Venera 11 and Venera 12 observations of e.u.v. emissions from the upper atmosphere of Venus

Ptmet. Per-n

SFUCZ Sci. Vol. 29, pp. 149-164 Press Ltd.. 1981. Printed in Northern

Iniand

VENIZRA 11 AND VENERA 12 OBSERVATIONS OF E.U.V. EMISSIONS FROM THE UPPER ATMOSPHERE OF VENUS J. La BERTAUX,

1. E. BLAMONT,

V. hi. IEPINB

Service d’A6ronomie du C.N.R.S., Venikres le Buisson, 91370 France and V. G. RURT,

Institute of Cosmic Investigation (XI),

N. N. ROMANOVA,

k

S. -0V

Academy of Sciences 88 Profsoyouznaya, Moscow 117485,

U.S.S.R.

(Received in final form 4 August 1980) Wet-The results obtained by two extreme ultra violet (e.u.v.) spectrophotometers flown near Venus on VENERA 11 and VENERA 12 in December 1978 are uresented. Detectors were placed at discrete wavelength positions to measure e.u.v. emissions from thk upper atmosphere of Venus while the spacecraft were drifting on their fly-by orbits. The emissions of HI 121.6 run (Ly-a), He1 58.4 nm, and 01130.4 nm were measured with unprecedented sensitivity and spatial resolution. An 01 signal of 500 Rayleigh (R) measured outside the disc suggested the existence of a large bulge of oxygen atoms. The e.u.v. emissions of two ionic species, 01183.4 nm and He11 30.4 nm, were measured for the first time in the atmosphere of Venus. The zero order detector of VJZNERA 12 indicated the presence of a very intense e.u.v. emission (28 kR) lying between the monitored wavelengths. This emission, which

was only 3 kR for VENERA 11, is likely to be associated with the solar wind-ionosphere interaction. An attempt to measure ArI and NeI resonance emissions failed.

The Lyman alpha (Ly-a) inte~lane~ background was 4 to 5 times larger than expected, suggestive of a very intense solar flux or an increase of the inte~I~et~ density. The dis~bution of hydrogen indicates two populations with temperatures of 400 and 700 IC.

similar to the e.u.v. airglow spectrophotometer flown on MARINER-10 (Broadfoot et al., 1977). Some differences as described below made it more sensitive to true signal and less sensitive to stray light, at the expense of larger dimensions and weight (15 kg). The general specifications of the instrument are listed in Table 1.

INlRODUCllON The VENERA 11 and VENEFW 12 space vehicles reached Venus in December 1978. Each descent probe separated from the spacecraft and executed a soft-landing on the planet while the remaining part of the spacecraft, passing by the planet at a distance of -4X lo4 km, was used as a telecommunication relay for the descent probes. Each fly-by spacecraft contained an e.u.v. spectrophotometer to measure resonance radiation of H, He, He’, 0, Oc, Ar, Ne, C, CO at discrete wavelengths between 30 and 170 nm. Both spectrophotometers, of nearly identical design, were operated during two to three hours around the pericenter and functioned properly. A preliminary discussion of the results is presented here, in order to permit early comparison with experimental results from other space missions to Venus.

TABLE 1. VEXJZRA 11 AND VENFRA TROPHOTOMETER

Weight Size Power Field of view Collimator Grating

INSIlZUMElW DESCRIPTION The multi-channel extreme ultra-vioIet spectrophotometer uses a mechanical collimator, an objective grating and 10 detectors placed in the focal surface of the grating at the focusing positions of 10 wavelengths of interest (Fig. 1). Its design is

Detectorslit Data rate 149

12 E.U.V. SPEZ-

SPECIFICATIONS

15 kg H=180mm; W=24Omm; L=57Omm 15w Colliiator: 10’ half-width X P Detector: 20’ x 1.1” total 9 grids; useful area: 58x 80mm; peak transmission: 50% Instrument SA (Jobin Yvon), holographic type; size: 90 x 75 mm; useful area: 86x 71 mm; 890.1 mm radius; aluminium made; 3200 lines mm-l; 40 nm blaze; dispersion in focal surface: 0.71 nm mm-l; platinum coating 2.8 x 9.5 mm; bk = 2 nm spectral resolution 100 bps

J. L. BERTAUX

150

et al.

Slits

570 mm

t

-7

FIG. 1. OPTICALSCHEMEOFTHESPECTROPHOTOMJTERFLOWNONVENERA~~ANDVENEFLA

12.

The holographic grating is tilted. Individual detectors are behind the slits.

In order to make the best possible use of the holographic grating in the available space, the focal line is tilted by -7” from the plane of symmetry of the mechanical collimator (Fig. 1). A careful raytracing analysis was performed to optimize the recording of the holographic grating (with the interferences of two laser beams) and the position of the ten detectors. The grating, with 3200 lines mm-l was blazed for 40 nm and provided a linear dispersion of 0.71 mn mm-’ in the focal surface. The mechanical collimator, of overall length 100 nm, contained 9 plates with lands and voids of 0.28 mm. As seen from the grating, the acceptance angle of the collimator was 10’ by 3” half-width. However, the size of the field of view (FOV) was limited by exit slits placed in front of each detector, of size 2.8 x 9.5 mm corresponding to a 20’~ 1.1” total FOV. The width of the slit covered 2 nm in wavelength. The grating, coated with platinum, had

a typical efficiency of 3.5% at 58.4 nm 2% at 121.6 nm, and 1.5% at 161 nm in the first order. The zero order, where the grating acts as a mirror, at most wavelengths had a greater efficiency than the first order. Ten identical detectors (electron multipliers) were placed behind the ten exit slits, at the wavelengths of the resonance radiation of a number of species: H, He, Ar, Ne, He+, O+, C and a molecular emission of CO (see Table 2). One detector was placed in the zero order and was therefore sensitive to the whole emission spectrum of Venus, as limited by the quantum efficiency curve of the electron multiplier (Fig. 2). Windows made of MgF, (transmitting for A > 115 nm only) and CaF, (transmitting for A > 125 nm only) were placed in front of H and 0, CO, C channels respectively (see Table 2). This precaution was taken to suppress any diffused light of short wavelength, or

TABLE 2.

Detector No.

