Solar wind measurements near Mars and their implication in the Red Planet environment

Pltrrwt. .spw

sci.. Vol. ‘Id. No. 2. pp. I I7 117. IYYh Copvright I 19% Else\ ier Science Ltd Printed in Great Britain. All rights rcservcd 0032 aii Y6 $1 .5.00 + 0.00

Pergamon

0032-0633(95)00071-2

Solar wind measurements near Mars and their implication in the Red Planet environment .J. G. Trotignon,’

R. Grard,’ S. Barabash,j R. Lundin’ and E. Dubinin’

’ Laboratoirc dc Ph>Gque et Chimie de I’Environnement. CNRS. 3A avenue de la Recherche Scientifiquc. I ‘-1_5071 OrlPan\ cede\ 01. France ’ Space Science Department. ESA ESTEC. Keplerlaan 1. Postbus 29Y. NL-7200 AG Noordwijk. The Netherland\ ’ Swedish Institute of Space Physics. Box 812, S-981 28 Kirunu. Sweden ’ Space Rcsrarch Institute. Profsoyrnayt X&37. 117810 Mosco~c. Russia Received 7 Octc~bcr I YYl : revised

IO March 1995 : accepted .S 4pril l9Y 5

The Phobos 2 spacecraft was operational around Mars for almost two months. From 29 January to 27 March 1989, the Martian bow shock was crossed at least 200 times and the upstream regions extensively probed. The solar wind density and velocity deduced from the a-c. electric-field measurements performed with the PWS plasma and wave experiment are compared with those of the ASPERA 3-D plasma composition experiment. A reasonably good agreement is found between the 12 min averages of the plasma density numbers derived from PWS and ASPERA (considering the poor precision of the instruments) : they differ by at most a factor of four. The solar wind density is found to decrease when the flow speed increases, in agreement with the stream structure of the solar wind, which results from heliomagnetic latitude dependencies of solar wind parameters. The solar wind ram pressure, which is essential to study the solar wind-Mars interaction, is also estimated : the absence of a clear relationship between this dynamic pressure and the bow shock location is confirmed. The bow shock observations performed during very high solar activity do not support the hypothesis that Mars possesses a significant intrinsic magnetic field. Abstract.

1. Introduction

On ?Y Janu:iry IYXY, six months after its launch from Baikonour (Kazakhstan). the Phohos -7 spacecraft was

injected into an equatorial Martian orbit with a period of 76.5 h, and periapsis and apoapsis altitudes of 850 and 80.000 km. respectively (Sagdeev and Zakharov. 198’)).

On IX February 1989, when the orbit W;LScircularized in preparation for the close flyby of Phobos, the natural satellite of Mars, Phobo~s -3 had crossed the Martian bow shock eight times (four times inbound near the subsolar point and four times outbound in the tail). Until 27 March. when the mission ended prematurely. the spacecraft described a circular orbit in the equatorial plane of the planet with a radius of 9670 km and a period of X h. Thus, during two months Phohos _7crossed the Martian shock about 200 times and performed extensive observations of the upstream solar wind. This paper presents the solar wind parameters, density. velocity and ram pressure, derived from two of the instruments that form the Phohos 2 scientitic payload : the threedimensional plasma composition experiment, ASPERA. and the plasma and wave analyser, PWS. We also make limited comparisons between the PWS;ASPERA measurements and already published data collected with TAUS, a proton. alpha particle and heavy ion spcctrometer based on hemispherical clectrostatlc analysers with magnetic deflection systems (Rosenbauer C/ trl., 1989). Particle and field data are evidently crucial to study the planetary or cometary bow shocks. as well as the equilibrium between the pressure (thermal and/or magnetic) exerted by the plasma of planetary or cometary origin and the solar wind ram pressure. In particular. the knowledge of the latter is essential to set an upper limit to the magnitude of planetary intrinsic magnetic tields. This presentation is divided into three main topics. Brief descriptions of the ASPERA and PWS instruments arc given in Section 2. Section 3 is mainly devoted to a comparison between the solar wind parameters obtained from the two experiments. Finally, in Section 4. we study the relationship between the properties of the solar wind and the Martian bow shock position.

