another, suchas from the outer magnetosphere to the ionosphere on Earth,or from Io to the jovian ionosphere. These currents can leadto electric potential drops along 5eld lines. These potentialdrops can accelerate charged particles into the atmosphereto excite aurora and to decouple the magnetosphere fromthe ionosphere

The Dynamics of Planetary Magnetospheres

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