Element

1 2 3 4 (W&A (vENAA12) 6 7 8 9 10

all He+ He Ne 0+

Sensitivity of detector 1 at wavelength A N, (counts s-l R-’ )

Sensitivity N(counts s-l R-’ )

all 30.4 58.4 73.6 83.4

1.6x lo-* 7.2 x 10-l 1.1 1.2

8.2 x lo-’ 5 x10-l 4 x 10-l 3.2 x 10-l

86.9 104.8 121.6 130.4 150 165.7

1.3 1.48 1.2x 10-l 7.2 x lo-* l.lxlo-z 1.2 x 1o-3

A (mn)

Filter transmission Tf

11) Ar Ar H 0 co C

MgFz CaF, CaF, CaF,

0.52 0.48 0.65 0.72

2.9x 1.7 x 2.6 x 1.04x 1.2 x 1x

10-l 10-l lo-* 1o-2 1o-3 10-a

Observations of e.u.v. emissions from upper atmosphere of Venus second order light, and to investigate the very intense emissions reported for MARINER- 10 (Broadfoot et aE., 1974) longward of Ly-a (121.6nm). There were two differences between the instruments placed aboard VENERA 11 and VENERA 12. Detector No. 5 was placed for Ar (86.9 nm) observation on VENERA 11, and for 0’ (83.4 nm) observation on VENERA 12. The slit in front of detector No. 9, for CO (1.50 nm) observations was 6 mm wide for VENERA 11 instead of the standard 2.8mm. Otherwise, overall performances of both instruments, as checked during calibration, were quite similar. For each detector the photon-induced pulses were added in a 12 bit register. Each second the last 8 bits in each register were transmitted to the telemetry system, except for the zero order detector, where 8 bits from 23 to 21° were transmitted. The real number of counts is, in this case, 8 times the number transmitted through the telemetry.

151

s

1 DETECTOR

L

FIG.

2.

&y

WINOOW TRANSMISSION

4

I

’

30

90

60

120

EFFXIEN~~CURVES FOR

150 h lnml

THE DIFFERENT

I

180 m

OF THE INSTRUMENT.

The number N of counts s-l is, for a monochromatic intensity I of one Rayleigh (R):

The

Quantum

Efficiency curve is typical of the curve of each detector.

Countsse' Rayleiqhm'

N = IS’r,r,rrQR

I

where 0 is the solid angle of the FOV: 0.56X 10m4sr, and the transmission r= of the collimator is 50%. The useful area of the collimator S = 46.4 cm-‘, and I = 1 R = (106/&r) phot s-’ cmm2sr-‘. After passage through the collimator, N’ photons R-I of a given wavelength are directed toward the exit slit of the detector: Nr=gSrGO=

103photons

‘\\ B

R-‘.

The efficiency 7%of the grating was measured in the first order and in the zero order, and displayed in Fig. 2 as a function of A. All the detectors had similar quantum efficiency curves Q(A), as displayed in Fig. 2. The transmission TVof the filters placed in front of detectors Nos. 7, 8, 9 and 10 are listed in Table 2 and plotted in Fig. 2. Values of rf equal for VENERA 11 and and ‘I-, were VENERA 12 within a few per cent. The sensitivity of the ins~ment N may be calculated and is recorded in Table 2 and in Fig. 3 as is the sensitivity No in the zero order. Because the grating is more efficient in the zero order than in the first order for A >55 nm, the instrument is more sensitive at these wavelengths in zero order.

I

I

30

60

90

CURVES FIG.3. SENSITIVITY AS THE COUNTING STRUMENT

RATE

IS UNIFORMLY

120

I

150 180 h fnm!

OF TIIE INSTRUMENT,

OF EACH

DElECl’OR

ILLUMINATED

TIC EMISSION AT THE RELEVANT

DEFINED

WHEN THE IN-

BY A MONOCHROMA-

WAVEJZNGTH,

FILLING THE

FIELD OF VIEW. curve marked zero order is the sensitivityof detector No. 1, placed at the position where the grating acts as a mirror. These curves are calculated from curves of Fig. 2 and other optical parameters.

The

J.

152

L. BERTAUX

Calibration of the entire instrument was carried out in a vacuum chamber, where a collimated beam of monochromatic light was directed toward the instrument and toward a e.u.v. diode, absolutely calibrated and obtained from U.S. National Bureau of Standards. The absolute sensitivity derived from these me~urements was a factor of 2-4 lower than the calcuiated sensitivity of Table 2. This could be due to a slight misalignment of the instrument collimator and the collimated beam. Time was not available to repeat twice the calibration procedure. The large signal which was observed by both instruments during the cruise from Earth to Venus, particularly on the Ly-ar interplanetary background, caused us to suspect that the real sensitivity of the instrument is best represented by the calculated sensitivity of Table 2. This vahre will be used in the discussion of the results. The observation of Ly-or from the extended I-I cythereocorona yielded, for VENERA 11 and VENERA 12 instruments, a ratio of sensitivities in zero order and first order: 3.6 and 3.8 respectively, in fair agreement with the ratio of calculated sensitivities of Table 2, 4.6. GEOMETRY OF OBSERVATTONS AND DATA TREA’IMRNT

Each spacecraft, after separation of the lander, was injected into a hyperbolic orbit with a peri-

Parameter Tie

of encounter(p¢er) GJlJ

Pericenter distance(km) Orbital vetaity at pericenter(km,s-l) Inclinationof orbital plane on orbitai plane of Venus Angular radius of Venus as seen from the pericenter (“1 Celestial coordinates of line of sight a: (“) (Epoch 1975) 6 (“f (Epoch 1975) Angle of line of sight with Venus-Sun line (“1 Minimum impact parameter of the line of sight (km) Maximum impact parameter (km> Tie of fust data point (UT) Tie of dark limb crossing (UT)* Tie of terminator crossing (UT)* Tie of bright limb crossing (UT)* Phase angle (Sun-Venus-S/C at pericenter 0

et al.

SUN

FXG.4. T&EG~O~~YOFOB~RVA~ONSIS~US~~. The lineof sighthas a fixedorientationand sweepsacross the illuminateddisc of Venus while the spacecraftdrifts alongits f&-byorbit. center distance of 4~ lo4 km. The geometrical characteristics of the observations, quite similar for the two spacecrafts are summarized in Table 3 and displayed in Fig. 4. VENEBA 12 reached the planet on 1978 December 21, followed by VENERA 11 four days later. The spacecraft was kept in a fixed orientation, with the line of sight sweeping across the planetary disc during the orbitai drift. The orbital plane was at -45” incl~ation to the ecliptic plane, the spacecraft coming from above the ecliptic plane. The pericenter was below the

VENERA 11 03h23minS8s Dec. 25,1978 40,412 6.03

VENBRA 12 03h26min42s Dec. 21,197s 40,355 6.02

46.78”

38.6”

8.52

8.53

142.3 18.04

139.4 19.6

162.6

159

3,380 50,000 2h34min13s 3hOlmin36s 3h02min30s 3h30min30s 19.23

* computed at a radius of 6200 km (or 170 km of altitude).

2,344 26,600 2h4Omin4gs 3h13minlOs 3h14min21s 3h44min56s 24.08

Observations of e.u.v. emissionsfrom upper atmosphere of Venus

FIG.

5.