J. G. Trotignon

ASPEfXA Field of View (Scanning / Spinning)

ASPERA Field of View (non-scanning / Non-spinning)

-90 -180

1’1(I/. : Solar wind measurements

180 sczLimuthOAn,,,~

-180

40

0

Solar AzimuthAngle

90

180

(deg)

tield of view in the tixed setting mode (left) and the scanning/spinning mode (right). For fixed setting. the 5 x 360 field of view lies either perpendicular or parallel with the ecliptic plane. .I.}.. A typical solar wind offset of IO is considered for illustration (Lundin ct (I/.. 1993) Fig. 1. ASPERA

2. Description of the two instruments

1.2. Tlw pltrsrmr II’L/~‘P.sj,strm. P F’S

The ASPERA experiment was a joint Swedish. Soviet, and Finnish enterprise supported by grants from the Swedish National Space Board and the Academy of Finland. Significant contribution and technical support were provided by the staff of the Space Research Institute of the Soviet Academy of Sciences. The instrument has been designed to measure three-dimensional electron and ion distributions with the help of two spectrometer systems. Ions are analysed within the energy range 1 eV e-‘-74 keV e-’ and electrons from 1 eV to 50 keV. As the combined field of view of the detectors. 5 x 360 . lies in planes parallel to the Phohos 2 spin axis. the 471 space can be covered in half a spin period of the spacecraft (4-6 min). On the circular orbits. only half of the 10 ion sensors were used. each of them covering an angular sector of 36 , so that a full spin period was necessary to obtain a full coverage. When the spacecraft was three-axis stabilized (elliptical orbits 3 and 4, elliptical to circular transition orbit, and some of the circular orbits), a full coverage could still be achieved in about 2 min (I 80 turn) by activating a scan platform, whose rotating axis was aligned along the nominal spacecraft spin axis. When the scan platform was not working. a two-dimensional cut of the distribution function could still be obtained every 13 s at best and 2 min at worst. Figure I shows the ASPERA field of view as a function of the solar elevation and azimuth angles. assuming a 10 offset of the solar wind with respect to the .\--axis of the Mars solar orbital coordinates (Lundin rt cl/., 1993). For fixed setting, the cuts in the .v_~‘andXT planes are mutually exclusive and depends on the ASPERA mode of operation. The first three moments of the distribution function. density, mean flow velocity and pressure tensor. were calculated on board in order to reduce the amount of transmitted data. For a more detailed description of ASPERA. see Lundin et N/. ( 1989).

The PWS instrument was developed by the Space Science Department (ESA’ESTEC) and the Laboratoire de Physique et Chimie de I’Environnement (CNRS. France) with the collaboration of the Space Research Institute of the U.S.S.R. Academy of Sciences (IKI, U.S.S.R.), the Institute of Geophysics and Planetary Physics (UCLA, U.S.A.). and the Space Electronics Laboratory of the Polish Aviation Institute (IL. Poland). The main function of the plasma wave system is to measure the electric component of natural waves at frequencies up to 150 kHz. The ambient electric field is derived from the difference in potential detected between two solid spheres. IO cm in diameter, separated by a distance of 1.45 m. The spectrum of the signal delivered by the dipole antenna is analysed with a set of 24 adjacent and logarithmically spaced filters, which cover the frequency range from 5 Hz to 150 kHz. Some low-frequency channels are also available. As the 24 filter channels are not sampled at the same time and the signal level fluctuates over a short time scale (less than 7 s) in the Martian plasma environment, data processing has been applied on the ground to each spectrum to take these effects into account. First, a cubic spline interpolation process has been applied to each of the four sets of six channels that are commutated at the same instant. Second. the four spectra thus obtained were averaged. This data processing is fully explained in Trotignon 01 trl. (1991a,b). An interpolation process is also included so that smoothed spectra made of 72 artificial frequency channels are produced. It is worth noting that this interpolation is essential to better estimate the plasma density as explained in Section 3. I. PWS has two other functions : to continuously monitor the voltage between one of the spherical probes and the spacecraft structure and to perform plasma diagnostic5 using a Langmuir probe. Several temporal resolutions were used during the Phohos 2 mission. For example. an electric field spectrum was returned in 3, 20 or 80 s depending on the PWS mode of operation. The location of the PWS elements on the Phoho.~ 2 spacecraft is dis-

I. G. Trotignon

cm/u/.

