IL NUOVO CIMENTO VoL. 4 C, N. 6 Novcmbre-Dicembre 1981

Quiet-Time Solar-Wind Ionic Composition.

V. FORMISAN0

Sl~ace Science Department, E S A - E S T E C - Noordwijk, The Netherlands

S. ORSINI

.Laboratorio Plasma Spazio, CNR - .Fraseati, Italia

(ricevuto il 26 Gennaio 1981; manoseri t to revisionato ricevuto il 3 Dicembre 1981)

S , ,mmary . - - The ISEE-B EGD solar-wind p lasma exper iment often observes heavy ions in the solar wind. Selecting two periods of quiet solar wind with no parMcles observed backstreaming from the Ear th ' s bow-shock, we show tha t different ion species follow similar t rends in terms of changes in the bulk speed and changes in the part icle number density. In par t icular , quiet-solar-wind ionic abundances in these two cases are (relative to II +) aIIe2+= (1 - -5 ) -10 -2, 8He2+= 1.5.10 -a, 0 +e= ( 3 - 1 5 ) . 1 0 -4, O+V=3.10 -~, S i = l . 1 0 -~, F e = 0 . 5 . 1 0 -~. We show that , if ~He ~+ abundance increases, the same occurs for the Si, 0 +7 and 0 +e abundances. F ina l ly the O+e densi ty fluctuations seem to be anti- correlated with 4He~+ densi ty fluctuations.

1 . - I n t r o d u c t i o n .

I o n s h e a v i e r t h a n p r o t o n s a n d d o u b l e - c h a r g e d h e l i u m h a v e been d e t e c t e d

in t h e so lar w i n d b y u s i n g m o s t l y e l e c t r o s t a t i c a n a l y s e r s (1-~). I n f o r m a t i o n

on ~ e a n d nob le -gas a b u n d a n c e s has also been o b t a i n e d b y e x p o s i t i o n of

(1) S . J . BAM~, A. J. IIu~DHAUS~N, J. R. ASBRIDGE and L. B. STI~O~G: Phys. Rev. .~ett., 2@, 393 (1968). (2) S. J. BAME, J. R. ASBRIDGE, A. J. HUNDHAUSEN and M. D. MONTGOM:ERY: J. Geol~hys. l~es., 75, 6360 (1970). (~) M. B. CATTAX~O, V. FORmSANO, G. MOI~.~O, E. PALM:OTTO, F. PALUTA.W and P. SA~AC~NO: Sol. Phys., 17, 468 (1971).

682

QUIET-TIME SOLAR-WIND IONIC COMPOSITION 6 ~

a lumin ium foil to the solar-wind part icles (4) or b y count ing pits in an exposed mica foil (5). There is no doubt , however, t h a t the informat ion avai lable on ions heavier t h a n hel ium in the solar wind is ra ther scarce. B ~ . et aI. (e) based their s tudy upon 17 hour ly measurements in the solar wind to find t h a t on the average

and

(N(Fe)/N(H)) = 5.3.10 -5 , (N(Si)/N(H)) = 7.6.10 -5

(2V(O) /2V(H)) ---- 5 . 2 . 1 0 - ' ;

the evaluat ion of the oxygen abundance was less cer ta in due to H e +2 contamina- t ion. These abundances were observed to va ry b y a fac tor 10, 4, 4, respectively, wi th some evidence t h a t Si and Fe abundances m a y be correlated. BAlm, et al. (7) and FERVOR-. (9) have pu t the heavy- ion observat ions in relat ion with solar-wind and solar-corona propert ies .

Due to ve ry high sensi t ivi ty of the solar-wind exper iment flown on board of ISEE-2 , m a n y solar-wind heavy- ion spectra have been c o l l e c t e d ~ t h e pre- l iminary s tudy of which we present in this paper .

The E G D solar-wind p l a sma exper iment has been described in (9); here we repea t the re levant informat ion. The electrostat ic analysers provide 64 energy windows wi th AE/E ___ 5 % width. The coverage is continuous in the energy range 350 eV, 4260 e u 8 more energy windows are displaced between 55 and 350 eV and other 8 energy windows are displaced between 4260 and 11000 eV. I t is, therefore, evident tha t , due to the low t empera tu r e of the solar-wind ions, while in the best cases (low solar-wind speed) we can observe ions wi th M/Q up and above 10, in mos t cases (solar-wind speed ~ 400 l~m/s)

detai led s tudy is possible up to M/Q ~__ 5, the other 8 channels (between 4260 eV and 11 000 eV) providing some possible informat ion, bu t not all the needed informat ion, on the presence of heavy ions.