THETRACKOFTHELINEOFSIGHTONTHE

PLANETARY

DISC IS REPRESENTED

153

IN PROJECTION ON THE

DISCOFVENUS,~THEHJILSIZEOFTHEFElDOFVEW(0.33X1.10). The

dark limb is crossed just before the terminator.

ecliptic plane on the day side of the planet. At this point, the angle S/C-planet center-sun was 24” for VENERA 12 and 19” for VENEBA 11. Because of this orientation, the line of sight, nearly parallel to the ecliptic plane, crossed fhst the dark limb of the planet, then the terminator, reached the point of minimum distance Pti to the planet center, and crossed the bright limb. The geometry of the observations is illustrated in Fig. 5 where the field of view of the instrument is projected on the planetary disc. Unfortunately, the axis of the FOV, perpendicular to the ecliptic plane, was not parallel to the limb at the position of bright limb crossing, but at an angle of 79” for VENEBA 11 and 61” for VENEBA 12. The size of the FOV projected on the planet was 775 x 117 km. The drift velocity of the line of sight was 6 kms-‘, or 360 km min-’ projected on the planetary disc. Various parameters relevant to the geometrical observations are plotted in Figs. 6 and 7 respectively for VENEBA 11 and VENEBA 12. Confusion of the respective parameters of the VENERA 11 and VENERA 12 instruments in a publication of preliminary results (Kurt et al., 1979) is corrected in Table 3. The content M’ of the 8 bit counter associated with each detector, was transmitted each second through the telemetry system, without zero reset.

,031 0

FIG. 6.

20

hPACTP

Lo

60

80

ARAMETEiR

SIGHT TO THE CENTER

loo

(MINlhlUM

120 140 160 Time (mid DISTANCE

OF LINE OF

OFTHEPLANETJASAFUNClTONOF

TIMEAETERTURNONOFTHEINSTRUMFl?4T.

J. L.

154 Solar

zenith anqle Pf

BERTAUX

f

I

I

I

40

30

,

et ai.

I

I

,

I

50

60

70

Time

(mink

FIG. 7. SOLAR 7.ENlTH ANGLE AT THE POINT OF INTERSECI’ION OF THE LINE OF SIGHT WITH A SPHERE OF 6200 km OF RADIUS AS A FUNCXYONOF TIME AFTER TURN ON.

The number of counts h4 counted from t to t+ Is is: M=h4’ft+ l)-M(t)+256k’ where k’ is unknown, but is such that A4 is positive. The minimum possible value of M is MO: M,=M’(t+l)-M’(t)

(mod256)

and M =M,+256k.

For all detectors, the successive values of MO for VENERA 11 are displayed in Fig. 8 with various linear vertical scales. It is obvious that Mm&f, at the beginning of the experiment, when the line of sight is off of the planetary disc. For detectors 1 and 7 (respectively zero order and Ly-cr channels), it is clear that when the line of sight approaches the planetary disc, M becomes larger than 256, and k is ~0. By continuity from the beginning of the experiment, it is rather easy to determine the value

FW.8. FOREACHCHANNELALL. DATA

POINTS OF f@,mR/% 11 INSTRUMENI’ ARE PLOlTED AS A FlJNCI’iON OF TIME FOR THE FIRST 70 MrNOF OPERATION.

The vertical scale is adapted to the measured number of counts for each channel. Note that for zero order and Ly-q the signal goes down to zero after reaching 256 counts.

Observations of e.u.v. emissions from upper atmosphere of

OF THE RESULl3

HI Ly-a! (121.6 mn)

There is a great similarity between the H Ly-cu curves (detector No. 7) recorded with VENERA 11 and VENERA 12 (Fig. 9) both on the disc, where the intensity is rather flat, and far away from the planet. At the time the instruments were turned off, the impact parameters of the line of sight were respectively 57 x lo3 and 29 x lo3 km for VENERA 11 and VENERA 12. At this large

10

t -7 -‘i,

HI

20 I

40 .

., 50 .>

.) 60

VENERA 11

l21.6nm

155

observation distance, all the emission (indicated in Table 4) comes from resonance scattering of solar photons of interplanetary H atoms. Our calculated calibration curve indicates that the counting rate of -23 counts s-’ would correspond to an emission rate of 885 R. This is larger than expected from a model calculation of interplanetary emission which predicts only 190 R, from the position of Venus (ecliptic longitude 114“) and in the direction 01= 142O.3 and S = 18”.04C in celestial coordinates). The interplanetary model assumed an H density in the nearby interstellar medium of n,= 0.1 atom cm+, a temperature T = 8.8 x lo3 K (Bertaux et al., 1977), a relative velocity of 20 km s-l of the solar system in the direction a = 252”, 6 = -15” (Weller and Meier, 1974), and a solar flux F, = 3.32~ 10” phot em-2A-1 s-l at the line center of solar Ly-a line. This reference value of F,, which is relevant to a medium solar activity, is such that solar gravitation is exactly compensated by solar radiation pressure. One possible explanation for this discrepancy by a factor 4.7 is that the solar flux was much higher in December 1978 than this reference value. Strong solar activity was reported in December 1978. Higher Ly-a! intensities than expected were also reported both from the disc of Jupiter and the interplanetary background in March 1979, (Broadfoot et al., 1979), implying in part also a large

of k to assign to each me~urement. Then, for each detector 30 successive points were averaged together, and the dispersion u was calculated. The ‘error’ bar is *ta/&% The results, for each spacecraft and each detector, are displayed in Figs. 9-18 on a logarithmic scale. Each vertical bar is centered on the average, its length being 2c&/%. Four detectors (Nos. 3, 4, 5, 6) of VENERA 11 showed a large counting rate, slowly varying, even when the line of sight was off the planetary disc. It is possible that gas released by the attitude control system of the spacecraft might be a source of rare gases (He, Ne, and Ar) which could be illu~nated by the Sun with resonance radiation. This problem also exists for the same detectors on VENERA 12 but to a much lesser degree. However, at the present time there is no completely satisfactory explanation of this effect. DISCUmON

Venus

7a

80

90

L2’.

:

P 5

k0: f-

e

Ll 10 0

Frcj. 9.

10

20

%-X3? SIGNAL OF DETECWR

SO

VENERA !Z LO

50

No. 7 (Ly-CU)IS

TIME min

Lf 60

PLO’l-IED

70

00

AS A FU?4Cl’ION

90

Id,

OF TiME ASTER TURN ON.

The time scale for VENERA 11 is displaced in order to have approx~ately coinciding disc measurements. Limb crossings are indicated by Ll and L2. Each point represents an average of 30 data points. The vertical bar represents the dispersion aroung the average vahze. Only the first 100 min of data are represented in all Figs. 9-18. The value of the interplanetary background, measured at the end of the experiment, is indicated.