: Solar uind measurements

I I’)

X

t to

the

sun

PHOBOS SPACECRAFT

*

---

\ solar panels

2.2m

electric

$ z

unit

Fig. 2. A sketch 01’the P/whos spacecruft showing the location <)I‘the two electric \ensors. the Langmuir probe and the plasma \\,~\e system electtx>nic unit (Trotignon c’t ol.. 199la)

played in Fig. 2 (Trotignon technical information about < irard c’t ol. ( I YXY).

CI al.. 1991a) ; additional the instrument is given in

3. The solar wind density and velocity derived from the ASPERA and PWS observations

A typical example of electron plasma oscillations detected by PWS in the upstream region of the Martian shock is shown in the top panel of Fig. 3. The electric-field intensit! scale (shown on the right) is logarithmic. with a dynamic r\mge of 55 dH, extending from about 1.3 /IV rn-’ Hz- ’ ’ to 0.7 mV m ’ H/ I’. P/who.c .? was entering its third elliptical orbit around Mars: it left the solar wind and crossed the planetary bow shock at 0535:50 UT. Sporadic electric field enhancements are seen in the frequency range from IO to I5 kHz. These electrostatic waves are observed when the spacecraft lies in the electron foreshock and is magnetically connected to the bow shock. that is to say \vhen the interplanetary magnetic field line passing through the spacecraft position intersects the shock front. The waves are known to be excited by a plasma instability associated with suprathermal electrons. which are rcflectcd and accelerated at the shock. These particles

subsequently backstream into the solar wind away from the shock along the interplanetary magnetic field line connected to the shock (Filbert and Kellogg. 1979). The upstream waves which develop in the environment of Mars have been reported and studied by Skalsk!, or ol. (1993). Associated enhancements of electron flus have been simultaneously detected in the energy range from 100 to 530 eV (Skalsky it r/l.. 1993). A correlation between backstreaming electron fluxes and electron plasma oscillations detected by ASPERA and PWS. respectiveI>. XI:, also reported by Barabash CT trl. ( 1993). The peak I‘rcquency of the electron plasma oscillation5 is knoti n to hc close to the ambient solar wind plasma frcquenc!’ (LAcombe (11(I/. ( 1985) and references therein). bvhich enables us to estimate the solar wind density (Trotignon c’t I//.. 1992). The middle panel of Fig. i shows with dols the inferred plasma frequency (right-hand scale) M ith ~hc associated plasma density (left-hand scale). The mean value of this solar wind density ih in ;I good agreement with the ion density measurements of the TAlJS cxpcriment which are represented by ;I solid lint (Verigin o/ L/I.. 1991 ). The bottom panel sho\vs the proton denhit!; delivered by the ASPERA instrument. The dispersion ofthe PWS results i, partly due to errors inherent to the density determination procedure. as WC‘ assume that we are observing normal Langmuir wave5 in ;I thermal plasma. Lacombe vt ~1. ( IYXS) hake i;ho\vn that. in the presence of ;I backstreaming electron p~~pulalion. :I ne\\. branch ot‘ the dispersion rcla~ion appears. ;I beam mode. which an be unstable. This tic\\ branch exlst5 a~ frequencies which can be either will below\ or well :tbo\,c the normal local Langmuir c\;~vc plama I‘lcquency. WC know from expectations at the Earth that the frequency offset I‘rom the electron plasma frequency does no1 actsally exceed 10%. Moreover. thi’r large otTset occurs at ;I small connection depth : then it decreases with distance from the tangent interplanetary magirtic lield lint. The dispersion of the PWS density estimations is also causal by the impulsive nature of the waves. Indeed. the electric field intensities appeared to fuctuate considerably o\ct time scales less than halfa second. the shortest time needed by PWS to deliver six ximultaneou5 freqiiunc~-channel output5 : ;ih ;I full spectrum is compo5cd of 74 channel\. it is returned in 2 s at best. 1%hich is too long to avoid temporal variation of the spectrum amplitude. .A modulation is ;tl\o observed in the ASPERA tncasuremcnts whenever the spacecraft is spinning and the scan platform not used : thih results from the narro\< acceptance angle of the detectors in one dimension (5 ) and the dc\iation of the spacecraft spin axis from the solar wind direction