Here we present a pre l iminary s tudy of our da ta , a iming to cover some

(4) J. GEISS, P. ]~BF-RHARD, F. BUHLER, J. MEIST~R and P. SIGNER: J. Geophys. Res., 75, 5972 (1970). (s) E. ZINNER, R. M. WALKER, J. BORG and M. MAUR]~TT]~: Measurement el heavy solar-wind particles during the Apollo 17 mission, in Solar Wind Three, edited by C. T. RUSSEL (Los Angeles, Cal., 1974), p. 27. (~) S . J . BAME, J. R. ASBRIDGE, W. C. F]~LDMAN and M. D. MONTGO~]~I~Y: Sol. Phys., 43, 463 (1975). (~) S . J . BAMv., J. R. ASBRIDG~, W. C. F~LDMAN, E. E. F~IMORE and J. T. GOSLING: Sol. Phys., 62, 179 (1979). (s) E. E. FE~IMORV.: Solar-wind ]lows associated with hot heavy ions, in press in Astrophys. J. (1979). (9) C. BONIPAZI, P. C~.RULLI-IR~LLI, A. EGIDI, V. FORMISA~O and G. MORENO: IEEE Trans. Geosei. Electron., GEO-16, 243 (1978).

684 v. FORMISANO and s. ORSINI

of the gaps present in previous studies. We first discuss (sect. 2) the charac- teristics of the heavy-ion spectra, then present results of the analysis for day 310-311 ]977 (sect. 3), in which for more than 20 hours heavy-ion data were collected, then a prel iminary s tudy on statistical characteristics will be given (sect. 4).

2. - Some properties o f the h e a v y - i o n spectra.

I t is well known tha t af ter the passage of solar-flare-related in terplanetary shock waves, the solar wind shows an increase of helium abundance (see the review by F O~SANO and MOREiNO (t0)). In those cases not only the density of H + and He +~ is larger, bu t also the densi ty of the heavy ions increases (see(l,a,7)). To have an unper turbed spectrum, however, we need to have the magnetic field passing through the spacecraft not connected with the Ear th ' s bow-shock; in this way no backstreaming particles will be present, and no per turba t ion of the solar-wind spectrum will occur.

This is the case for the two spectra shown in fig. 1. ~ o particles back- streaming from the Ear th ' s bow-shock were observed in ~his period. The data where taken on day 8, 1978 between 08.00-09.00 U.T.; an in terplanetary shock hit the Ear th ' s magnetosphere on day 5, ]978 at 19.00 U.T.; our data refer, therefore, to the solar-wind conditions (low bulk speed, very low tem- perature, see table ]) characteristic of the post-flare plasma. In our spectra we note, indeed, another feature of the post-tlare plasma: the presence of a two-peak proton energy spectrum, as we can see from fig. 1.

As seen from table I, the solar-wind bulk speed was V----- 415 (or 419) km/s, consequently not all the possible ion species were detected in the two spectr'~ shown in fig. 1, due to the part ial energy coverage above 4.3 keV. Still the data show, above the H + peak, a peak at energies ].56, 2.00, 2.39, 2.78, 3.07, 3.68, 4.28, (4.73), (5.97), 7.71, 9.41. In brackets we have the location of two possible peaks tha t in tig. I are separated not by a valley but by an uncovered energy range; we have reasons to believe, from other measurements, t ha t they correspond to two actual peaks. Ahnost all these peaks were observed before by ]3A~E et al. (1) and were identified as due mainly to 3He+~, 4He+~, O +7, 0 +6, Si +9, Si +8, Si +7 (plus Fe+:a), Fe +'~, ]~'e +~~ Fe +8, respectively.

The last peak at E/Q of 9.41 has never been reported before. I t appears in our data several t imes in these periods also with fluxes higher than the one shown in fig. 1 at 08.40 U.T. :If in terpre ted as an iron ion, this peak should be due to ) 'e +~.

The two spectra shown in fig. 1 have been analysed, for ions other than H+ and 4He2+, with a zero- temperature approximation, i.e. all the counts in

(10) V. FOR~lSA~O and G. MolcE_~o: Riv. N,~wvo Cimento, 1, 365 (1971).