J. L. BERTAUX etal.

156

TABLE

Element

A (nm)

Background far from the planet (counts s-l) V-11 v-12

; 3 z

;t+ He Ne

30.4 all 58.4 73.6

0.6 lo-25 -45

(VEN. 11)

O+

83.4

-37

(VE:.

Detector No.

12)

0.5 2.5 3.5

4.

Maximum intensity on the disc v - *iRayleigh) v-12 3x103 100 280 31

28 x73103 270 22.5

156

Ar

a6.9

6 7 8 9

Ar H 0 co

104.8 121.6 130.4 150

-23 22 2 1.5

10 26 4 2

10

C

165.7

1.5

2

4.5

(t)Absorption by CC+ okm- )

‘*‘solar flux at 1 a.u. (phot cm-a s-r A-‘)

Ekcitation factor at Venus &-“)

1.4x lo-l7 2.9 x lo-r7 2.3 x lO-‘7

1.44x lo’o’*’ 4.3 x 109 f*) 7.5 x 106

1.38X 10-4 5.2x 10-s 5.8x lo-9

7.4 x lo-=

1.27 x lop t*)

8.2 x 10-s

55

1.8 x lo-”

1.96x 10’

5.2 x lo-”

135 37.7 x 1oa 6.2 x lo3 2.7 x 103

129 41.9 x 103 6.2 x lo3 2.1+103

2.5 x lo-r7 7.3 x lo-a0 1.1 x 10-1s 5 x lo-‘9

0.56x 64x 6:8x 7.5 x

2.7 x 1O-7 6.6 x 1o-3 2.7 x lW4 Not apphcable

10x 103

15 x 103

0.6x lo-t9

lo* 1011(*) 10’ (*) 10s

9.2X t09’*’

(*I Solar flux appropriate for 19 February, 1979, according to Hinteregger (1979). For He+, He, O+, H, 0 and C, the total line flux in phot cm-* SC’is indicated. Appropriate solar linewidths have been taken into account to compute the

solar flux at line center and excitation factor. (t) From Sun and Weissler (1955).

increase of solar Ly-ar. Indeed, the solar flux measured on February 19, 1979 (Hinteregger, 1979) (Table 4) was 6.4~ lO1’ phot cme2 s-l, which corresponds roughly to 6.4 x 10” phot cms2 s-l w at line center, a value larger than any value previously reported. Nevertheless, our interplanetary Ly-a observation implies a solar flux of 15.5 x 10” phot crn-‘~-~ A-‘, which is larger by a factor of 2.4 than the Hinteregger value. Since it is difficult to imagine how the instrument sensitivity could be larger than calculated, this result means either that at the time of VENERA observations, in December 1978, the solar flux was at this extremely high value, or it means that the density no in the nearby interstellar medium has increased significantly in the last ten years, a rather extraordinary happening. If we reject this hypothesis, the use of no = 0.1 atom crnw3 will provide a scaling factor for the absolute determination of the H density in the upper atmosphere of Venus. When this background is subtracted from atmosphere observations off the limb, the measurements cannot be fitted by a single scale height distribution, but rather with two scale height of 400 and ==700 K. A refined analysis will be performed in order to determine if the hot population of atoms comes from the exobase, as suggested by a number of authors (e.g. Anderson, 197% or is created at a high altitude, as was suggested by Bertaux et al. (1978) from Iine profile measure-

ments performed with a hydrogen absorption cell on VENERA 9. The intensity is quite flat on the disc. There is no horizon brightening at the limb, which implies an optically thick medium. The intensity near the center of the planet and subsolar point is 38 and 42 kR for VENERA 11 and VENERA 12 respectively, whereas the U.V. spectrometer on Pioneer Venus reported a maximum intensity of 31 kR (Stewart et al., 1979). However, the solar zenith angle for the Pioneer Venus observations was larger than ours and comparisons should be made cautiously. The zero order light curve is very similar to the Ly-a light curve off the pIanet~ disc, implying that Ly-a is the principal emission observed by the zero order detector. The ratios of counting rates (No. l)/(No. 7) are 3.4 and 3.8 respectively for VENERA 11 and VENERA 12 in reasonable agreement with the calculated ratio of sensitivities of 4.6. On the disc, there are differences between detectors No. 1 and No. 7 light curves which will be discussed later. Hef (58.4 rim) Detector No. 3 of VENERA 11 showed a high background counting rate of unidentified origin outside the disc of Venus (Fig. 10). For VENERA 12 the counting rate was 2.5 cs-“. If a ‘dark count’ of 0.5 as measured by the detector No.

Observations of e.u.v. emissions from upper atmosphere of Venus 2 off the disc is subtracted it yields an intensity of 4 R coming from He atoms in the interplanetary medium which is in the general range of previous measurements, though it is lower than the 10 R reported for MARINER 10 observations. The maximum disc intensity is nearly the same for both spacecraft: 275 R. It is less than the 500 R reported for MARINER 10 observations (Kumar and Broadfoot, 1975), but the ratios of disc intensity to interplanetary He emission are approximately equal for both missions. This comparison suggests a difference in instrument calibrations or in solar flux intensities at line center. However, the solar flux of Hinteregger (1979) (found in Table.4) is higher by a factor of 1.65 than the one which was assumed by Kumar and Broadfoot in order to fit their data. It is likely that there is rather a difference in calibration of the two instruments. There is a possible contamination by 011 emission at 58.1 nm, as suggested by Delaboudinitre (1977). From a rocket e.u.v. spectrum of the terrestrial dayglow (Gentieu, Feldman and Meier, 1979) (Fig. 19), an upper limit of 15 to 25 R was assigned to 01158.1 nm, whereas an intensity of 620 R was assigned to 011 at 83.4 nm. Assuming the same excitation sources for 011 emissions in the two planets (electron ionization and excitation of atomic oxygen and direct photoionization) the ratio of 01158.1 nm to 01183.4 nm should be the same. Since our instrument measured 156R at 01183.4nm (detector No. 5 of VENERA 11) on the disc of Venus, it is likely that 01158.1 nm emission in Venus is smaller than 5 R. rd

lo b E

20

30

40

L\ He I M.4 nm

157

There is an important variation of He1 58.4 nm intensity across the disc, with a gradual increase from the dark limb, which lies very near the moming terminator. This might be an indication that the helium abundance above the COZ absorbing level is smaller at night than during the day, as a result of lower temperature during the night. Near the bright limb a small increase of 10 and 15 es-l is observed respectively for VENERA 12 and VENERA 11, with about the same morphology as the increase observed in the zero order for the two instruments. However, the effects are likely to have a different origin, since the zero order bulge in the VENERA 12 counting rate is 6.7 times larger than the one for VENERA 11. Detailed radiative transfer calculations will be necessary to interpret these helium light curves, since the optical thickness is larger than 1. OI (130.4 nm) The emission feature of oxygen at 130.4 nm is thought to be due to resonance scattering of solar photons and to excitation by collison with electrons. (Strickland and Thomas, 1976). The maximum intensity on the disc was 6.4 kR, both for VENERA 11 and VENERA 12. This is to be compared to a rocket measurement (Rottman and Moos, 1973) of 5.5 kR, MARINER 10 measurement of 17 kR (Broadfoot et al., 1974) and the recent measurement by Pioneer Venus of 3.8 kR on December 7, 1978 (Stewart et al., 1979), two weeks before VENERA measurements. These last authors argue that an intensity of 3.8 kR is too high 60