(Kallio C/ t/l.. I YY3: Lundin (It r/i.. I YY7I. In lhc X0 min period shown in Fig. 3. no modulation is observed because P/roh,~~ 3 was operated in ;I three-auis c~rientalic~n mode. To remove these fluctuations. the PWS and ASPERA data have been averaged over intcr\ala 01‘ I2 min. which correspond approximately to the spin period. Figure 4 displays in blue and in red. respectively. 1I! min a\‘eragea of the solar wind density deri\,ed from the PWS and ASPERA data. between I February and ?6 March Ic)XY. The average \xlue is shown by ;I dot and the height of a bar i4 equal to twice the standard devi;ltlc~n. A running

120 average filter has been applied, instead of a simple averaging process, so dots appear every 4 min (and not 12) when data are available. Even if some discrepancies appear here and there, a good agreement is generally found between the PWS and ASPERA solar wind estimations when both data are available simultaneously. In order to quantify the compatibility of the two sets of measurements, we have plotted in Fig. 5 the ASPERA density versus the PWS density. As can be seen, most of the density estimations are in a ratio less than two and greater than a half: these limits being indicated by the two dashed lines. Moreover, four and a quarter appear to be maximum and minimum values, respectively, for this ratio. The latter result may be considered as satisfying as far as the ASPERA instrument was not originally designed to perform accurate measurements in the solar wind. As an example, the moment computation can be spoilt by a bad estimation of the spin period as well as a mixing of the signals generated in the ion foreshock by the inward solar wind protons and the outward ions reflected by the bow shock. The large density enhancements recorded near Mars on 9-10 March 1989, are the signatures of coronal mass ejections. These coronal mass ejections were observed on Earth and gave rise to auroras and a strong magnetic storm (McKenna-Lawlor et rzl., 1991 ; Dolginov and Zhuzgov, 199 1).

3.2. Solur n-indrelocit?. and r{~mrnic pressure Another important parameter of the solar wind is the flow speed. Figure 6 presents 12 min averages of the solar wind proton velocity (top panel) derived from ASPERA. As in Fig. 4, the vertical bars show the standard deviations. As can be seen, the average solar wind velocity near Mars varied between less than 200 km SC’ and up to 1200 km S _I, during the period from 1 February to 26 March 1989. ASPERA provides only two-dimensional cuts of the distribution function when the spacecraft is three-axis stabilized. This bad sampling of the distribution function explains why the mean flow velocity moment computed on board differs from that of the solar wind during the three-axis orientation mode. The velocity component along the Sun-Mars line can be restored, however. directly from the maximum of the proton energy spectrum. Note that 22% of the data plotted in Fig. 6 are derived from the maximum of the proton energy spectrum. On the other hand, a 47~ coverage is achieved when the spacecraft is spinning, which makes possible a three-dimensional flow velocity determination (Kallio et al., 1993. 1994). Knowing the solar wind density and velocity, it is now possible to evaluate the ram pressure. This is obtained by combining the PWS density and the ASPERA velocity determinations. The PWS density estimation has been chosen because it is obtained from a natural wave experiment and its accuracy is believed to be better than 30% (Trotignon et NI., 1992). Natural wave experiments are not subject to disturbances coming from the floating potential of the spacecraft and are known to give reliable and accurate determination of the electron density. The