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~ V. F O R M I S A N 0 a n d 8. 0 R S I N I

a given peak have been assumed to be due to the same species with no contam- ination from others, therefore they represent the integral flux of each ion

species. The results of this analysis are shown in table ]. During this disturbed period (helium enriched) 4He+~ went up to 22 % and O +6 up to g 1% of the

10 4

10 3

{2 :b

810:

101

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a) i-

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10 2 10 3 10 a 10 2 10 3 10 a

Fig. 1. - Solar-wind energy spectra showing heavy-ion enrichment in a piston of an interplanetary shock wave. The ions identified are indicated in the figure. Note the two peaks in the proton energy spectrum, typical of post-shock solar-wind flow. a) Day 8, 1978, 08.40 U.T.; b) day 8, 1978, 08.19 U.T.

H + number density. All the other ions do show a similar behaviour, and have

abundances much higher than those reported by BA~_E et al. (8). F rom fig. 1 we see tha t the O +e flux m ay change very much (a factor 3) on a shoit t ime

scale. This density change must be real, because the adjacent channels do not show any significant change of counts. The ahem+ abundance was on day 8,

QUIET TIME SOLAR-WIND IONIC COMPOSITION 0~7

1978 at 08.19 U.T. , 15.10 -4 cm -3, i.e. for the ra t io of the number densities

~(aHe2+~/N(H+~ ~ 10- ' and

N(3tte,+>

in close agreement with the 310 -4 found by BAME et al. (7). Other two spectra more typ ica l of the solar wind are shown in fig. 2 for

day 321, 1978, 08.39 U.T. and day 310, 1977, 17.32 U.T. I n the first case an in te rp lane ta ry shock hi t the magne tosphere on day 3:[6, i.e. 5 days before our observat ion, while in the second case the last in te rp lane ta ry shock t h a t hi t

10 5

a)

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10 2 10 3 10

b) H*

4He2"

0 +7

_ 1 ~

I Si 9 Ilsr ~

10 10 I~ ener'gy/Q (eV)

Fig. 2. - Solar-wind energy spectra showing heavy ions in an , average, situation. The ions identified are indicated in the figure, a) Day 310, 1977, 17.32U.T.; b) day 312, 1978, 08.39 U.T.

688 v. FORMISANO a n d s . ORSlNI

the magnetosphere was observed 11 days before. Bo th these cases, therefore, are not flare-related events , and m a y be considered, specially the second one, as typ ica l solar-wind conditions wi th bulk speed and the rma l speed smaller t h a n average. A simple inspection of fig. 2 tells us how much var ia t ion of ion abundances we should expect in the solar wind.

The quan t i t a t ive analysis of the da ta (see table I) tells us tha t , on day 310, ~He ~+ has a num ber densi ty _~ 6 t imes smaller t h a n on day 321, 1978, while the O +~ numbe r densi ty is prac t ica l ly the same in the two cases (6.7.10 -3 and

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10 3

i 10

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a)

He +2

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II Si

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H+ b) [ H+ c) ]

He +2

'1 I I 0 +~

I I I Io-

Fe

| i~1 I I I IIII 1r ld 10' lo ~ 10'

energy Q(eV)

Fig. 3. - Fluctuation of the heavy-ion abundance on a (< short ~ time scale. Note the differences for O+e abundance between 01.30 (a)), 0217 U.T. (b)) and 02.44 U.T. (c)) on day 311, 1977.

7.4. :10 -3 P / cm 3, respectively). The heavier ions do not show a similar behaviour , bu t r a the r they go like the ~He 2+ num ber density. I t is interest ing to note tha t , a l though t a k e n in an (< average ~ solar-wind condition (from a densi ty .point of view at least, see tab le !) , the spec t rum of day 310, 1977 shows the SHe~+ peak giving us a rat io of N<SHe~+>//q<'He2+> ~ 2.2-10 -3, which is ra ther

QUIET-TIME SOLAR-WIND IONIC COMPOSITION ( ~ 9

higher t h a n the value obta ined for fig. 1 cases, bu t which m a y be more typ ica l of <( average )) solar wind (see below in this paper) . I t is clear t h a t the abundances we are going to derive f rom inspection of single spectra m a y change ve ry much due to s tat is t ical f luctuations, t h a t are ve ry i m p o r t a n t when we s tudy the abun- dances of ions heavier t h a n 0+% and due to real oscillations of which we have evidence in the 4He~+ component , bu t which m a y concern also other ions, a t least 0 +6 . An indication in this sense we have f rom inspection of fig. 3.