M

VENERA 1,

70

00

90

L!2

,,.,,Y” ,,,,,,,.. .,,,%.,. ,,,,,,., 1.,-%, ,,,’

FOGS.10 to 13. THESIGNALIS PLOTTJSD FOR EACH DETECTOR As ON FIG. 9. Note the high background counts off the disc for 58.4, 73.6, 83.4 nm and 104.8 nm detectors VENERA 11.

of

J. L. BEIRTALJXet al.

158 10

I 30.

20

Li 01

lso.4 “In

40

50

80

m

60

VENtRA ,*,,,,W”%.,,1.,, L2 ,,J ,,,,,,,11,11,111 ,11,

90

11 _

-1

w

VENERA 20

10

60

30

FIG.

10

L

I ’

7 H.II 30.4nm -ln

60 I y L2

50

VENERA 11

Ll

m

60

im

90

80

11.

40

30

20

12

SO

80

70

90

: 10.

FIG. 12.

10

t - 011

M

20

LO

03.4 “In 86.9 nm

m

60

30

so

90

L2

VENEFtA 11

Ll

-b

Ll I ‘0

10

2u

30

L2

I

VENERA t2 40

so

FIG. 13.

60

, m

TIME 30

min 90

im’

Observations of e.u.v. emissions from upper atmosphere of Venus to be explained with a quantity of 4% of 0 at 140 km of altitude, and that another source of oxygen emission excitation resulting from the photodissociation of CO, must be significant. The oxygen atmosphere is optically thick at 130.4 nm. Therefore we should expect, if the emission were due only to scattering of solar resonance radiation a rather flat light curve across the bright disc. This is not the case, especially for VENERA 11 (Fig. 11). Electron excitation must be an important mechanism, but the mechanism must have a variable intensity across the disc of Venus in order to account for the profile across the disc. There is a strange feature in the light curve of VENERA 11, after the crossing of the bright limb which is not present for VENERA 12. Over a range of -5 x lo3 km of altitude a nearly constant counting rate of 5 counts s-’ above the background can be seen. Up to now we have not been able to propose any instrumental effect that would explain such a feature, absent from all other channels of both instruments (except perhaps for La curves). This counting rate, if attributed to resonance scattering of 130.4 nm emission, implies an emission of 500 R, which would correspond to a column density of 1.9 x lo*’ atom cm-*. An excitation factor of g = 2.7 X lo-* s-l atom-’ was assumed, derived from the solar flux estimate for the day February 19, 1979. (Hinteregger, 1979). This column density, spread over two planetary radii, would imply an average exospheric density of 1.6 X lo3 atom cme3 of oxygen. It could be the result of a sporadic “hot” component of exospheric oxygen, existing preferentially on the day side of the planet and at high latitudes (afternoon). (Fig. 20). Since the U.V. spectrometer aboard Pioneer Venus is spinning and scanning in wavelength, this weak emission of 500 R would probably be difficult to find in the data, unless specifically searched for. HeII (30.4 nm) The curves of signal delivered by detector No. 2 placed at A = 30.4 nm for observation of He’ are very similar for VENERA 11 and VENERA 12. (Fig. 12). Far away from the disc, they both present a very low background counting rate: 0.5 cs-*, the lowest of all channels. Data points are not displayed on the figures if the counting rate is lower than 1. On the disc, the counting rate reaches 6 cs-I, but there is a significant signal near the planet, outside the disc (Fig. 12). It could be stray light diffused from the grating. Ly-C-Kis the most likely candidate. One way to put an upper limit to the contribution

159

of stray Ly-a to the signal of detector No. 2 is to write: Nz=nN,+NZ where Ni is the counting rate of detector i, N,’ is the counting rate relevant to true signal, and UN, is the Ly-ar stray light contribution. Application of the condition that N2’ must be everywhere a0 sets an upper limit for alv, and a lower limit for N,’ of =4.5 cs-’ on the disc. ,Even after subtraction of the maximum value for uN~, there is some emission left outside the disc, extending up to 4.6~ lo3 and 3.2~ lo3 km of the bright limb side for VENERA 11 and VENERA 12 respectively. The scale height, outside the disc is -3200 km. The value of 4.5 es-’ corresponds to 55 R. Since MARINER 10 was not able to detect an emission rate lower than 300 R (Kumar and Broadfoot, 1975) there is no discrepancy. However, taking a g factor of 1.4~ 10e4 s-l atom-” (emission rate per atom) the observed He+ intensity implies a column density of 3.9~ 10” at cmm2. In situ measurements of He’ ion densities have been made with the Orbiter Ion Mass Spectrometer and Bus Ion Mass Spectrometer on Pioneer Venus (Taylor et at., 1979), showing less than 2X 10’ ions crne3 between 200 and I700 km, the ionopause altitude, resulting in only 3 x 10” at cm-’ in the vertical This discrepancy could imply: (1) We see an emission which is not related to He” but to some other species. (2) Measurements were not made at the same place and time. The orbit of Pioneer Venus was nearly polar. (3) The g factor was much greater than assumed, owing to a large solar activity increase. (4) Mass spectrometer measurements are wrong for He’ on Pioneer Venus. (5) Another excitation mech~ism. 011 (83.4 nm) We report here the first e.u.v. observations of 011 emission in the atmosphere of Vends. Since the solar flux and the 0” densities are much too weak to produce significant radiation, the mechanism of excitation is probably electron impact ionization and excitation of 0, and direct photoionizationexcitation. A very intense line (620 R) was observed in the e.u.v. spectrum of the Earth (Gentieu et al., 1979), looking upward from an altitude of 210 km. 011 is optically thick in the ionosphere of both planets, Earth and Venus. If resonance scattering were the

3. L. BERTAUX et al.

160

dominant process of excitation, we would see twice more signal at Venus than at Earth. In fact, the maximum intensity (Fig. 13) seen on the disc of Venus is 156 R only, which proves that resonance scattering is of minor ~po~~~, if of any. In addition, the 011 83.4 nm is not uniform across the disc, with fluctuations of *20% over distances of 10” km or less. These fluctuations are likely to be related to variations of excitation mechanism rather than to variations of 0 atoms. Photoionization by the solar flux should be rather constant over the planet, whereas electron impact ionization and excitation of 0 could vary substantially, The decrease of light on both limbs is rather sharp, with a scale height of less than 180 km at the