J. G. Trotignon et ul. : Solar wind measurements 12 min averages of the solar wind dynamic pressure and their associated standard deviations are displayed in the bottom panel of Fig. 6. A similar profile, which covers a short time interval, between 28 February and 27 March 1989. was published by Schwingenschuh ct trl. (1992). Their results are from the TAUS measurements and are plotted in the bottom panel of Fig. 7. while the PWS; ASPERA curve is shown at the top. Both sets of results show a reasonable qualitative agreement but a closer comparison is not possible due to the fact that the times of samplings differ and the TAUS pressure is measured in arbitrary units. On 11 March 1989, at 2050 UT. the solar wind ram pressure reaches 23.5 x IO ’ dyn cm ‘: it is the sole point out of the bottom frame of Fig. 6 and the top frame of Fig. 7. Figure 8 shows three electric field spectra recorded during three different inbound crossings of the Martian bow shock. The upper frequency component, beyond 100 Hz, is the signature of Doppler-shifted ion-acoustic waves (Trotignon ct ul.. 1991a.b; Grard et rd.. 1993). These waves exhibit an upper frequency cutoff‘ less than C;,&&, where C;, is the solar wind speed and i., the Debye length (Gurnett. 1985). It appears in Fig. 8 that the frequency shift of the ion-acoustic waves varies with the solar wind dynamic pressure ; the maximum Doppler-shift frequency is indeed proportional to (P,,/T,)’ ‘, where Psw is the ram pressure and T, the electron temperature. Unfortunately. the Doppler shift does not yield the solar wind speed. because the angle made by the wave vector and the flow speed is unknown (the Doppler-shift frequency is given by the scalar product of the wave vector by b’sw).

4. Relationships between the solar wind density and velocity and the Martian how shock location Figure 9 shows the scatter plot of the plasma density versus the flow speed measured in the solar wind at 1.5 AU by PWS and ASPERA, respectively. An inverse relationship similar to that found in the interplanetary medium at 1 AU is clearly identified (Hundhausen et ul.. 1970). It is worth noting that this pattern is relative to quiet solar wind conditions. because the data corresponding to the solar flare events have been discarded. The variation of the solar wind density as a function of the flow speed finds its explanation in stream structures and heliomagnetic latitude dependencies first proposed by AlfvCn (1977). The separation between positive and negative solar magnetic field lines gives rise to the heliomagnetic current sheet. The plasma density reaches its maximum near the current sheet surface and decreases away from it, whereas the behaviour of the flow speed is exactly the opposite. As the apparent solar magnetic dipole axis is tilted with respect to the rotation axis. the distance from the spacecraft to the current sheet varies continually. allowing thus the observation of the solar wind velocity and speed variations. In the absence of strong solar perturbations, the mean variations of the solar wind density and speed may be attributed to the stream structure in the interplanetary medium.

et al. : Solar wind measurements

J. G. Trotignon

13

ELECTRIC FIELD SPECTROGRAM

0.1 4

,

0400

0410

0420

0440 0450 0430 Universal Time

Density

0510

0520 I\/

FEB 08, 1989

PHOBOS 2 Electron

0500

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Frequency

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0520

Time

Fig. 3. (Top) Electric field spectrogram recorded in the electron foreshock of Mars by PWS on board Phobos 2. Intense electron plasma oscillations are seen between 10 and 15 kHz. (Middle) The dots give the solar wind density (left-hand scale) inferred from the plasma frequency at which the electron plasma oscillation reaches its maximum (right-hand scale). For comparison. the density measured by the TAUS spectrometer (Verigin et al., 1991) is plotted as a solid line. (Bottom) Proton density derived from the ASPERA data

J. G. Trotignon et al. : Solar wind measurements

22

Proton Density from ASPERA EJec:tron 1)crlsit.y from PWS 20.0

I

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0225

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0313

(at 0000 1989

0321

0329

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12 min averages of the solar wind density deduced from the PWS electron plasma oscillation observations, from 1 February to 26 March 1989. The proton density measured by ASPERA has also been averaged over 12 min intervals and is plotted in red. A bar that joins the average value minus the standard deviation and the average value plus the standard deviation is associated to each dot. Dots appear every 4 mitt, instead of 12, because a running average filter has been applied

.I (3. Trotignon c’/~1.: Solar wind measurements ROj

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01‘3 half and tmo (and one). respectively