Here we show three spectra observed on day 311, 1977. The H + temper - a ture was ra ther small, and several energy channels with no counts separa te the p ro ton peak f rom the 4He~+ one. Heav ie r ions are present in all three spectra, however, due to s tat is t ical f luctuations, Fe ions are absent in the first spect rum, Si ions are absent in the second spect rum, while the two groups appear together in the th i rd spec t rum.

3. - Shor t - t ime- sca l e behav iour o f the so lar -wind h e a v y ions : days 310-311, 1977.

On this day the solar wind was r a the r cons tan t for 20 hours with a bulk speed around 300 km/s , a t he rm a l speed ra the r small ( g 1 6 km/s equivalent to 1.104 K) and a p ro ton n u m b e r densi ty around 8 P / c m 3 with occasional de-

crease to 6 P / cm 3 or increase up to 11 P / c m a (see fig. 4). The heavy ions are observed clearly th roughou t all this period because of the absence of back-

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Fig. 4. - Solar-wlnd parameter for day 310-311, 1977. The gaps in the data are caused by high-time-resolution data, not included in the figure. ISEE-2SW experiment.

~0 V. FORMISANO a n d 8. ORSINI

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QUIET-TIME SOLAR-WIND IONIC COMPOSITION 6 9 1

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6 9 2 v . FORMISANO a n d s. ORSI~I

s t reaming part icles for this long t ime interval . I n fig. 3 we have shown some ex t reme s.w. spectra observed in this period. ~ o s t of the t ime, however, the f luctuat ions ale less pronounced, and we can follow the t ime deve]opment of a t least 8 peaks ident i fying the hydrogen, helium, oxygen and silicon ions. The iron ions are also present mos t of the t ime, however they m a y disappear below our background level (1 count level) because of the i r lower fluxes. We have analysed the da ta in order to obta in the part icle n u m b e r densi ty and bulk speed for all the observed peaks. Figures 5 and 6 show two short periods of this day. I n fig. 5 we present the bulk speed (a)) and the number densi ty (b)) for H +, 4He2+, 0 +e, 3He2+, 0 +7, Si +~, Si +8, Si +7. I n the n u m b e r densi ty plot the

Fe ions are also shown added toge ther (like the Si ions). This period has been selected because it shows a change in the bu lk speed of all ions f rom 300 to 315 km/s (probably a tangent ia l discontinui ty) a t 16.30 U.T. The n u m b e r densities of all ions show a similar t ime behaviour on the average decreasing f rom a m a x i m u m at 15.10 to a m i n i m u m a t 16.50 U.T. The values of the number densities of the heavy ions f luctuate very much, and cer ta inly much more t h a n the H + n u m b e r density, p robab ly due also to larger error values (at least for the less abundan t ions).

I n fig. 6 similar da ta are shown for the period 01.30-03.30 of day 311, 1977. I n this per iod the bu lk speed was more or less cons tan t a round 310 kin/s, and fig. 6a) shows t h a t the bu lk speed ha4 the same behaviour for all the ion species. The t I + number densi ty, oll the other hand (see fig. 6b)), showed an increase f rom ( 5 - - 6 ) P / c m 8 a t 0].50 U.T. to ( 1 0 - - 1 1 ) P / c m 3 (at 03.00 U.T.).

An increase in the n u m b e r dens i ty is revealed in all the ion species. Fe ion densi ty changes f rom 3.10 -4 to 15.10 -4 P / e m 3, Si changes f rom (2--3) .10 -4 to 10.10 -4 P / c m 8, 8He2+ changes f rom 3.10 -4 to 8.10 -4, O +e and O +7 change

f rom 10 -8 to 4.10 -a and 2.5.10 -8, respectively. The mos t str iking fea ture of fig. 6b), however, is the observed ampl i tude of the f luctuations t h a t is ve ry large. These f luctuations m a y be of s tat is t ical na ture for the heavier ions bu t not for the 4He2+ ions t h a t show a ra ther high peak in our data . No te t h a t no part icles backs t reaming f rom the bow-shock were present in all these data . These f luctuations should, therefore, be real oscillations for the 4He2+ and, as we shall see in the nex t section, also for O +e. B y compar ing fig. 5b) and fig. 6b), we see t h a t be tween 15.30 U.T. (day 310) and 03.00 U.T. (day 311) the abundance of 4He2+ and O +e has decreased b y a fac tor of 5 wi th small changes of the abundances of the other ions.

4 . - S o m e s t a t i s t i c a l p r o p e r t i e s .