Outside the bright disc, there is a great similarity of shape between the curves of zero order detector and Ly-a detector, which is normal since Ly-a! is the only strong emission present in the exosphere and in the inte~lanet~ back~ou~d. This allowed us to measure the ratio of sensitivities in the zero order and at Ly-cr, 3.4 and 3.8 respectively for

VENERA 11 and VENERA 12 in fair agreement with the ratio of 4.6 which was calculated from performances of optical elements. The zero order curves are also very similar to each other in level and shape outside the disc. On the contrary, when the line of sight crosses the bright disc, there is an enormous difference between the IWO zero order curves. A maximum counting rate of 8.4~ lo3 and 3.8 x 10” counts s-’ is reached respectively for VFNERA 11 and VENERA 12. (The numbers of Fig. 14 have to be mul~p~ed by 8 to obtain the exact counting rate). It is particularly clear on the zero order curve of VENEXA 12 that au additional emission, which is not Ly-u, appears on the disc. There is a sharp change in slope at the places indicated by an arrow. It should be recalled that Ly-ar scarcely varies across the disc, whereas the zero order radiation clearly changes. One can estimate the number of

first limb. Zero order detector

The detector No. 1 is placed at a position where the grating acts as a mirror. The whole emission spectrum is seen by the detector with a relative sensitivity curve iudicated on Fig. 2 which drops very rapidly above 130 nm. The signal is displayed in Fig. 14 for both ~s~rnents.

_n

10

ZERO

20 ORDER

100 ..

23 WtiER~l2

‘, I!2 ,

nG. 14. TW.rSIGNAL

OF DETJXTGR

No. 1 (ZERO

ORLXXR) IS PJ.XYITED

AS A FUNCMON

OF TIME AFTER

TURN

ON.

Unlike the other figuresthe curve for VENERA 12 is placed above the curve for VENERA 11. The

time scale for VENERA 11 is displaced in order to have coincidingdisc measurements.The arrows indicate a change of slope at the limb, indicativeof a strong emission(speciallyfor VJBEFU 12) on the disc. The vertical bar represents the dispersion around the average value. otf the disc the signal comes mainly from Ly-a emission. For this detector No. 1 only the true number of counts is 8 times the number indicated on the vertical scales. Only the first 1OOminof data are represented in ail Figs. 9-18.

Observations of e.u.v. emissions from upper atmosphere of Venus counts in the zero order channel which come from the emissions seen by the other detectors, by using the calculated ratio of sensitivities of zero order to other detectors of Table 2 (except for Ly-a! where the value 3.8 is taken). We find 4.9x lo3 es-l, in which Ly-ar accounts for 76%. The difference between zero order counting rate and this total is 3.5 x lo3 and 3.3 x 10” es-l respectively for VENERA 11 and VENERA 12. This is most likely the result of an additional emission, not looked at by the tixed position detectors. Since the average sensitivity in the zero order channel is - 1.2 count s-l Rayleigh-’ below 110 nm, the intensity of this extraneous emission is 3 kR for VENERA 11 and 28 kR for VENERA 12. The characteristics of this emission are the following: (1) it is restricted to the disc of the planet; (2) there is a bump near the bright limb, extending over -3x103km, both on VENERAll and VENERA 12; (3) it varies across the disc (a factor of -4 for VENERA 12) in a way which excludes long wavelength Rayleigh scattering of solar light; (4) it is very intense and highly variable a factor of 9 in four days. From this last characteristic we can probably exclude resonance scattering on neutral species. When the terrestrial e.u.v. spectrum is examined (Fig. 19), it is tempting to attribute our Venus e.u.v. emission to lines of 01, 011 and NII, all species present in the atmosphere of Venus. The total of the terrestrial spectrum, outside the lines that are observed by our Venus instrument, is 1.5 kR up to 110 nm, out of which the 01 98.9 nm line accounts for 56%. This line is produced by electronic excitation of 01 (Gentieu et al., 1979). This intensity is not too far from 3.5 kR for VENERA 11 but is very far from the 28 kR for VENERA 12. It is unfortunate that 01183.4 nm was measured only on VENERA 11 and not on VENERA 12; however, the structure found on this emission for VENERA 11 is not at all reflected in the zero order line, which is very smoothly varying across the disc. The extraneous emission is unlikely to be 011 58.1 nm, since on detector No. 3 (monitoring He1 58.4 nm) there is practically no difference between VENERA 11 and VENERA 12. Whatever the species responsible for the emission are, it is most likely connected with the direct interaction of the solar wind with the atmosphere of Venus. In this respect, it can be remarked that there was a large increase in the velocity of the solar wind protons measured by Pioneer Venus

161

orbiter the very day of VENERA 12 observations (Wolfe et al., 1979). This emission would also explain the MARINER 10 observations of Venus in the zero order channel where it was sensitive in the range 20-150 nm. There were two zero order channels in the MARINER.10 instrument, the second one being sensitive in the range 115-170 nm. Both recorded a very intense signal. The interpretation of the authors (Broadfoot et al., 1974) was to attribute the signal to a very intense emission longward of 135 nm, because two other detectors, placed at 148 and 165.7 nm recorded a rather high counting rate. We propose another interpretation. In the MARINER 10 instrument, the long A order (115-170 nm) detector, the 148nm and 165.7 nm were all equipped with a CuI cathode extending the quantum efficiency to longer wavelength than a bare channeltron surface. We suggest that Rayleigh scattering of solar light at A > 200 nm in the atmosphere is responsible for the high count rates found in the CuI detectors, through reflexion and diffusion on the grating. CuI cathodes may have a ‘tail’ of sensitivity at A > 200 nm. In contrast, for the zero order channel not equipped with CuI cathode, the high signal could be accounted by the same emission as our e.u.v. extraneous emission, below 110 nm. CO (150nm) The detector No. 9 was centered at 150nm wavelength to measure the intensity of the emission of the CO fourth-positive system A’B + X1x’. The number of counts above the background measured at the center of the disk (Fig. 15) is 3.3 and 2.6 counts s-l respectively for VENERA 11 and VENERA 12, which correspond to intensities of 2.7 and 2.1 kR, according to the sensitivity of Table 2. This detector was placed here to discriminate between two previously reported and conflicting observations: the rocket spectrum of Rottman and Moos (1973) and the MARINER 10 observations. For MARINER 10, a very high intensity of 55 kR was reported for the discrete channel placed at 150nm, (and 30 kR at CI 165.7 nm), and an enormous emission of 4X lo3 kR for the total bandwidth A 115-170 nm detected in the zero order channel 202. The present observations are in complete contradiction with MARINER 10 results. It can be pointed out that three channels of MARINER 10 for which unexpected high intensities were reported had a cathode with CuI, which might present a ‘tail’ of quantum efficiency at A >