The existence of an intrinsic magnetic field at Mars is still an open question (Luhmann, 1992; Schwingenschuh c’t ~1.. 1992). Assuming that Mars possesses a weak intrinSIC dipole tield. the gas-dynamic model of Spreiter c’t ~11. ( 1966) may be applied to give an estimation of its magnetic moment. This exercise has already been performed by looking at the subsolar stand-off distances of the Martian shock crossed by Plwhos -7 during its first three elliptical orbits (Slavin c’t a/., 199 1 ; Russell et ul., 1992 ; Trotignon ct L/I., 1993). In Trotignon ct (I/. (1993) an upper limit of 2.2 x IO“ T m’ has been obtained by artificially situating the obstacle clohe to the bow shock, the true position of the obstacle being unknown. As stated by Spreiter ct t/l. (1966). the estimate of the magnetic moment depends upon the solar hind dynamic pressure and the position of t tie effecti\,c obstacle : if,,

= 1.927 x lO‘(P,,D:,“)’

2

(1)

\vhere the magnetic moment 121,. the solar wind ram pressure P,bv. and the subsolar distance D,, from the centre of Mars to the obstacle are expressed in T m3. dyn cm ‘. ,md km. respectively. To illustrate this result. we have plotted in Fig. IO the magnetic moment vs. the obstacle altitude for three dynamic pressure values. The solid lint corrcxponds to 0.83 x loo-’ dyn cmp2, a mean \aluc of the upstream dynamic pressure observed by l’iwhos 2 during its revolutions around Mars. The vertical ;Irrow points out the lowest shock altitude observed dur~ng the Phho.v 3 mission and the two horizontal arrows Indicate the magnetic moments estimated by Dolginov ;md Zhuzgob (1991) and Grafe (1992). From Fig. 10, we can see that the moment claimed by Dolginov and Zhuzgov (1991). say I.22 x IO” T m3, implies a subsolar obstacle altitude of 720 km. and the 0.7 x IO” T m3 value derived by Grafe (1992) gives a very low altitude of 30 km. Moreover, gas dynamic modelling of the solar wind-

Fig. 6. Twelve min averages of the solar ~c~nd proton \elocitk (top panel) and dynamic pressure (bottom panel) measured nwr Mars during the f%oho.~ 2 mission. between I k’ebruary and 26 March 1989. The proton Lelocity comes I‘rom rhe iZSPERA plasma composition experiment and the solar wind denxity th;tt is used to compute the ram preswre is deduced from the PWS wave data. The vertical bars gi\e cslim:ttion\ of the slandard deviation

Mars interaction predicts an obstacle height of the order of 500 km (Slavin ct rrl.. 1991). It is clear that the diversity ofmagnetic moment estimations published to date reflects the lack of a consensus within the scientific community about the mere existence of this field. In any case. the Phohos 2 observations have proved that such ;I weak intrinsic field, if it ever exists. should not play any significant role in the solar wind-Mars interaction, during periods of high solar activity (Breus (‘t (I/.. 1989). A mean Martian shock model has been determined b) least-square titting a conic section t0 the shock cro>\sing locations identified in the PWS data (Trotignon c’t trl.. 1991~. 1993). This model is a function of two main parameters, namely the eccentricity 8: and the semi-talus rccturn L. The latter is nothing else than the distance I’rom the focus of the conic to the shock front. measured perpendicularly to the solar wind velocity vector: the focus is located at 0.5 Mars radius from the planet centre along the .\--axis of the aberrated Mars solar orbital co-ordinate\

134

et (11.: Solar wind measurements

J. G. Trot&on

Solar

Wind Dynamic

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Fig. 9. Solar wind density

as a function of velocity. These two solar wind parameters have been measured on Phoho.~ -7. in front of Mars. Solar flare events have been discarded

0228

Month

0309 0318 * 100 + Day (at 0000 PHOBOS 2 1989

Fig. 7. Comparison between the computed from the PWS and that derived from the TAUS Schwingenschuh rt ul. (1992)). March 1989 time interval

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(a 4 aberration is assumed). L is connected bow shock distance D,, as follows :

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solar wind dynamic pressure ASPERA measurements and observations (see Fig. 5 in during the 28 February-27

~-

L = (&-0.5RM)(

MAR 09,

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(2)

R, is the Mars radius; R, and E are equal to 3390.5 km and 1.03, respectively. With the help of Fig. I I we shall try to answer the following question: is there any relationship between the bow shock position and the solar wind dynamic pressure‘? A particular relationship would suggest that a Martian magnetic field influences in a sigwhere