Withou t a iming to provide a real stat ist ical s tudy, which will be a rgumen t for a fu ture s tudy, it is of interest to ex t rac t f rom the da ta presented previously some other informat ion b y s tudy ing t h e m f rom a s tat is t ical point of view.

QUIET TIME SOLAR-WIND IONIC COMPOSITION ~ 9 3

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6 9 4 V. FOI~,MISANO a n d s . ORSIN]

F o r example~ h i s t o g r a m s of t h e ionic a b u n d a n c e s w i t h r e s p e c t to p r o t o n s h a v e

been p l o t t e d in fig. 7. The h i s t o g r a m s of t h e 4He~+ a n d O +s a b u n d a n c e show

two p e a k s as t h e y ref lec t t h e d i f ference a l r e a d y m e n t i o n e d before , b e t w e e n

t h e fig. 5 p e r i o d a n d t h e fig. 6 one.

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Fig. 8. - Scatter plot of 0+% 0 +7 and Si abundance v s . ~He 2+ abundance. Note a c l ea r ' t r end for larger abundances when 4IIe 2~- abundance increases. A linear best fit (in a log-log plane) is also given for the O+e-dHe 2+ abundance relationship.

The a b u n d a n c e of 4He2t was (1 - -5 ) .10 - ' - , for O +e i t was ( 0 . 3 - - 1 . 5 ) . 1 0 -3,

t h e 0 +7 a b u n d a n c e was 2 .10 -4, t h e Si a b u n d a n c e is 3 . ] 0 -3. I n fig. 7 t h e a r rows

i n d i c a t e t h e a b u n d a n c e o b s e r v e d d u r i n g t h e p o s t - s h o c k f l a r e - r e l a t ed flow on

d a y 8, 1978. I t is i n t e r e s t i n g to n o t e t h a t a l l t h e ions e x c e p t one show lower

Q U I E T T I M E S O L A R - W I N D I O N I C C O M P O S I T I O N 6 9 5

abundances in the period considered t h a n in the flare-related solar wind of day 8, 1978. The ra t io of the 8He2+ to 4He:+ abundance is lower in the flare-

re la ted solar wind t h a n in this period (5.10 -4 against 3-10-3). This fac t m a y in p a r t be due to the difficulty in detect ing 3He~+ wi th an electrostat ic analyser (its peak m a y be often hidden in the H + high-energy tail).

I t should also be noted t h a t there are measuremen t s (like those shown in fig. 3) in which only an upper l imit of (5 - -10) '10 -4 can be pu t for the ra t io _Y(3He~+)/.Y(4ne~+).

Another i n t e r e s t i ng stat is t ical result, different f rom wha t previously re- por ted b y BA~E et al. (6), is shown in fig. 8. Here we show scat ter plots of the heavy- ion abundances vs. 4He~+ abundances. We see t h a t O +~ and Si ionic abundances show a s t rong correlat ion with the 4He~+ abundance on short t ime

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10 - 3 10 -2 10 - I

N (~He2") /N (H + )

Fig. 9. - Scatter plot of 0 +s abundance relative to aHe2+ vs. the ahem+ abundance relative to H +. The two symbols refer to the two periods considered with -f- data points referring to the period 1500-1700 U.T. on day 310, 1977. For one period a linear best fit to the points in a log-log plane has been drawn to guide the eyes.

scales; the difference wi th the B a m e et al. results is unders tood in t e rms of dif- ferent t ime scales. Here we note t h a t the sca t te r of the points decreases when we increase the ionic abundances , which, however, does not imply t h a t the dispersion of the points is caused b y s tat is t ical problems only, because the 4He~+ abundance also shows a large dispersion, while there should be no problem for the stat is t ics of counted part icles.

We have finally studied the f luctuations of the number densities of O +e and 4He2+. Inspec t ion of fig. 5b) and 6b) shows t h a t the densities of these two

ions seem to be ant icorrelated in their oscillations. This is clearly shown in

6 ~ Y. F O R M I S A ~ O and S. ORSINI

fig. 9, in which a scatter plot of the rat io N(O+6)/2g(qte 2+) v s . * H e ~+ abundance is shown. The data f rom the two periods of fig. 5 and 6 have been shown with

different symbols. A linear best fit through the first group of points is also shown. I t is clear t ha t the O +s abundance with respect to helium decreases

(increases) when the 4He*+ abundance increases (decreases) in the fluctuations

of each period. We may perhaps conc]ude tha t density oscillations of large amplitude and opposite in phase are present in the O +6 ions as well as in *tie*+.