162

J. L.

BEXTAUX

200nm, where the solar spectrum is Rayleigh scattered by the upper atmosphere of Venus. Reflected light in the zero order and scattered light by the grating could likely be the explanation of the high intensities reported by the e.u.v. spectrometer of MARINER 10. The CO fourth-positive system was first identified in Venus from the rocket spectrum. At 150 nm the measured intensity was -1 kR nm-*. The slit placed in front of the detector No. 9 of VENERA 11 was 6 mm wide (AA= 4.2 nm) instead of 2.8 mm (AX = 2 mm) wide for VFNERA 12. The VENERA 12 intensi~ of 2.1 kR is equal to the rocket spectrum intensity. For VENERA 11 the intensity is only 1.3 higher whereas the effective bandwidth is a factor of 2.1 larger. If the emission was more or less a continuum, as shown by synthetic spectra obtained by Rottman and Moos, we would expect a signal ratio equal to the bandwidth ratio. Since it is not the case, it is an indication that the emission bands are more narrow than predicted by the synthetic spectra and should present a low signal between the peaks. Indeed, the Pioneer Venus U.V. spectrum shows a strong structure in this region of the spectrum (Stewart, 1979, private co~unication). The synthetic spectra of Rottman and Moos (1973) were computed to simulate the various possible emission mechanisms: resonant and fluorescent scattering of solar radiation, direct electron excitation of CO, and dissociative recombination of CO,‘.

et ai.

(X(165.7

nm)

The detector No. 10 was placed at 165.7 nm to measure the line emission of CI, produced by resonance scattering of solar radiation and possibly dissociative excitation of CC&. The sensitivity of the bare detector is very low at this wavelength. The counting rate is, at the center of the disc, -1 cs-’ for VENERA 11 and -1.5 cs-’ for VENERA 12, yielding intensities of 10 and 15kR respectively. (Fig. 16). This is quite lower than the 30 kR reported for MARINER 10 observations, but higher than the value reported from the rocket spectrum of 4 rLr1.5 kR. However, in this case the estimate of the CI intensity was dependent on the intensity assigned to CO emission, since the spectral resolution was not very high. The spectrum taken by Pioneer Venus, with its better resolution (Stewart, 1979, private communication}, shows that the CI line stands well above the band emissions. It can be estimated that -85% of the signal measured in VENERA instruments may be assigned to CI line, yielding an intensity of -8.5 to 13 kR. It is likely that the MARINER 10 detector for CI emission which was equipped with a CuI cathode, suffered from the same problem as the CO detector. ArI (86.9 and 104.8 nm) The detector No. 5 was placed at 86.9 nm to observe ArI resonance lines at 86.6 and 87.6 nm (on VENERA 11 only), and detector No. 6 was

9. f+GS. 15tO 18. ~ES~GNALISPLO~~~REACHD~~RASGNFfG. Note the high backgroundcounts off the disc for 58.4, 73.6, 83.4 and 104.8nm detectors of VFNFRA 11.

Observationsof e.u.v. emissionsfrom upper atmosphere of Venus 10

20

30

40

30

m

60 I

Cl

Ll

W.?nm

60

163 90

L2

VENERA 11

-. m B 5

‘: to*,* g 3

,‘,I

u 1

10

20

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t2 40 . VENER.+SO

Fzo. 17.

60

VENERA I2

L2

TIME

,

70

min do

w

14

164

J.L. BERTAUX et al.

placed at 104.8 nm to observe another resonance line of ArI. The measured intensities (displayed on Figs. 13 and 17) are surprisingly high, reaching 55 R at 86.9 nm and ~130 R at 104.8 nm (for both instruments), when compared to a simple estimate achieved in the following manner. We consider an atmospheric constituent i distributed with a scale height Hi, and concentration nio at a certain level of reference z = zo, where the concentration of CO2 is no and H is the CO* scale height. The distributions of COz and of constituent i are n(z) and n+(z): (1) n(z) = noe-(r-rd’H (2) n,(z) = ~oe-“-‘J’H~. We assume a plane-parallel atmosphere and consider only primary scattering. The optical depth of CO, at altitude z is: (3) T(Z) = m(z)H

= o-noHe-“-‘~‘H

where u is the absorption cross-section of CO, at the wavelength considered, respectively 1.8 x 10-l’ and 2.5 x 10-l’ cm-’ (Sun and Weissler, 1955) at 86.9 and 104.8nm. From Pioneer Venus results of the Bus Neutral Mass Spectrometer, no= lO*l crnm3 at zo= 135 km and H =4.3 km (Von Zahn et al., 1979). The vertical depth of penetration of the solar flux in the absorbing atmosphere of CO, is defined at the altitude where T = 1, and is 133.9 km at 86.9 nm and 135.3 km at 104.8 nm. The column density of Ar above these levels depends on the mixing ratio of Ar/C02. Results of the Bus Neutral Mass Spectrometer were first reported to show a value of 80 ppm both for A? and Ar4’ (Von Zahn et al., 1979). However, a further interpretation of the measurements same (Mauersberger, Von Zahn and Krankowsky, 1979) yielded only an upper limit of 41 ppm at 135 km (total for both isotopes) and no positive identification. Assuming a value of 100 ppm at 135 km, a rough estimate of the resonance line emissions of Argon can be established. The column densities are 6.2 x 10” and 4.6 x 1O1* atoms cm-’ above the absorption levels at 86.9 and 104.8 nm respectively. The solar flux F,, measured in February 1979 at one AU (Hinteregger, 1979), was 1.96~ lo* and 0.56X 10’ phot cm-* s-l A-’ at these two wavelengths. Taking the oscillator strength f = 0.2 and 0.254, the excitation factor g is, at Venus, 5.2 x lo-’ and 2.7 x 10m7phot s-l respectively at 86.9 and 104.8nm. The expected intensities I(Rayleigh) = 10e6gN are 3.2 and 1.3 R respectively at 86.9 and 104.8nm for a mixing ratio of 100ppm. These expected values are much lower