Dynamic Pressure

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Fig. 8. Electric-field spectra of ion acoustic waves observed the Martian shock boundary. The upper frequency cutoff these waves depends upon the solar wind ram pressure

in of

I

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I

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1600

800

0

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Fig. 10. Martian magnetic moment as a function of obstacle altitude and solar wind ram pressure. assuming that the solatwind&Mars interaction is dominated by the presence of a planetary magnetic field. The two magnetic moment estimations given by Dolginov and Zhuzgov ( 1991) and Grafe ( 1992) are indicated by horizontal arrows. The vertical arrow points out the lowest subsolar altitude at which the Martian shock was rw crossed by Phho.~ -7

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with

2 mission.

( 1’393) ; the main dill‘ercnce hct\vcen thei] I 1 is due to the insutlicient time co\cragc 01’

plot and Fig. the PWS

the P/&o.s

Note that the strength

to detect its changes.

VerigAn

tield is compatlhle

during

place during

was maximum.

sic field

intrinsic

performed

bind

Vcnt15 \+~nd

that the Venus

wind

interaction

I() the solar

and

arc‘ ionospheric

ha\ also

in

ccrt;Lin


(Hreu\

(‘I t/l..

~comet inlcraclivn

lY8Y ). hut it i\ not the uubjt’ct of (hi\ paper

5. Conclusions The

M;II.C and plasma

dimensional

plasma

have probed

the solar

Februar)

instrument. wind

to 36 March

lY8Y.

nificant

manner

the interaction

.tnd the plancl. .ipproxiniatcd

Assuming

between

the solar

that the shock

by a conic scction.

surfxe

wind

may be

The

said above. it is pos4ble to asociale an .!_ value to each bow shock crossing by looking for the conic that passes through it : to do that WC have lo tis : to its mean value 1.02. Figure I I shows

\vith dots these \emi-latus ram pressure. from

the dadmi

apected iattcr

for

rectum

dots plotted line

a I .(I

cur\‘c has

Knowins an

The

which

I

been

obtained

the general

formula

in

the

bow shock

obstacle

(Spreiter

wind

I differ marked11

I x IO” T m1 magnetic

the subsolar

the empirical

values vs. the solar

in Fig.

shows

the mean subsolar

cstimatc

as

tendency

moment. following

distance

distance

(‘1 (I/.. lYh6)

The way.

DH\, \ve

DCIH bv using

:

_

between the solar essenlial

was by

comparing

for

plasma

prcssurc

bog

hhock.

the stud!

flow

The

solar

Mach number

and ;’ the ratio

we

\xlidatc

our

them

uith

those

ohtalned

buring

or determination

sit\, pro\ ided b!, PWS

appears

due to the impulsive

the fat

intenhith,

main

(onI!,

5

wind

ASPERA

‘Take ;’ = 2 and 121, = 8. we then have

is spinning.

The

ASPERA

locates I),,, the obstacle stituting

’.It I .57 R,, (Trotignon :I[ ;I subsolar for

1’1al..

altitude

1993).

of about

crossings this places

470 km.

Sub-

II

equation (1) we obtain c)I< into III \, = I .()I h 10” T m’ , assuming a 0.83 x IO-” dyn cm ’ mean solar wind ram presxurc. Applying equation (I ) again.

wc can nov,

given dynamic

tion (3) NC have xxxss Finally.

calculate

pressure.

equation

1I. the diwrcp;tnq

the obstacle

Substituting

for

to the subsolar

(2) gives

L.

Let

position II,,,

into

for

;I

equa-

bow shock distance. us come back to Fig.

bctwcen the dots

and the model

con-

mea~uremcnt\ by

difl’crenc

variation\

problems

;i\xx~ged

flop.

tn term\

To t)\ er

nialcl\~. AInio4t in ;I ratio

on 5horl

are linhcd

of thl\

cspeciall!,

minimize

linic and

t>i‘the \\;i~c‘ u5cd I0 dcliLer

the error

it

to rhc

111011~ direction I‘rom the

directIon

cause ;I stron g modulation

mcauremcn~s. dcnhit>,

01‘ mx‘I he den-

01‘ lhc instabilit>

met by ASPERA

They

and ASPER.

pr~xctlurcs.