5. - S n m m a r y .

I n summary we have shown in this prel iminary s tudy tha t with the EGD plasma experiment on board I S E E - B we are able to detect several ion species

in the solar wind. We can follow the t ime development of eight of the ob- served peaks, and we have shown tha t the bulk speed of these ions (H +, *He 2+, 3He2+, O +', 0+% Si +9, Si +8, Si +') are very close to each other and change all

together when a discontinuity is observed. Occasionally, however, small changes in the bulk speed of some of these ions are observed. Also the number densities

roughly follow the H + number density, al though fluctuations are also present. The abundance of these ions in a quiet solar wind for the two periods studied

are 1 (It+), (1--5).10-* (*He*+), (1.5).10-*(Fe) in general agreement with previous findings. These abundances, however, m a y va ry very much during

flare-related post-shock flows. The abundances of *He *+, O +e and Si ions in

general increase together. The O +6 and *He *+ ions appear to fluctuate opposite in phase like 4He2+ and H + have been observed to do.

I n conclusion we stress the point tha t this is only a prel iminary s tudy of

ISEE-2 heavy-ion observations, more results will be obtained by a deeper

s tudy of the vast amount of data available.

�9 R I A S S U N T O (*)

L'csperimento sul plasma del vento solare dell'EGD ISEE-B spesso osserva ioni pesanti nel vento solare. Scegliendo due periodi di vento solare quieto in cui non si osservano particellc che rifluiscono dall'onda d'urto della Terra, si mostra chc differenti specie di ioni hanno andamenti simili delle variazioni della vclocit~ complessiva e dei cambia- menti nella densit~ numcrica dclle particelle. In particolare le abbondanze ioniche nei periodi di vento solare quieto sono in questi due casi (relativamente all'H +) aHe2+=(1--5).10 -2 , aHe2+=1.5.10 -4 , 04e=(3--15).10 -4 , 0+v=3.10 -4 , S i = l . 1 0 -~, Fe = 0.5.10 -4. Si mostra che, sc l'abbondanza dell'dHc 2+ cresce, lo stesso avviene dalle abbondanze di Si, 047 e 0 ~e. Infine, le flut~uazioni della densit~ dell'O +6 sembrano essero an~icorrelate a quellc dell'~He ~+.

(*) T r a d u z i o n e a eura de l la R e d a z i o n e .

QUIET TIME SOLAR-WIND IONIC COMPOSITION 697

Hovmbn] COCTa~ co~meqHOrO BeTpa B cnogoiim~n] nepao~.

Pe31oMe (*). - - I I p n HpOBe~eHHH m-xa3Metmoro arcnepHMenTa I S E E - B E G D qaCTO O6Hapy~HBamT TYI)KCJ~Ir HOHI~I B COJIHe~HOM BCTpC. OT6Hpa~ ~ a n e p ~ o ~ a CIIOKO~tHOFO

COJIHeqHOFO BCTpa, M~ IIOKa31~IBaeM, qTO pa3Jmq~LbIe HOHHI~IO KOMIIOHeHTI~I HO~tIHHHIOTCff

aHaHOrtxqHbIM T e H , ~ e ~ B H3MeHeHHHX O6~M~OH c ropocT . . B H3MeHeHH,qX riYIOTHOCTH

KHCYIa qaCTHI~. I-[pH OCO6OHtIO CHOKOff[HOM COHHeqHOM BeTI3~ HOHt~le KOH~eHTpaHt4H B

yxa3armsxx nByX cnys aax COCTaBHgIOT (rio OTHOmeHIRO K H +) 4 H e 2 + ~ ( 1 - - 5 ) ' 1 0 -z, 8Ho~+= 1.5" 10 -4, 0 + ~ 10 -4, 0 + 7 = 3 �9 10 -4, S i = l " 10 -4, F c = 0 . 5 " 10 -4. M~,I noxa- 3MBaOM, "qTO e,C.llH pacnpocTpanennocTb 4He ~+ yBeau,taBaeTca, TO liMeeT Me~"TO yBeau,~enne pacnpocTpaaem~ocTe~ Si, 0 +7 K 0+0. ~ n y K T y a u - - nnoTnOCTn 0+0, no-Barm~oMy, aHTH-

roppennpyIoT c q b n y ~ T y ~ nnoTnOCTa 4He~+.

(*) Hepeaec)eno pec)ar~ue~.