than the measured values (55 and 13OR), by a factor of 17 at 86.9nm and 100 at 104.8nm. The most likely explanation is that there are other very strong emissions in addition to ArI emissions in the spectrum of Venus. The terrestrial e.u.v. spectrum shows a line of -10 R near 86.9nm, which has been tentatively attributed to Nz at 86.7 nm or to 011 at 43 nm seen in the second order of the grating. The holographic grating of our e.u.v. spectrometer has a very low efficiency in the second order and 01143 nm is therefore an unlikely candidate. The concentration of N2 has been measured in the upper atmosphere of Venus (Niemann et al., 1979) with Nz = lo8 crnm3 at 155 km. Extrapolating these measurements down to the level of CO* absorption, a column density of 8.7 X lOI4 mol crne2 of N, is available for such emission. Detailed calculations of various possible emission mechanisms will be required to investigate further this possibility. Around 104.8 nm in the terrestrial e.u.v. spectrum (Fig. 19), there is an emission of 27 R attributed by Gentieu et al. (1979) to Nz emissions between 105.2 and 105.8nm, rather than to ArI resonance line. The terrestrial ratio of Nz 105.2 nm to N2 86.7 nm emissions is 2.7, quite near the mean ratio observed in Venus of 2.4 in the present experiment. However, the shape of the curves recorded for VENERA 12 at 86.9 nm and 104.8 nm are not rigorously identical, and the ratio changes from place to place, which is indicative that the recorded emissions cannot be entirely explained by the same mechanisms. Ne I (73.6 nm) The detector No. 4 was placed at 73.6nm to observe NeI resonance line. The measured intensities are 31 and 23 R near the center of the disc (Fig. 18) respectively for VENERA 11 and VENERA 12. This is much more than can be expected from a simple estimate similar to that made for Ar. The absorption cross-section of CO:! is 2.3~ 10-‘7cm2 (Sun and Weissler, 1955) and the depth of penetration of radiation at 73.6 nm is 135 km. If a mixing ratio of Ne/CO, = 4.3 ppm is taken at this altitude, a column density of 3.2X 10” at cm-’ is available for resonance scattering. The mixing ratio is based on the low atmosphere measurements by mass spectrometry (Istomin et al., 1979) from VENERA 11 and VENERA 12 descent probes and the altitude profile of CO, as measured by Pioneer Venus, with a turbopause at 137 km (Von Zahn et al., 1979) is taken into account.

165

Observations of e.u.v. emissions from upper atmosphere of Venus

FIG. 19. THE E.U.V.

SPBCT’RUM

RJXCORDEXI IN THE TERRESTRIAL

DAYGLOW

BY A ROCKET

BORNE

INSTRIJ-

and ~@ER, 1979). The position of detectors Nos. 2-8 of our VENJXA e.u.v. spectrophotometers is indicated, as well as the wavelength bandpass of each of them. G-u,

MENT (AFIBR

FFXDMAN

t(R)

ItRl

At this waveIength, the terrestrial e.u.v. spectrum shows no identified emission, and the signal is less than =5 R. Because of the large background on VENERA 11 detector No. 4, it is difficult to compare from Fig. 18 the shapes of the two curves.

10s

CONCLUSION i0'

:

0

10

20

30

40

50 60

70 00 ME

FXG. 20.

90 loo (Mint

FOR 130.4nm OF VENEBA 11. The orbit passes behind the planet. The line of sight is directed towards the reader. The bulge is located at high latitudes in the afternoon side. GBOMETRY

OF OBSERVATION

The solar flux at 73.6nm is 7.5X lo6 phot cm-* s-l A-’ (Hinteregger, 1979, private ~ommunicntion) and the excitation factor is g = 5.8~ 10-ps-‘. Therefore a signal of -2x 10M3R only is expected from NeI resonance, to be compared to the observed 27 R. Another source of emission is responsible for the Venus observations.

Results obained by the two spectrophotometers flown to VENFRA 11 and 12 were very similar, an indication of their overall satisfactory performance. The emissions of HI at 121.6 nm, He1 at 58.4nm, and 01 at 130.4nm, which had been identified in earlier observations were measured with an unprecedented sensitivity and spatial resolution. These measurements should be useful in a detailed study of the exosphere of Venus. A bulge in the emission (0.5 kR) at 130.4nm has been discovered, extending up to 5 x lo3 km of altitude. Emissions of CO at 150 nm and C at 165.7 nm were detected. These can be studied in greater detail by the U.V. spectrometer on the Pioneer Venus Orbiter. For the first time, e.u.v. emissions of ionic species have been measured in the Venus atmosphere: 0183.4 nm and He11 30.4 nm, this latter emission extending unexpectedly very far from the disc. A detailed study of the result will allow some insight on the dynamics of the ionosphere.

166

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An attempt to measure ArI and NeI resonance emissions failed, because of serious cont~ination by other very strong e.u.v. emissions of unknown origin. The zero order detectors showed, particularly for VENERA 12, that there is a large e.u.v. emission between the wavelengths where the detectors were placed. Because of its sporadic character this emission is probably connected with the solar wind interaction. Many questions which are still open after this inves~gation could probably have been solved if a full e.u.v. spectrum had been obtained. This could have been achieved with this ins~ment, or could be achieved in the future, with a similar ins~ment if the grating were given a rotating capability. This was not possible in this experiment for reasons independent of the will of the authors. Acknowledgements-This experiment was a cooperative effort between the Laboratory of UV Astronomy at the Institute of Cosmic Research (IKI) of the Academy of Sciences of USSR and Service d’A6ronomie du C.N.R.S. in France. We wish to thank particularly J. F. Brun, who managed the project, and C. Taufemesse, for the opticat design of the ~s~ment. Data reduction was performed under the r~ponsib~i~ of V. M. Pokrass at the Computer Center of the Institute for Space Research, Moscow, M Monchy at the Division Ma~~matique du C.S.T., CNES Totdouse. and J. C. Lebrun at the Service d’Aeronomie. A preliminary data anaIysis was made by J. Afsa and J. P. Chassetiere as a part of their courses at the Ecole Polytechnique. We are grateful to T. M. Donahue for the revision of the manuscript. The French effort was sponsored by CNES under contracts Nos. 76 CNES 201 and 77 CNES 201.

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N. W., Hunten. D. M.. and Cariaaan, 6. R. (1979). Science, Wash. 203, 77% Rottman, G. 3. and Moos, H. W. (1973). J. geophys. Res. 7% 8033. Sun, H., and Weissler, G. L. (1955). .Z. Ckem. Pkys. 23, 1625. Stewart, A. I., Anderson, D. E., Jr., Esposito, L. W. and Barth, C. A. (1979). Science, Wask. 203,777, StrickIand, D. J. and Thomas, G. E. (1976). PIunet. Space Sci Zq 313. Taylor, H. A., B&ton, H. C., Bauer, S. J., Hartle, R. E., Donahue, T. M., Cioutier, P. A., Michel, F. C., Daniell, R. E. and Blackwell. B. H. (1979). Science, Wash. 203, 752. Von Zahn, U., Krankowsky, D., Mauersberger, K., Nier, A. O., and Hunten, D. M. (1979). Science, Wash. 203, 768.

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