Io Iluclu:ilc

nuturc

j and the de\,iation

I),,,,

best lit model to the shock

par-

but \omc mincer ditl?r-

(he>, ma!’ he c\plaincd

technlqueh

scales

given by PWS

in good agreenicnl.

enccs remain:

of specific heat\. = ().7X1),,,.

sol:~r

These

and the red planet. and it

that

Mind densities

are gcnsrall~

solar is the free-stream

of the

of the interaction

narro\\ acccptancc angle of the detcctor3

where ,2/,

I

perfornxxl

techniques.

The

(3)

ASPERA.

I .5 ALJ. from

P/~oh~jv _’ yield cs(imalion\

of the Martian

ametcrh arc fundamental

at

ob>er\ations

\.elocity and dynamic

\\ind upstream

and the three-

c\;periment.

near Mars,

bv these tn o devices on board oi the density.

PWS.

composition

ol‘ the

\I hen Ihc \pacecral‘t h:\ry on the PWS

determinations.

111~ data

have

and been

I3 min inter\ ~12. the \pin period approsia\‘a-age dcnGt\ e\:1lu:lt1on4

all nl‘thc

between four

and

;I quarter.

;II’C

126

acoustic wave spectra usually detected in the shock ramp depends on the solar wind speed: this has been verified on the PWS data. Unfortunately, we have no access to the wave vector and its deviation angle from the solar wind flow direction that are necessary to compute the exact Doppler-shift frequency. It implies that this ionacoustic wave behaviour cannot be used to give a valuable estimation of the solar wind speed. As expected from the heliomagnetic latitude dependence of solar wind properties, the solar wind density appears to vary inversely with the flow velocity. The plasma density decreases and the flow speed increases when the distance to the heliomagnetic current sheet increases. A plot of the solar wind dynamic pressure that prevailed around Mars during the Phobos 2 mission was presented. There is no correlation between the ram pressure and the semi-latus rectum, the distance of the Martian shock to the s-axis of the aberrated areocentric solar ecliptic system. This distance is measured from the focus of the mean shock model which is located (along the s-axis) at 0.5 Mars radius from the planet centre. This feature is similar to that observed with Piotww Vmrs Orbiter in the environment of Venus. Like Venus, Mars does not seem to possess any consistent intrinsic magnetic field. The magnetic moment of a possible dipole field is not believed to exceed 2.2 x IO” T m3. To set the idea, a magnetic moment of 10” T m’ corresponds to a magnetic field of 26 y at the surface of Mars, on the equator, that is to say about 1000 times less than at Earth. Comparisons with published solar wind density and ram pressure data obtained by the proton. alpha particle and heavy ion spectrometer, TAUS, have also been presented ; the TAUS spectrometer possesses a large field of view in the solar direction (40 x 40’ , as shown in Fig. 1) and should yield reliable solar wind measurements. In view of the comparisons shown in this paper, there is no drastic disagreement between the PWS, ASPERA and TAUS mean solar wind measurements.

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IX, 365 36X.

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of the Martian

observations. Verigin, 111. I., Gringauz. K. I., Kotova, G. .A., Shutte, N. M.. Rosenhauer, H., Livi, S.. Richter, A., Riedler, W’., Schwingenschuh, K. and Szegii, K., On the problem 01‘ the Martian 2 ‘TjZI’S 5p~ctroiiitAt’r atmohphere dwipation : P/IcJ/Jc~.\ results. .I. +Y~I/I,IY. RCA. 96, 19.315 IO_370, 194,I Verigin, M. I.. Gringauz, K. I., Kotova. (;. A.. Remizot. A. I’., Shutte, N. RL, Rosenhauer, H., I.ivi, S., Richter, A., Riedler, W.. Schwingenschuh. K., SzegS, K., Apath?. I. and Tatrallyay, M., The dependence ot‘ the Mxtian Magnetopause and bon shock on solar wind ram prczsure according lo /‘/whf~\ _’

TAUS ion qxx%rometer Ii03 I io9. IV))7

me;lsuremcnt\.

.I ~CVI/J/~I~FL,\ 9X.